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SOMITE (IMS) 



mil 

EXPUUIATORV SUPHEMHIT 

Edited by 

C. A. Beichman, G. Neugebauer, H. J. Habing, P. E. Clegg, 
and T. J. Chester 


Prepared under the auspices of 

The Joint IRAS Science Working Group 


| Scientific and Technical Information Division 1988 

I \J# \ National Aeronautics and Space Administration 

Washington, DC 


The Joint IRAS Science Working Croup was composed of: 



G. Neugebauer 

1977-1984 


R. J. van Duinen 

1977-1982 


H. J. Habing 

1977-1984 

H. H. Aumann 

1977-1984 

T. de Jong 

D. A. Beintema 

1977-1984 

F. J. Low 

N. Boggess 

1977-1984 

P. L. Marsden 

J. Borgman 

1977-1981 

S. R. Pottasch 

P. E. Clegg 

1977-1984 

B. T. Soifer 

F. C. Gillett 

1977-1984 

R. G. Walker 

M. G. Hauser 

1977-1984 

J. P. Emerson 

J. R. Houck 

1977-1984 

E. Raimond 

R. E. Jennings 

1977-1984 

M. Rowan-Robinson 


Co-chairman 1977-1984 
Co-chairman 1977-1982 
Co-chairman 1982-1984 


1977-1984 

P. R. Wesselius 

1982-1984 


1977-1984 

B. Baud 

1982-1984 


1977-1984 

C. A. Beichman 

1982-1984 

J 

1977-1983 

T. N. Gautier 

1982-1984 

1 

1977-1984 

S. Harris 

1982-1984 


1977-1984 

G. K. Miley 

1982-1984 


1979-1984 

F. M. Olnon 

1982-1984 

- 

1979-1984 

E. Young 

1982-1984 

- 

1979-1984 



- 


Library of Congress Cataloging-in-Publication Data 

Infrared Astronomical Satellite (IRAS) catalogs and atlases 7 prepared 
under the supervision of the Joint IRAS Science Working Group, 
p. cm.— (NASA reference publication : 1190) 

Contents: v. 1. Explanatory supplement / C.A. Beichman ... [et 
al.] — v. 2. The Point source catalog declination range 90° 

[greater than delta greater than] 30° — v. 3. The Point source 
catalog declination range 30° [greater than delta greater than] 0° 

—v. 4. The Point source catalog declination range 0° [greater 
than delta greater than] -30° — v. 5. The Point source catalog 
declination range -30° [greater than delta greater than] -50° — 
v. 6. The Point source catalog declination range -50° [greater than 
delta greater than] -90 a — v. 7. The Small scale structure catalog 
/George Helou and D. W. Walker. 

1. Infrared astronomy— Observations— Catalogs. 2. Infrared 
astronomy— Atlases. 3. Infrared source— Observations— Catalogs. 

4. Infrared sources — Atlases. 5. Infrared Astronomical Satellite — 
Catalogs. 6. Infrared Astronomical Satellite— Atlases. I. Joint 
IRAS Science Working Group. II. Series. 

QB470.154 1988 
523.8 9 — del 9 


For sale by the Superintendent of Documents, U.S. Government Printing Office 
Washington, D.C. 20402 


TTin 


PREFACE TO THE SECOND EDITION 


The Explanatory Supplement was first released in November 1984 as a preprint (Jet Pro- 
pulsion Laboratory D-1855) to accompany the initial distribution of the IRAS Catalogs and 

Atlasses. Since that time a small number of errors have been found in both the data products 
themselves and in their description in the Supplement. Rather than completely rewrite the Sup- 
plement to remedy its shortcomings, we have chosen to fix only a few glaring errors in the main 
body of the text and to insert as Chapter XII a discussion of differences between the initial and 
present versions of the data. As a result of this choice, some inconsistencies in references to 
Chapters XII, XIII and XIV may be found in the text. 

This version of the Supplement describes the data available as of March 1987 at the 

National Space Science Data Center (NSSDC) at the Goddard Space Flight Center; the 

interested reader is referred to the NSSDC for access to the IRAS data. 


C. Beichman 
Pasadena 
March 1987 


PREFACE TO 1985 EDITION 


The attached Explanatory Supplement is being released in a preliminary version with the 
various catalogs and atlases because the information contained in this Supplement is essential 
for using the IRAS data. Many sections, but especially Chapters VII, VIII and XI, could not be 
written before the production of the catalogs was completed. As a result, the analysis is neither 
as complete nor as detailed as might be desired. Undoubtedly, the Supplement and catalogs 
contain errors. We have chosen to release the Explanatory Supplement in preprint form at this 
time. A hard-bound version of the Supplement and the Catalog will be issued shortly after 

February 1, 1985. 

Any reader who finds errors in the text or the Catalog before February 1, 1985, is 
requested to communicate such errors to: 


C. Beichman or G. Neugebauer 
c/o IRAS Project Office 
California Institute of Technology 
Mail Stop 100-22 
Pasadena, California 91125 


TABLE OF CONTENTS 


Page 

l. INTRODUCTION 

A. General Overview I- 1 

A. 1 The IRAS Mission I- 1 

A. 2 The Explanatory Supplement 1-1 

A. 3 Cautionary Notes 1-2 

B. Summary Description of Catalogs and Atlases 1-3 

B. 1 Point Sources 1-3 

B.2 Small Extended Sources 1-3 

B.3 Sky Brightness Images 1-3 

B.4 Low Resolution Spectra 1-3 

B.5 The Extragalactic Sub-catalog 1-3 

C. Overview of Infrared Sky 1-4 

n. SATELLITE DESCRIPTION 

A. Introduction II- 1 

B. The Spacecraft II- 1 

B. 1 Onboard Computers and Software II-2 

B.2 Attitude Control II-2 

B. 3 Communication II-2 

C. Telescope System Overview II-3 

C. 1 Cryogenics II-3 

C.2 Thermal Control II-6 

C.3 Optics n-7 

C.4 Focal Plane Assembly II-9 

C.5 Electronics 11-19 

Appendix II. 1 Data Compression 11-25 

m. THE IRAS MISSION 

A. Requirements III-l 

B. Constraints III-2 

B. 1 Introduction III-2 

B.2 Attitude Control III-2 

B.3 Solar Radiation III-2 

B.4 Earth Radiation III-2 

B.5 Moon and Planets III-2 

B.6 South Atlantic Anomaly III-3 

B.7 Station Passes III-5 

B.8 Constant Sun Angle III-5 

B. 9 Eclipse Operations III-8 

C. Design III-9 

C. 1 Basic Strategy m-9 


v 




C.2 The Second Six Months 

III- 13 


C.3 Scan Rate 

III- 13 


C.4 Strategy during South Atlantic Anomaly Passage 

III- 13 


C.5 Moon and Jupiter Avoidance Strategy 

III- 13 


C.6 Strategy of Attitude and Photometric Calibration 

m-17 


C.7 Realization of Survey Strategy 

m-17 


C.8 Half-Orbit Constraint 

III- 17 


C.9 Lune Constant 

III- 17 


C.10 Hole Recovery Strategy 

m-17 


C. 1 1 Pre-Survey Observations 

m-17 

D. 

In-Right Modifications 

m-19 


D. 1 Introduction 

III- 19 


D.2 Polar Homs 

m-19 


D.3 Operations Problems 

m -19 


D.4 Saturation 

III-20 


D.5 The 5° Gap 

III-21 


D.6 Early Eclipse and Warm Up 

III-22 

IN-FLIGHT TESTS 


A. 

Detector and Focal Plane Performance 

IV- 1 


A. 1 Detector Sensitivity and Responsivity 

IV- 1 


A.2 Detector Reliability and Anomalies 

IV-3 


A.3 Cross-scan Response 

IV-3 


A.4 Verification of Linearity 

IV-3 


A. 5 Baseline Stability 

IV- 11 


A.6 Particle Radiation Effects 

IV- 11 


A.7 Effects of Bias Boost 

IV- 12 


A.8 Photon Induced Responsivity Enhancement 

IV- 12 


A.9 Feedback Resistor Nonlinearity Analysis 

IV- 15 

B. 

Spectral Passband Verification 

IV- 16 


B. 1 Verification of the Relative Consistency 

IV-17 


B.2 Verification of the Nominal Inband/Out-of-Band Transmission 

rv-i7 

C. 

Optical Performance 

C.l Optical Cross Talk due to Bright Sources 

IV- 19 


Crossing the Focal Plane 

C.2 Optical Cross Talk from Sources not directly 

IV- 19 


on the Focal Plane 

IV- 19 


C.3 Out-of-Field Rejection Monitoring 

IV-20 

D. 

Internal Reference Source Stability 

IV-22 

DATA REDUCTION 


A. 

Overview 

V-l 


A. 1 General 

V-l 


A.2 IRAS Catalogs and Atlases 

V-l 


A.3 Processing Summary 

V-l 

B. 

Pointing Reconstruction 

V-6 

C. 

Source Detection 

V-9 


C. 1 Square Wave Filter 

V-9 


C.2 Noise Estimator V-10 

C.3 Timing Estimate V-ll 

C.4 Correlation with Point Source Template V-l 1 

C.5 Determination of Templates V-l 3 

C.6 Low Signal- to-Noise Detections V-l 3 

C. 7 Source Shadowing V-l 3 

D. Point Source Confirmation V-l 3 

D. 1 Processing Overview V- 1 3 

D.2 Overview of Seconds-Confirmation V-l 6 

D.3 Band-Merging V-24 

D.4 Known Source Correlation V-26 

D.5 Overview of Hours-Confirmation V-28 

D.6 Overview of Weeks-Confirmation V-31 

D.7 Auxiliary Processing for Low Resolution Spectra V-33 

D.8 Flux and Confusion Status Words V-33 

D. 9 Conversion of Position Uncertainties to Gaussian Approximation V-34 

E. Overview of Small Extended Source Data Processing V-34 

E. l Potential Detections V-35 

E.2 Seconds-Confirmation V-36 

E.3 Source Construction and Hours-Confirmation V-36 

E.4 Cluster Analysis Processing V-37 

E.5 Weeks-Confirmation V-38 

E.6 Band-Merging V-39 

E.7 Optimizing the Processor V-40 

E.8 The Small Extended Source Catalog V-46 

F. Asteroids and Comets V-47 

G. Extended Source Products V-48 

G. 1 Processing Overview V-48 

G.2 Quality Checking, Selection, and Weights V-48 

G.3 Phasing, Sorting, and Gaps V-49 

G.4 Conversion to Surface Brightness V-49 

G.5 Compression and the Time-Ordered Files V-49 

G.6 Destriping V-50 

G.7 Projection into Sky Maps V-50 

G.8 Consistency Checking and Removal of Bad Data V-5 1 

G. 9 Final Image Generation V-52 

H. The Point Source Catalog V-52 

H. 1 Processing Overview V-52 

H.2 Clean-Up Processing V-53 

H.3 Neighbor Tagging V-53 

H.4 Cirrus Flagging V-54 

H.5 Average Flux Computation and Variability Analysis V-55 

H.6 High Source Density Regions V-56 

H.7 Catalog Source Selection V-64 

H.8 Low-Resolution Spectral Associations V-64 

H.9 Associations V-64 


VI. FLUX RECONSTRUCTION AND CALIBRATION 

A. Data Processing: Removal of Telescope Transfer Function VI- 1 

A. 1 Digital Electronics VI- 1 


vii 



A.2 Analog Electronics Amplifiers VI-2 

A.3 Trans-impedance Amplifier VI-2 

A.4 Removal of Coherent Detector Noise VI-3 

A.5 Feedback Resistor Vl-4 

A. 6 Summary VI-4 

B. Determination of Relative Flux VI-5 

B. 1 Overall Procedure to Determine Relative Photometry VI-5 

B.2 Photometry of Point Sources and Small Extended Sources VI-6 

B.3 Photometry of Extended Emission VI-6 

B. 4 Problems VI- 12 

C. Absolute Calibration VI- 1 9 

C. 1 General Philosophy VI- 1 9 

C.2 Point Source Calibration VI-20 

C.3 Color Correction VI-27 

C.4 Absolute Calibration of Extended Emission VI-28 

D. Comparison of IRAS Observations with Ground Based Observations VI-28 


Vn. ANALYSIS OF PROCESSING 

A. Overview 

B. General Statistics of the Point Source Processing and Catalog 
B. 1 The Generation of Reliable Point Sources 

B. 2 Distribution of Sources in the Catalog 

C. Positional Accuracy 

C. 1 Positional Accuracy of Catalog Sources 

C. 2 Accuracy of Scan-by-Scan Pointing Reconstruction 

D. Photometric Accuracy 

D. l Absolute Calibration Uncertainty Checks 

D.2 Relative Photometric Accuracy 

D.3 Variable Sources 

D. 4 Discrepant Fluxes 

E. Point Source Processing Considerations 

E. 1 The Nature of Rejected Sources 

E.2 Bright Source Problems 

E.3 Sources of Incompleteness 

E.4 Effects of Failed Detectors 

E. 5 Setting the Seconds-Confirmation Threshold 

F. Asteroids and Comets 

F. 1 Number Present in Catalog 

G. Associations 

H. Meaning of Point Source Flags 
H. 1 Confusion Flags 

H.2 Cirrus Flags 

I. The Small Extended Source Catalog 

J. Extended Source Products 

J. 1 Zodiacal Emission Effects 
J.2 Effective Resolution 

J.3 Tests of Extended Source Calibration Consistency 


viii 


VII- 1 I 

VIM 

VII- 1 1 

vn-2 \ 

vn-2 l 

VII-2 

VII-10 ; 

VII- 11 

VII-11 s 

VII- 13 

VII-22 : 

VTI-24 

VII-25 t 

VII-25 ! 

VII-28 l 

VII-31 

VII-32 1 

VII-32 f 

w 

VII-33 ' ! 

VII-33 * 

VII-35 j 

VII-36 

VII-36 i 

VII-37 ’ 

i 

VII-38 l 

m 

VII-38 l 

VII-38 

VTI-39 k 

VII-39 ! 

i 

1 

f 

! 


Min 



m 


vra. SKY COVERAGE, CONFUSION, COMPLETENESS AND RELIABILITY 

A. Introduction VTII- 1 

B. Sky Coverage ' VTII-1 

C. Point Source Confusion VIII-2 

D. Reliability and Completeness of Point Source Catalog VTII-4 

D. 1 Definitions, Assumptions and Limitations VTII-4 

D.2 Formalism for the Determination of Completeness and Reliability VTII-5 

D.3 Estimation of Parameters VTII-7 

D.4 Completeness and Reliability Outside of the Galactic Plane VIII-8 

D.5 Completeness and Reliability in the Galactic Plane VIII- 10 

D.6 Galactic Plane Shadow VIII- 1 1 

E. Completeness and Reliability of the Catalog of Small Extended Sources VIII- 1 1 

IX. THE LOW RESOLUTION SPECTRA 

A. Instrumentation IX- 1 

A, 1 Introduction IX- 1 

A. 2 Optical Properties IX- 1 

A. 3 Electronics IX-3 

A.4 Effects of the Zero-Clamp IX-3 

A. 5 Summary of Instrumental Characteristics IX -4 

B. Performance and Calibration IX-4 

B. 1 Detectors IX-4 

B.2 Wavelength Scale IX-6 

B.3 Cross-Scan Responsivity IX-6 

B.4 Wavelength-Dependent Responsivity IX-6 

B.5 Radiation Effects IX-6 

B.6 Multiplexer Glitches IX-6 

B.7 Confusion IX-8 

B.8 Photon Induced Responsivity Enhancement IX-8 

B.9 Memory Effects IX-9 

B.10 Linearity Checks IX-9 

B. 1 1 Overall Flux-Density Scale IX-9 

C. Data Processing IX-9 

C. 1 The Database IX-9 

C.2 Processing the Individual Spectra IX- 10 

C.3 Averaging Spectra, Quality Checks IX- 1 1 

C. 4 Final Selection of Spectra IX- 1 2 

D. Classification IX- 1 3 

D. l Introduction IX- 13 

D.2 Classification Scheme IX- 13 

D. 3 Performance of the Classification Scheme IX- 19 

E. Some Characteristics of the Catalog IX-20 

E. 1 Completeness IX-20 

E.2 Checks on the Shape of the Spectra IX-20 


IX 


X. THE FORMATS OF THE IRAS CATALOGS AND ATLASES 


A. Introduction 

B. Point Sources 

B. 1 The Machine Readable Version of the Point Source Catalog 
B.2 The Printed Version of the Point Source Catalog 
B.3 The Working Survey Data Base 


X-l 

X-2 
X-2 
X-10 
X-l 2 


C. The Small Extended Source Catalog 


X-29 


D. Extended Emission 

D. 1 Introductory Comments 

D.2 Map Projections and Transformation Equations 

D.3 16.5° Images 

D.4 Galactic Plane Maps 

D.5 Low-Resolution All-Sky Maps 

D.6 Zodiacal Observation History File 

D.7 Coordinate Overlays 


E. Low-Resolution Spectra 
E.l Catalog Header File 
E.2 Spectra Records 

Appendix X.Ap. 1 Regions of High Source Density 
Appendix X.Ap.2 Location of 16.5° Image Fields 


Appendix X.Ap.3 Sample FITS Headers 

Appendix X.Ap.4 Zodiacal Observation History File (ZOHF) Format 


X-30 

X-30 

X-30 

X-32 

X-36 

X-37 

X-37 

X-37 

X-38 

X-38 

X-38 

X-45 

X-49 

X-51 

X-61 


XI. KNOWN PROCESSING ANOMALIES 


A. 

Processing of Extended ("Cirrus”) Sources as Point Sources 

XI- 1 

B. 

Instability and Lag of the Noise Estimator 

XI- 1 

C. 

Frequency Dependence of Responsivity with Amplitude 

XI- 1 

D. 

Errors in Cross-Scan Uncertainties Related to Failed Detectors 

XI-2 

E. 

Photon-Induced Responsivity Enhancement 

XI-2 

F. 

Artifacts in the Digital Image Data Base 

XI-2 

G. 

Photometric Processing 

XI-3 

H. 

Insufficient Specification of HCON Coverage 

XI-3 

I. 

Position Uncertainties 

XI-3 

J. 

Overestimated Weak Fluxes 

XI-4 

K. 

Minor Problems 

XI-5 


x 


I 1 11 



XII. ERRATA AND REVISIONS AS OF 1987 XII- 1 

A. Version 2.0 of the Point Source Catalog XII- 1 

A.l The Flux Overestimation Correction XII- 1 

A. 2 Additional Flux Density Changes XII- 12 

A. 3 New and Deleted Sources XII- 12 

A. 4 Revised Completeness Estimates for Version 2.0 XII- 14 

A. 5 Associations XII- 19 

A. 6 Source Names XII-20 

A. 7 Revised Positional Uncertainties for Bright Sources XII-21 

A. 8 Correction of Point Source Neighbor Counts XII-2 1 

A. 9 Spurious 25 pm Only Sources XII-21 

A. 10 Working Survey Data Base and Ancillary File XII-21 

B. Total Intensity Data XII-2 

B. 1 Total Intensity Maps XII-2 

B.2 Version 2.0 of the Zodiacal Observation History File XII-2 

C. Low Resolution Spectrometer XII-22 

D. Other Anomalies Fixed in This Release X 11-23 

XIII. CONTRIBUTORS TO IRAS (Formerly Chapter XII) XIII- 1 

XIV. AREA COVERAGE PLOTS (Formerly Chapter XIII) XIV- 1 

XV. INDEX XVI 


xi 


lo lo to 



INDEX OF TABLES 


Page 

II. SATELLITE DESCRIPTION 

C.I Telescope System Physical Characteristics II-5 

C.2 Telescope Optical Characteristics H-8 

C.3 Characteristics of Survey Array 11-12 

C.4 Survey Array Optical Characteristics 11-16 

C.5 Spectral Response H-18 

C.6 Electrical Characteristics of Survey Array 11-20 

in. THE IRAS MISSION 

C. 1 Mission Chronology HI- 1 8 

IV. IN-FLIGHT TESTS 

A. 1 Detector Data Based on NGC6543 Scans IV-8 

B. l Out-of-Band Rejection Test Sources IV- 18 

B.2 Out-of-Band Rejection Test Results IV- 18 

V. DATA REDUCTION 

D. l Confirmation Summary V-16 

D.2 Input Data for In-Band Seconds-Confirmation V-17 

D.3a Confusion Status Bit (CST AT) Assignments V-21 

D.3b Common CST AT Values V-21 

D.4 Order of Band-Merging V-24 

D. 5 Flux Status (FSTAT) Values V-33 

E. la Small Extended Source Processing Results in Region A V-47 

E.lb Small Extended Source Processing Results in Region B V-47 

H.l Catalogs Used for Associations with IRAS Sources V-65 

VI. FLUX RECONSTRUCTION AND CALIBRATION 

B. l TFPR Model Parameters VI- 10 

C. I Difference Between Survey and Pointed Observations VI-21 

C.2 Comparison with Ground-Based Observations VI-22 

C.3 IRAS Magnitudes of Calibration Stars in Pointed Mode VI-23 

C.4 Color Temperatures of Asteroids between 25 and 60 pm VI-24 

C.5 Ratio of Observed to Model Fluxes of Asteroids VI-25 

C. 6 Color Correction Factors VI-26 

D. 1 Comparison with Ground-Based Observations VI-29 

VII. ANALYSIS OF PROCESSING 

B.l Number of HCONs in WSDB and Final Catalog VII-3 

B. 2 Spectral Classification of Catalog Sources VII-3 

C. 1 Absolute Position Difference Statistics VII-5 

C. 2 IRAS-SAO Position Differences at Seconds-Confirmation VII- 1 1 

D. 1 IRAS Survey vs. Ground Based Magnitudes VII- 1 2 

D.2 Point Source Catalog Relative Photometric Uncertainties VTI- 15 


xii 


1 II: 



INDEX OF TABLES - Continued 


Page 

D.3 Gaussian Fits to Distributions of Photometric Ratios VII- 18 

D. 4 Relative Flux Changes of Variable Sources at 12 and 25 pm VII-22 

E. 1 Reasons for Rejection of a Band in High Source Density Regions VII-26 

E. 2 Bright Source Neighbors Suppressed as Cross-Talk VII- 3 1 

F. l IRAS Names of Sources Contaminated by Numbered Asteroids VII-34 

F.2 Frequency of Flux Contaminaion by Asteroids VII-34 

F.3 Cirrus Sources Possibly Contaminated by Asteroids VII-35 

vm. SKY COVERAGE, CONFUSION, COMPLETENESS AND RELIABILITY 

D. 1 Completeness and Reliability Data in 7 HCON Area VIII-8 

D.2 Number of HCONs in a 7 HCON High Source Density Region VIII- 1 1 

IX. THE LOW-RESOLUTION SPECTRA 

B. 1 Detector Characteristics IX-5 

D. 1 Spectral Classification Scheme IX- 1 5 

X. THE FORMATS OF THE IRAS CATALOGS AND ATLASES 

A. l Format of Header Files X-l 

B. 1 Format of Point Source Catalog Tape X-3 

B.2 Meaning of Hex Encoded Flags X-6 

B.3a The Catalog Working Survey Data Base (WSDB) X-l 7 

B.3b Ancillary WSDB File X-l 8 

B.4 Meaning of the Source Association Fields X-20 

B.5 Known Source ID’s X-25 

B.6 Detector Number X-25 

B.7a CLEAN Bit Assignment X-26 

B.7b BRIGHT/ ACCEPT Bit Assignment X-26 

B.7c HSDPROC High Source Density Processor Rags X-27 

B.7d MISC Bit Assignment X-28 

D. l Plates Missing From The Third Sky Coverage X-33 

E. la Header Information for Catalog of Spectra X-40 

E. lb Format of Spectra in Catalog X-41 

Apl.l Format of File of High Source Density Bins X-47 

Ap2.1 16.5° Field Centers (Equinox 1950) X-49 

Ap3. 1 Sample FITS Headers X-5 1 

Ap4. 1 Format of Zodiacal History File X-62 

Ap4.2 Table of Missing SOPs X-62 

XI. KNOWN PROCESSING ANOMALIES 

F. l Fields Deleted from 3rd HCON XI-3 

J.l Relation between Fraction of Possible Sightings and Flux Overestimate XI-5 


XIII 



INDEX OF TABLES - Continued 


Page 

XII. ERRATA AND REVISIONS AS OF 1987 

A. Version 2.0 of the Point Source Catalog XII- 

A. 1 IRAS Data Products XII- 

A.2 Flux Overestimation Parameters XII- 

A.3 Revised Portion of Catalog Tape Format XII- 1 

A. 4 New or Deleted Point Sources XII- 13 

A. 5 Bright Sources Missing from the Point Source Catalog with 161 > 10° XII- 14 

A. 6a Flux Densities for a Given Completeness Level (PSC-1) XII- 16 

A.6b Flux Densities for a Given Completeness Level (PSC-2) XII- 1 6 

A. 7a 12 pm Completeness from SSC XII- 17 

A.7b 60 pm Completeness from SSC XII- 17 

A.8 Best Estimate of Completeness Level (PSC-2) XII- 19 

A.9 Changed Gliese Associations XII-20 

A. 10 Changed SAO Associations XII-20 

D. Other Anomalies Fixed In this Release XII-23 

D. 1 Status of Anomalies From First Release XII-23 


t o VO l o I O 



INDEX OF FIGURES 


Page 

I. INTRODUCTION 


C. 1 Sky Coverage of the IRAS Survey 1-5 

C.2 Distribution of IRAS Sources with Stellar Characteristics 1-6 

C.3 Distribution of IRAS Sources with Extragalactic 1-7 

Characteristics 

C.4 Distribution of IRAS Sources Detected only at 100pm 1-8 

n. SATELLITE DESCRIPTION 

A. 1 Spacecraft Configuration II- 1 

B. 1 Spacecraft Control Axes II- 3 

C. l Telescope System Configuration IT-4 

C.2 Cross-Section View of Main Cryogen Dewar II-6 

C.3 Cross-Section View of Optical Subsystem 11-7 

C.4 Internal Reference Source Assembly II-9 

C.5 Rejection Performance of Telescope System II- 10 

C.6 IRAS Focal Plane Schematic II- 1 1 

C.l Infrared Subarray Module 11-14 

C.8 Focal Plane Array II- 1 5 

C.9 Response vs. Wavelength of Optical Components II- 1 7 

C.10 Preamp and Bias Supply Schematic 11-19 

C. 1 1 Sample Resistance vs. Voltage Curve II-2 1 

C.l 2 Focal Plane Array Infrared Channel Data Flow 11-22 

Ap.l Data Compression 11-25 

m. THE IRAS MISSION 

B. 1 Schematic of Orbital Geometry III-3 

B.2 Solar Radiation Constraint HI-4 

B.3 Earth Radiation Constraint III-4 

B.4 Earth Infrared Constraint m-5 

B.5 Combined Constraints at Two Epochs III-6 

B.6 Sample Orbital Tracks Through SAA III-7 

B.7 Attitude Control Coordinate System III-7 

B. 8 The "Banana Effect" III-8 

C. l Lune Strategy III- 10 

C.2 Scan Coverage Within a Lune HI- 1 0 

C.3 Typical Day’s Survey Coverage III- 11 

C.4 Lune Coverage Scheme III- 1 2 

C.5 Half-Circle Scans on Celestial Sphere III- 12 

C. 6 First, Second and Third HCON Coverage III- 1 4 

D. 1 Lune Coverage Constraints III-20 

D.2 The SAA and Polar Homs III-2 1 

D.3 Schematic of Final Coverage of Sky Coverage Prior III-22 

to First Eclipse of Satellite 


xv 



INDEX OF FIGURES - Continued 


Page 

IV. IN-FLIGHT TESTS 


A. 1 Histograms Giving Sensitivities of Each Detector IV- 1 

A.2 Histograms of Uncorrected Responsivities IV-2 

A.3 Cross-Scan Response of NGC 6543 IV-4 

A.4.1 Measurements of 12 and 25 pm Response vs. Dwell Time IV-9 

A.4.2 Measurements of 60 and 100 pm Response vs. Dwell Time IV- 10 

A.5 Survey and Half Survey Rate Scans IV- 1 1 

A.6 Effects of Bias Boost on Responsivity IV- 13 

A.7 Photon Induced Responsivity Enhancement IV- 1 4 

A.8 Decay of Photon Induced Responsivity Enhancement IV- 14 

A.9 Photon Dosage Crossing Galactic Plane IV- 1 5 

A. 10 Demonstration of Feedback Resistor Linearity IV- 16 

C. 1 Map of Moon Glints IV-20 

C. 2 Out-of-Field Rejection IV-21 

D. 1 Internal Reference Source Stability IV-22 

V. DATA REDUCTION 

C. 1 The Square Wave Detection Filter V-10 

C. 2 Repesentative Point Source Templates V-14 

D. 1 Confirmation Decision Parameter V- 1 5 

D. 2 Confusion Processing V-19 

E. 1 Cluster Processing at High Galactic Latitudes V-4 1 

E.2 Cluster Processing at Low Galactic Latitudes V-42 

E.3 Effect of Weeks-Confirmation Theshold V-43 

E.4 Optical Threshold for Band-Merging V-44 

E.5 Optimal Thresholds V-45 

H.l High Source Density Bins V-58 


VI. FLUX RECONSTRUCTION AND CALIBRATION 


A. 1 Feedback Resistor Model VI-4 

B. 1.1 Variation of Total Sky Brightness at 12 and 25pm VI-7 

B.1.2 Variation of Total Sky Brightness at 60 and 100pm VI-8 

B.2 Effects of Photon Induced Responsivity Enhancement VI- 14 

B.3 Source Brightness on Ascending/Descending Scans VI- 15 

B.4 Responsivity Enhancement VI- 16 

B.5 Responsivity Enhancement Correction VI- 17 

B.6 Corrected Source Brightness on Ascending/Descending Scans VI- 1 8 

D.l Comparison of Uranus Measurements VI-30 

VII. ANALYSIS OF PROCESSING 

B. 1 Galactic Latitude Distribution of Sources VII-4 

C. l Position Differences (In-Scan) for Stars VII-6 

C.2 Position Differences (Cross-Scan) for Stars VII-6 

C.3 Position Differences (In-Scan) for Galaxies VII-7 

C.4 Position Differences (Cross-Scan) for Galaxies VII-7 


xvi 


l in 



INDEX OF FIGURES - Continued 


Page 


C. 5 Observed Position Differences vs. Quoted Uncertainties V1I-9 

D. 1 12 vs. 25 pm Flux Densities VTI- 13 

D.2 25 vs. 60 pm Flux Densities VTI- 14 

D.3 60 vs. 100 urn Flux Densities VTI- 14 

D.4 HCON-HCON Repeatability for Faint Sources VII- 1 6 

D.5 HCON-HCON Repeatability for Bright Sources VII- 17 

D.6 Normalized HCON-HCON in 12 pm band VII- 19 

D.7 Normalized HCON-HCON Repeatability for Bright Sources VTI-20 

D.8 Normalized HCON-HCON Repeatability for Faint Sources VTI-2 1 

D.9 Correlated HCON-HCON Flux Variations at 12 and 25pm VII-23 

D. 10 Number of Sources with a Given Probability of Variability VTI-24 

E. 1 Effects of High Source Density Criteria VTI-27 

E.2 Optical Cross-Talk: IRC+ 102 16 at 60 pm VTI-29 

E.3 Optical Cross-Talk: Area Around IRC+ 1 02 1 6 in WSDB VII-30 

Apl- Ditribution of Various Types of Sources in VII-40 

24 Galactic Coordinates -63 


VIII. SKY COVERAGE, CONFUSION, COMPLETENESS AND RELIABILITY 


C. l Differential Source Counts VTI-3 

D. 1 12 pm Completeness for 2 and 3 HCON Surveys VIII-9 

D.2 Wavelength Dependent Effects of Galactic Plane Shadow VIII- 1 2 

IX. THE LOW-RESOLUTION SPECTRA 

A.l Wavelength Dependence of Spectrometer Characteristics IX-2 

A.2 Optical Layout of Spectrometer IX-2 

A.3 AC-Coupling and Zero Clamping Circuit IX- 3 

A. 4 Effects of Background Slopes IX-5 

B. 1 Cross-Scan Responsivity of Detectors IX-7 

B.2 Wavelength Dependent Responsivity IX-8 

D. 1 Classification Scheme for Blue Spectra IX- 1 4 

D. 2 Representative Low Resolution Spectra IX- 1 6 

E. 1 Plot of log(number) vs. log / v for IX-2 1 

Spectrometer Sources 


X. THE FORMATS OF THE IRAS CATALOGS AND ATLASES 

B. 1 Explanation for Format of Printed Version of Point X-5 

Source Catalog 

D.l Map of 16.5° Image Boundaries in Equatorial X-34 

Coordinates 

Apl Scheme for Obtaining 1 sq.deg Bins in Ecliptic Coordinates X-48 


XVII 



INDEX OF FIGURES - Continued 


ERRATA AND REVISIONS AS OF 1987 

Page 

A. 1 

12 nm Flux densities before and after correction 

X 1 1-5 

A. 2 

25 urn Flux densities before and after correction 

XII-6 

A. 3 

60 nm Flux densities before and after correction 

X 11-7 

A.4 

100 (am Flux densities before and after correction 

XI1-8 

A. 5 

12 (am and 25 pm colors of stars before and after 
correction for flux overestimation 

XII- 1 1 

A. 6 

Differential source counts 

XII- 1 5 

A. 7 

Single HCON completeness 

XII- 1 8 


xviii 


1 1 n 



I. INTRODUCTION 


A. General Overview 
A. 1 The IRAS Mission 

The primary mission of the Infrared Astronomical Satellite (IRAS) was to conduct a sensitive and 
unbiased survey of the sky in four wavelength bands centered at 12, 25, 60, and 100 pm. The project 
was initiated in 1975 as a joint program of the United States, the Netherlands, and the United Kingdom. 
Launched in January 1983, IRAS ceased operations in November 1983 after having successfully surveyed 
more than 96% of the sky. 

The results of several portions of the IRAS mission are given in a catalog of infrared point sources, 
in a catalog of extended sources smaller than 8', in a catalog of low-resolution spectra, and in an atlas of 
absolute surface brightness images of the entire infrared sky. These catalogs give the characteristics of 
some 250,000 point sources and 20,000 small extended sources down to a limiting flux density, away 
from confused regions of the sky, of about 0.5 Jy at 12, 25 and 60 pm and about 1.5 Jy at 100 pm for 
point sources, and about a factor of three brighter than this for small extended sources. The angular 
resolution of the instrument varied between about 0.5' at 12 pm to about 2' at 100 pm. The positional 
accuracy of sources detected by IRAS depends on their size, brightness and spectral energy distribution 
but is usually better than 20". Approximately 5000 8-22 pm spectra of survey sources brighter than 10 Jy 
at 12 and 25 pm are available. 

A.2 The Explanatory Supplement 

This Explanatory Supplement is intended to be a complete and self-contained description of the 
IRAS mission in relation to the products of the survey. In Chapter II, the IRAS satellite, telescope and 
focal plane instrumentation are reviewed. The elements of the mission profile--the constraints, the design 
features, and the in-flight modifications to that design— are described in Chapter III and are accompanied 
by a chronology of the events of the mission. In-flight tests of those aspects of the performance of the 
instrument directly associated with the survey are presented in Chapter IV. Chapters V and VI describe 
the processing performed on the data; the summaries that precede the detailed discussions should be 
sufficient to acquaint the user with the contents of the catalogs. Since the flux reconstruction and calibra- 
tion of the instrument probably hold intrinsic interest for many readers, these are described separately in 
Chapter VI. 

A preliminary analysis of some of the statistical properties of the catalogs is given in Chapter VII. 
Emphasis is placed on general statistics, such as positional and photometric accuracy and on easily 
derived number counts. A preliminary analysis of the sky coverage and of the completeness and reliabil- 
ity of the catalog is given in Chapter VIII. The low-resolution spectrometer and the analysis of its meas- 
urements are described in Chapter IX. Chapter X explains the format and meaning of each of the entries 
in the catalogs. Each printed volume of the catalogs repeats the description of the formats of that catalog. 

In order to produce the catalogs in a timely fashion, some processing errors and anomalies could 
not be fixed; those which were discovered before the release of the data in November 1984 are described 
in Chapter XI. A compilation of the names of people who worked on the IRAS project comprises 


1-1 



Chapter XII. The last chapter provides a series of plots giving the details of the coverage of the sky by 
the IRAS survey. 

Each chapter of the Supplement was written by those members of the IRAS team whose names are 
appended to that chapter. The work described was obviously the result of efforts by many individuals 
and should not be ascribed to the authors alone. 

A. 3 Cautionary Notes 

While it is unlikely that all aspects of the instrumental performance or the data processing will be of 
interest to all readers, even casual users should familiarize themselves with the various caveats described 
in the chapters appropriate to the type of data in question. All users of IRAS data should be cognizant of 
the following crucial facts: 

a) The sky at 100 pm is dominated by filaments termed " infrared cirrus " which, although concen- 
trated near the Galactic plane, can be found almost all the way up to the Galactic poles (Fig. I.C.4). The 
primary, deleterious effects of the cirrus are that it can generate well-confirmed point and small extended 
sources that are actually pieces of degree-sized structures rather than isolated, discrete objects and that it 
can corrupt 100 pm, and occasionally 60 pm, measurements of true point sources (Sections V.H.4, 
Vffl.D.2). 

b) The spectral bandwidths of the detectors were sufficiently wide that the quoted flux densities 
depend on the assumed energy distribution of the source. For the catalogs, the energy distribution was 
taken to be constant in the flux per logarithmic frequency interval. If the source has a different energy 
distribution than this, a color correction, as large as 50% in extreme cases, must be applied to the quoted 
flux densities (Section VI.C.3). 

c) The survey is clearly confusion limited within about 10° of the Galactic plane and in several areas 
of the sky such as the Ophiuchus and Orion-Taurus regions. Considerable effort has been made to select 
only highly reliable sources in such areas, at the expense of completeness. The flags associated with 
sources with possible confusion-related problems should be examined very carefully (Sections V.D.8, 
V.H.6, VIII.C, Vm.D. and X.B). 

d) The algorithm used to estimate the detector noise suffered from a significant lag. This caused an 
under-estimate of the true noise when approaching regions of rapidly changing noise and an over- 
estimate of the noise when leaving such areas. Regions with large and rapidly varying numbers of sources, 
such as the Galactic plane, also produced this effect. Since the source detection algorithm (Section V.C) 
thresholded on signal to noise ratio, the overestimated noise level resulted in a dearth of sources, or a 
shadow, in the areas observed just after passage across the Galactic plane. At 60 and 100 pm, where the 
effect is worst, a "coverage hole" can extend as far as 2° from the plane. The density of detected sources 
can differ, totally artificially, by as much as a factor of ten from one side of the plane to the other due to 
this shadowing (Section VIII.D). 

e) While great pains were taken to confirm the reality of sources in the point and small extended 
source catalogs, no such attempt was made for the sky brightness images. Instead, separate images of the 
sky taken at times differing from weeks to months are given. It is the responsibility of the user to ensure 
that sources seen in the images are not due to transient sources such as asteroids. 


1-2 


I 1 1 


B. Summary Description of Catalogs and Atlases 

The IRAS data are presented in different ways depending on the angular sizes of the structures 
involved: 

B. 1 Point Sources 

Sources that appeared as point-like are presented in three different ways depending on their reliabil- 
ity and on the detail of information given for the sources. 

a) A catalog of some 250,000 well-confirmed point sources is available in both printed and machine 
readable forms. Positions, flux densities, uncertainties, associations with known astronomical 
objects and various cautionary flags are given for each object. The information available in this 
catalog should satisfy almost all users. 

b) A file known as the Working Survey Data Base, available only on magnetic tape, is intended for 
researchers requiring specialized information about the observational and processing history of a 
source in the catalog. 

c) A file of rejected sources, available only on magnetic tape, includes any sources that did not meet 
the reliability criteria of the catalog. Some of these sources will be wholly spurious and due, for 
example, to detector noise, space debris, radiation hits or processing errors; others will be solar sys- 
tem objects such as asteroids and comets; and some will be true extra-solar system objects that 
failed to meet the confirmation criteria due to their faintness or variability, or to confusion effects. 
Sources in that small portion of the sky which received only limited survey coverage will also be 
found here. 

B.2 Small Extended Sources 

Sources larger than point-like, but smaller than in angular extent are to be found in the catalog of 
small extended sources which is available in both printed and machine readable versions. 

B.3 Sky Brightness Images 

The overall view of the sky, and the repository of IRAS data for structures larger than 8 , is found 
in the images of 212 fields that cover the entire celestial sphere. The fields are 16.5° on a side and have 
been imaged in each of the four wavelength bands with 2' pixels and 4-6 ' resolution. As many as four 
images based on observations separated by a few weeks to a few months are presented. These data give 
absolute surface brightness and are available in both digital and photographic representations. 

These image data are also available in a galactic coordinate projection for galactic latitudes within 
10° of the plane, and, at degraded resolution, in an Aitoff projection over the entire sky. A file giving the 
time history of the total sky brightness measured by IRAS is available at 0.5° resolution. 

B.4 Low-Resolution Spectra 

Point sources which are bright in the 8-22 |im range may have been detected by the low resolution 
spectrometer (Chapter IX). Spectra are available in both printed ( Astronomy and Astrophyics Supplement 
Series 1985) and machine readable forms. 


1-3 



B. 5 The Extraealactic Sub-Catalog 

A catalog of well-confirmed sources that are positionally associated with previously identified extra- 
galactic objects is available in printed and magnetic tape versions. All the sources in this catalog are con- 
tained in either the catalogs of point or small extended sources, but additional information about the 
associated galaxies and quasars, obtained from a variety of astronomical catalogs, is presented here. 

C. Overview of Infrared Sky 

The various depths of coverage by the IRAS survey are displayed in Fig. I.C. 1. The clear areas 
in the middle plot were covered with at least two sets of confirming scans, while the clear areas in the 
bottom plot were covered with confirming scans three or more times. Because the basic requirement for 
inclusion in the IRAS catalogs was that an object had to be observed with at least two sets of confirming 
scans, the clear portion of the middle plot represents the basic area covered by the IRAS survey. The 
shaded areas in the top plot show the areas of sky that were missed entirely. 

The general distribution of well-confirmed point sources observed by IRAS is shown in Galactic 
coordinates in Fig. I.C.2 to I.C.4. Three classes of source covering almost all objects in the point source 
catalog can be defined according to spectral energy distribution: most of 1 30,000 sources that are brighter 
at 12 pm than at 25 pm are stars (Figure I.C.2); most of the 50,000 objects that are brighter at 60 pm 
than at 25 pm and which are located more than 20° from the Galactic plane are external galaxies (Fig. 
I.C.3); most of the 35,000 sources detected only at 100 pm are cold, dense clumps within the interstellar 
cirrus (Fig.I.C.4). 

Authors: 

G. Neugebauer and C. A. Beichman. 


ORIGIN At PAGE IS 
OF POOR QUALITY 


EQUATORIAL RIGHT ASCENSION AND DECLINATION 




Figure I.C. 1 Sky coverage of the IRAS survey. Three plots of the entire sky are shown with an equal 
area projection in equatorial coordinates (see text). 


1-5 




ORIGINAL PAGE IS 
OF POOR QUALITY 


ORIGINAL PAGE IS 

of: poor quality 








II. SATELLITE DESCRIPTION 


ORIGINAL PAGE IS 
OF POOR QUALITY 


A. Introduction 

The satellite consisted of two main parts, the spacecraft and the telescope system as shown in Fig. 
n.A.l. The overall dimensions of the satellite, with deployed solar panels, were: height 3.60 m, width 
3.24 m, depth 2.05 m. The spacecraft and telescope system are described here with sufficient detail to 
allow the user of the catalog to understand how observations were made and the various limitations to 
the accuracy of these measurements. A fuller description of the spacecraft is given by Pouw (1983). 

B. The Spacecraft 

The features of the spacecraft of most relevance to the acquisition of the astronomical data were the 
control of the satellite, the method of executing the observational program, and the storage and transmis- 
sion of data. To understand these, it is necessary to discuss the onboard computers and their associated 
software, the attitude control system, data recording and the communications links. 



Figure II. A. 1 Spacecraft configuration. 



B. 1 Onboard Computers and Software 

Two identical computers provided redundancy. Each computer had a central processing unit with 
32,000 16-bit words of random access memory, which could be accessed by either central processor, and 
3000 words of read only memory. Each read only memory contained the routines essential for the safety 
of the satellite, for command handling and for the generation and downlink of housekeeping data. In 
particular, transition to read only memory control of the satellite could be, and indeed was, triggered by 
anomalous software or hardware behavior. 

The random access memory contained the routines for executing a complete Satellite Operations 
Plan (hereafter denoted an SOP) for the ten to fourteen hour period between passes over the ground sta- 
tion and for generating the scientific data stream. Although 64,000 words of memory were available, this 
capacity was sometimes insufficient to store as large a program of observations as could be carried out 
during the observation period; this was a result of the high efficiency of the ground system at filling the 
observation time and applied particularly to the third coverage of the sky during the last four months of 
the mission (Section III.C.2). 

B.2 Attitude Control 

The satellite attitude was controlled by three orthogonal reaction wheels; excess momentum was 
dumped via magnetic coils to the Earth’s magnetic field as necessary. The attitude, and changes in atti- 
tude, were sensed by a combination of an horizon sensor, a sun-sensor and three orthogonal gyros. The 
z-axis gyro was used in all modes of control and was duplicated to provide a redundant backup. 

The spacecraft control axes are shown in Fig. II.B. 1 . In observational modes, the y-axis was always 
kept perpendicular to the satellite-Sun vector while the x-axis corresponded closely to the telescope 
boresight. Only two of the many possible control modes are described here. During normal operations, 
the signals to control rotations about the x- and y-axes were obtained from a two-axis Sun-sensor with 
3.5" x 7" resolution. Signals from the z-gyro were used to control the rotation about the z-axis at the 
rate necessary to achieve the desired rate of scan, dy/dt, across the sky (see Section III.C.3). Towards the 
end of the mission, the Sun was eclipsed by the Earth for a time during each orbit and the Sun-sensor 
could not be used. At such times, all three gyros were used to control the satellite, although with a 
marked loss of control accuracy. No scientific data were taken during eclipses (see Section III.B.9). 

Onboard attitude updates and ground attitude reconstruction were made using a two-axis star- 
sensor of the V-slit type in the focal plane of the telescope. Section V.B describes in detail the attitude 
reconstruction process. The absolute pointing accuracy for control purposes of the system was approxi- 
mately 30". The accuracy of reconstructed positions is discussed in Sections V.B and VII.C. 

B.3 Communication 

The data for successive SOPs were recorded alternately on two tape recorders. In record mode, ear- 
lier data were erased. During a ground-station pass, the data recorded during the previous SOP were 
transmitted to the ground from one recorder while the other was commanded into its record mode ready 
for the data from the next SOP. This procedure protected data from being immediately over-written on 
the occasions when it proved impossible to transmit all the data to the ground during the prime station 

pass. 


II-2 


f TOWARDS ECLIPTIC NORTH POLE 
X (S) 



TO SUN 


Figure II.B.l Spacecraft control axes labeled x, y and z. The axes x(s), y(s) and z(s) are fixed with 
respect to the Sun and the north ecliptic pole. 


C. Telescope System Overview 

The IRAS telescope system configuration is shown in Fig. II.C.l. The telescope system comprised 
the upper part of the satellite and was composed of a two mirror, Ritchey-Chretien telescope mounted 
within a toroidal superfluid helium tank, which in turn was mounted within the evacuated main shell. 
The optical system was protected from contamination before launch and during the first week of the mis- 
sion by an aperture cover cooled with supercritical helium. After the cover was ejected, the sunshade 
limited heat flow to the aperture by blocking direct solar radiation and reflecting away terrestrial infrared 
radiation. The telescope orientation was constrained to prevent sunlight from striking the inner surface 
of the sunshade and radiation from the Earth from illuminating the radiators around the telescope aper- 
ture. The telescope was cooled by contact with the superfluid helium tank to temperatures ranging from 
2 to 5 K. The surfaces of the sunshade which could be viewed by the telescope aperture were cooled by 
a three-stage radiator to about 95 K. 

The telescope system consisted of the cryogenics (Section H.C.1), the thermal control system (Sec- 
tion II.C.2), the optics (Section II.C.3), the focal plane assembly (Section II.C.4) and the electronics (Sec- 
tion n.C.5). The telescope system also provided interfaces at and behind the image plane for the low 
resolution spectrometer (Chapter IX) and the chopped photometric channel; the latter was not used for 
the survey. Key physical characteristics of the telescope system are listed in Table II.C.l. 

C.l Cryogenics 

The telescope cryogenic system provided a 1.8 K thermal sink for controlling the temperatures of 
the optics and detectors. As shown in Fig. II.C.2, the main cryogen tank was toroidal in shape and 


II-3 


EJECTABLE 
APERTURE 
COVER 


SUPERCRITICAL 
HELIUM TANK 


SECONDARY 
MIRROR 


EARTH SHIELD 

PRIMARY 
MIRROR 


FOCAL 

PLANE 

ASSEMBLY 


DUTCH 

ADDITIONAL 

EXPERIMENT 

ELECTRONICS 


COVER 

LOW 

THRUST 

VENT 



SUNSHADE 


OPTICAL 

BAFFLE 


SUPERFLUID 
HELIUM TANK 

EXPERIMENT 

ELECTRONICS 

DUTCH 

ADDITIONAL 

EXPERIMENT 

CRYOGENIC 
VALVES AND 
MAN FOLD 


HORIZON 

SENSOR 


SPACECRAFT 


Figure ILC. 1 Telescope system configuration. 


surrounded the optics and focal plane. Because maximum mission lifetime required isolating the cryogen 
from external heat loads, the tank was suspended from nine fiberglass straps to isolate it thermally from 
the exterior main shell. Three shields cooled by venting gaseous helium and 57 layers of multilayer insu- 
lation provided additional isolation between the cryogen tank and the main shell. The helium gas left the 
main cryogen tank through a porous plug made from densely packed sintered stainless steel. The plug 
allowed vapor to vent while retaining the superfluid liquid. The telescope and focal plane instruments 
were cooled through the attachment of the optics subsystem to the main cryogen tank near the primary 




Table II.C.l Telescope System Physical Characteristics 

CRYOGENICS 

Outer shell temperature 

195 K 

Main dewar capacity 

78 kg superfluid helium 

Cryogen temperature 

1.8 K 

Aperture cover dewar capacity 

6 kg supercritical helium 

THERMAL CONTROL 

Optics, Focal Plane 

Cryogenic 

Aperture cover 

Cryogenic 

Sunshade 

Passive radiator, heater 

Electronics 

Surface coatings, blankets 

Main Dewar 

Multilayer insulation, shading, 
passive radiator 

OPTICS 

Type 

Two mirror, Ritchey-Chretien 

Mirror material 

Beryllium 

Baffle material 

Aluminum 

Entrance pupil diameter 

57 cm 

Obscuration diameter 

24 cm 

Operating temperature 

2 to 5 K 

FOCAL PLANE ASSEMBLY 
Detector, feedback resistor 

operating temperature 

2.6 K 

JFET operating temperature 

70 to 80 K 

MOSFET operating temperature 

2.6 K 

Number of detectors 

62 infrared, 8 visible 

power dissipation 

14 mW 

Construction 

Modular: 8 infrared subarrays 
2 visible subarrays 

ELECTRONICS 

Preamplifier type 

trans-impedance 
amplifier, one per detector 

Number of subassemblies 

15 

Power consumption 

48.3 W 

Operating temperature 

Oto 15 C 

A/D sensitivity 

125 uV/data number 

1 Data rates 

Engineering 

128 bits per second (bps) 

Infrared data 

5888 bps 

Visible data 

128 bps 

MASS 

External thermal control 

73 kg 

Main Liquid helium dewar 

432 kg 

Liquid helium at launch 

73 kg 

Optics 

72 kg 

Focal plane instruments 

H kg 

Electronics and cables 

90 kg 

Structure and Miscellaneous 

58 kg 

Total 

809 kg 


FIBERGLASS MAIN CRYOGEN 



Figure II.C.2 Cross-sectional view of main cryogen dewar emphasizing components 
of insulation system. 


mirror. Heat loads from the aperture were coupled to the venting helium gas by a strap connecting the 
baffle assembly to a heat exchanger. The gas finally exited to space through two vent nozzles located 
symmetrically on the dewar exterior. 

Typical operating temperatures of the cryogen tank and exterior shell during flight were 1.8 K and 
195 K, respectively. The 73 kg of superfluid helium in the tank at launch gave approximately a 300-day 

lifetime. 

The aperture cover was an independent cryogenic system which included a cryogen tank, multilayer 
insulation blankets, a vapor-cooled shield, and a balanced vent system. The cover contained supercritical 
helium and operated between 6 and 15 K. A back-pressure regulator maintained the tank pressure to 
37 ±2 psia. After one week in orbit, the entire cover assembly was ejected from the telescope in prepara- 
tion for survey observations. 

C.2 Thermal Control 

The external surfaces of the telescope system were designed to minimize the main shell temperature 
and, therefore, heat loads to the cryogen. The sunshade protected the telescope aperture from solar radia- 
tion when the telescope was pointed more than 60° from the telescope-Sun line. The specular inner sur- 
face of the sunshade minimized heat loads into the dewar by presenting a cold surface (95 K) to the 
aperture and by reflecting away radiation from the Earth. The dewar’s location behind the solar panel 


n-6 



assembly also reduced the solar heat load. The Earth shield, located on the lower side of the telescope and 
facing away from the Sun, partially blocked terresterial radiation, while the dewar wall opposite the solar 
panels radiated unwanted heat to deep space. A large multilayer insulation blanket between the spacecraft 
and the lower dewar-shell minimized heat flow in that area. The signal processing electronics boxes were 
mounted on low conduction composite trusses and surrounded by blankets to reduce heat input into the 
dewar. The cables connecting the focal plane outputs on the main shell to the exterior electronics boxes 
were fabricated of low thermal conductivity stainless steel coaxial cables. For further discussion of the 
thermal performance, see Urbach (1984). 

C.3 Optics 

The optical subsystem (Fig. II.C.3) imaged the infrared and visible light onto the focal plane. The 
two-mirror Ritchey-Chretien telescope was made of beryllium to reduce mass and minimize thermal dis- 
tortion upon cooling to cryogenic temperatures. The secondary mirror was coated with aluminum to 
enhance its reflection at visual wavelengths. 

The telescope optical parameters and performance are given in Table II.C.2. The design goal for 
the image quality was that it be diffraction limited in all infrared bands. This goal was met except at 12 
pm. Since the telescope was intended to be a survey instrument rather than a high resolution imaging 
instrument, the poor image quality at 12 pm did not interfere with the mission. For further discussion of 
the optical system, see Hamed, Hamed and Melugin (1981). 



Figure II.C.3 Cross-sectional view of optical subsystem. 


II-7 



Table II.C.2 Telescope Optical Characteristics 

DESIGN PARAMETERS 

Primary mirror diameter 

60 cm 

Unvignetted field of view 

63.6 dia. 

System focal length(design) 

550 cm 

Back focal length 

18.35 cm 

TELESCOPE PRESCRIPTION 

Primary mirror vertex radius 

-180.0 cm 

Secondary mirror vertex radius 

-36.48 cm 

Primary eccentricity 

1.00569 

Secondary eccentricity 

1.43206 

Primary-secondary spacing 

74.74 cm 

TELESCOPE PERFORMANCE 

Entrance pupil diameter 

57 cm 

Central obscuration diameter 

24 cm 

Effective collecting area 

2019 cm 2 

System focal length (measured) 

545 cm 

System F/number 

9.56 

Plate scale at focal plane (measured) 

1.585 mm/ 

Diameter of 80% encircled energy 

12 pm 

25 

25 pm 

25” * 

60 pm 

60"* 

100 pm 

100"* 

Infrared surface reflectivity 

(all bands) 

96% 

* diffraction limited 


An assembly mounted behind the secondary mirror contained ten thermal calibration sources, 
hereafter called "internal reference sources", several of which were used to provide stable pulses of 
infrared radiation for use as a reference during the mission and for ground testing prior to launch. Figure 
II C 4 shows the location of the internal reference source assembly, the way in which a source illuminates 
the focal plane through a small hole in the center of the secondary mirror, and a cutaway view of an indi- 
vidual thermal source. The thermal source consisted of a 1 mm square diamond substrate coated with 
nichrome film and suspended by 0.05 1 mm diameter brass wires. During the mission an applied voltage 
ohmically heated the substrate to -200 K in 13/16 sec. Two optical sources were included in the call- 
bration assembly and used for ground testing of the star sensors. 


Out-of-field radiation was absorbed by aluminum baffle structures which were coated with Martin 
Optical Black. Figure II.C.5 shows the calculated out-of-field performance in the four wavelength bands. 
The survey strategy (Section ffl.C) limited the angle between the boresight and the Moon, Earth, Sun and 
Jupiter to greater than 24", 88", 60" and 5", respectively. At these angles the out-of-field radiation from 


11-8 





Figure II.C.4 


tntemd reference source assembly showing radiation path from 
and details of thermal source design 


source to focal plane 


these sources is thought to be negligible (see, however. Section III.B.5 and IV.C for a discussion of lunar 
(f n«) ). Further discussion of the out-of-Seld performance is included iu Hamed, Brcault, and Melngin 

C.4 Focal Plane Assembly 

The focal plane assembly contained the infrared and visible detectors, cold electronics, and associ- 
ated masks, filters and field optics. It consisted of 62 infrared channels and eight visible channels The 
infrared channels were divided into eight modules, two for each color band with each module containing 
either seven or eight detectors. Figure II.C.6 shows the layout of the focal plane and the numbers 


II-9 




Figure II.C.5 


Iculated out-of-field rejection performance of telescope system compared to the total 
photom«ric reference, or TFPR (Section VI.B), for 12, 25, 60 and 100 pm bands. 
; sunshade temperature was taken to be 95 K; the Earth was turned “ “ k * 
) K blackbody; the moon, Sun and Jupiter were taken as 370, 5000 and 133 K black 
lies with angular diameters of 31, 31 and 0.75 , respectively. 


II- 10 


WAVELENGTH BANDS, fjm 



STAR 

SENSORS 


Figure II.C.6 A schematic drawing of the IRAS focal plane. The numbered rectangles in the central 
portion each represent the field of view of a detector, filter and field lens combination. 
The image of a source crossed the focal plane in the Y direction as indicated. The 
filled-in detectors were inoperative while the cross-hatched detectors showed degraded 
performance during the mission. 


assigned to individual infrared detectors. Table II.C.3 lists the positions of the center of each detector 
mask relative to the boresight and the size of each mask projected through the optical system onto the 
sky. The detector masks were rectangular in aspect and infrared sources scanned across the focal plane 
parallel to the narrow dimension of the detectors in all observational modes. 

Figure n.C.7 shows an exploded view of the focal plane. Infrared radiation passed through the field 
mask and spectral filters and was focused by the field lens onto the aperture defining the detector cavity 
entrance. The detector cavities were constructed of Au:Pt alloy to provide local high-Z shielding to 
absorb y-rays with energies less than about 100 keV. In addition, for the 60 and 100 pm detectors, the 
cavities were designed as reflecting integrating cavities to increase photon absorption in the Ge:Ga 
material. 

The visible wavelength channels were similar to the infrared wavelength channels in construction, 
except that they used visible light filters, no field lens, and silicon diode detectors. The visible wavelength 
detectors were placed in a double "V" arrangement in order to provide two-axis spacecraft attitude infor- 
mation during star crossings. 


n-n 






Table II.C.3 Characteristics of Survey Array 



Det. 

Offset 

Nominal 

Mask Location 

Size 



Step 

Gain 1 

arc min 

arc min 


No. 

(0-7) 


Y 4 Z 5 

AZ 

AY 



— 

100 nm Band, Module B~ 



1 

3 

26.6 

27.87 8.71 

5.05 

3.03 

2 

3 

24.6 

27.80 0.04 

5.05 

3.03 

3 

3 

25.4 

27.86 -8.62 

5.05 

3.03 

4 

4 

24.1 

23.83 12.86 

4.68 

3.03 

5 

3 2 

26.2 

24.04 4.37 

5.05 

3.03 

6 

4 

25.9 

23.65 -4.29 

5.05 

3.03 

7 

3 

24.7 

23.78 -12.77 

5.05 

3.03 




-60 jam Band, Module B— 



g 

3 

21.2 

19.64 9.80 

4.75 

1.51 

9 

3 

20.2 

19.72 1.14 

4.75 

1.51 

10 

4 

21.3 

19.74 -7.53 

4.75 

1.51 

11 

3 

21.1 

19.70 -14.46 

1.30 

1.51 

12 

3 

21.3 

17.20 13.49 

3.45 

1.51 

13 

3 

23.1 

17.19 5.47 

4.75 

1.51 

14 

3 

21.0 

17.20 -3.20 

4.75 

1.51 

15 

4 

21.3 

17.20 -11.86 

4.75 

1.51 




-25 nm Band, Module B~ 



16 

4 

12.3 

14.01 8.71 

4.65 

0.76 

17 3 

3 

11.8 

14.04 0.04 

4.65 

0.76 

18 

4 

11.3 

14.04 -8.62 

4.65 

0.76 

19 

4 

13.1 

12.24 12.96 

4.48 

0.76 

20 3 

4 

11.7 

12.27 4.37 

4.65 

0.76 

21 

4 

11.8 

12.26 -4.29 

4.65 

0.76 

22 

4 

11.9 

12.27 -12.88 

4.48 

0.76 




- 1 2 nm Band, Module B— 



23 

4 

14.2 

9.47 9.81 

4.45 

0.76 

24 

4 

14.8 

9.46 1.14 

4.45 

0.76 

25 

4 

15.5 

9.47 -7.52 

4.45 

0.76 

26 

3 

14.3 

9.48 -14.50 

1.20 

0.76 

27 

4 

14.2 

7.71 13.55 

3.33 

0.76 

28 

3 

15.3 

7.71 5.47 

4.55 

0.76 

29 

4 

13.9 

7.70 -3.19 

4.55 

0.76 

30 

4 

14.5 

7.71 -11.86 

4.55 

0.76 



Table II.C.3 Characteristics of Survey Array (Cont.) 


Det. 

Offset 

Nominal 

Mask Location 


Size 


Step 

Gain 1 

arc min 



arc min 

No. 

(0-7) 


Y 4 

Z 5 

AZ 

AY 



-60 pm Band, Module A- 

- 



31 

3 

20.8 

4.56 

14.55 

1.28 

1.51 

32 

3 

20.8 

4.59 

7.61 

4.75 

1.51 

33 

3 

22.6 

4.58 

-1.06 

4.75 

1.51 

34 

3 

20.8 

4.59 

-9.73 

4.75 

1.51 

35 

3 

21.0 

2.06 

11.94 

4.75 

1.51 

36 3 

3 

20.7 

2.06 

3.27 

4.75 

1.51 

37 

4 

20.8 

2.11 

-5.40 

4.75 

1.51 

38 

3 

18.9 

2.10 

-13.41 

3.47 

1.51 



-25 pm Band, Module A- 




39 

4 

15.2 

-1.16 

14.05 

2.33 

0.76 

40 

4 

15.7 

-1.16 

6.55 

4.65 

0.76 

41 

3 

14.7 

-1.16 

-2.12 

4.65 

0.76 

42 

4 

16.1 

-1.14 

-10.78 

4.65 

0.76 

43 

3 

13.9 

-2.92 

10.88 

4.65 

0.76 

44 

4 

14.8 

-2.92 

2.22 

4.65 

0.76 

45 

3 

15.3 

-2.93 

-6.45 

4.65 

0.76 

46 

3 

15.4 

-2.92 

-13.95 

2.33 

0.76 



-12 pm Band, Module A- 




47 

4 

14.5 

- 5.67 

14.64 

1.18 

0.76 

48 

3 

14.0 

- 5.67 

7.65 

4.55 

0.76 

49 

4 

14.4 

- 5.67 

-1.02 

4.55 

0.76 

50 

2 

14.1 

-5.66 

-9.68 

4.55 

0.76 

51 

3 

14.2 

- 7.42 

11.98 

4.55 

0.76 

52 

3 

14.3 

- 7.43 

3.32 

4.55 

0.76 

53 

4 

14.4 

-7.43 

-5.35 

4.55 

0.76 

54 

4 

13.8 

-7.42 

-13.41 

3.36 

0.76 



-100 pm Band, Module A- 

- 



55 

4 

22.9 

-11.33 

13.95 

2.52 

3.03 

56 

4 

27.3 

-11.42 

6.55 

5.05 

3.03 

57 

3 

26.2 

-11.51 

-2.12 

5.05 

3.03 

58 

4 

27.6 

-11.41 

-10.79 

5.05 

3.03 

59 

3 

26.8 

-15.34 

10.88 

5.05 

3.03 

60 

3 

26.9 

-15.49 

2.21 

5.05 

3.03 

61 

4 

26.8 

-15.40 

-6.46 

5.05 

3.03 

62 

4 

26.5 

-15.38 

•13.85 

2.53 

3.03 


1 The ratios of nominal to low gain and nominal to high gain are 7.18, 0.107; 5.98, 0.109; 10.8, 0.102; 
and 13.4, 0.100 for the 12, 25, 60, and 100pm channels. 

2 Offset step changed after launch to 2 

3 Channel inoperative during mission 

4 During survey scans sources move from -Y to +Y 

5 Negative Z corresponds to larger angle to the Sun (0) 


11-13 






- APERTURE AND | 
FILTER BLOCK 


SPECTRAL 
FILTERS 
(8 SETS) 


LENS BLOCK 


FOCAL PLANE 
OPTICAL SYSTEM 
(FPOA) 


FIELD LENS 


KAPTON 

ISOLATION SHIM 


ANODIZED SHIM 


DETECTOR 


APERTURE 


DETECTOR MODULE 




DETECTOR 





<£> O/ 
O/ 


JFET FRAME 


DETECTOR 
CAVITY 
COVER 
PLATE ^ 


9 


m i - 

! THERMAL ISOLATION 
^ SUSPENSION FILAMENTS 




Figure II.C.7 Infrared subarray module exploded view showing module components 






Figure II.C.8 shows the focal plane filter/lens combinations and configurations for each color band 
and Table II.C.4 summarizes the optical characteristics of the survey array. The detailed optical system 
transmission, detector spectral response and overall relative spectral response for the four infrared bands 
are shown in Figs. II.C.9.a,b and are listed in Table II.C.5. These parameters were determined from pre- 
flight measurements of sample filters, field lens material and, with the exception of the 100 pm detectors, 
spare flight detectors. The response of the Ge:Ga detectors used at 100 pm was assumed to be the same 
as that of the 60 pm detectors even though the material came from a different source. Additional details 
of the focal plane optics can be found in Bamberg and Zuan (1984) and Darnell (1984). 


12 BAND 


25 Mm BAND 



MULTI-LAYER 
BANDPASS FILTER 



MULTI-LAYER 
BANDPASS FILTER 


60 ^.m BAND 




1 mm 

SAPPHIRE 




/ 


\7 


SHORT X 
BLOCKER 



14^.m OF 
ZnO POWDER 


100 Mm BAND 



15 OF 

DIAMOND 

POWDER 


Figure II.C.8 Focal plane array filter and lens components and configuration. 


11-15 




Table II.C.4 Survey Array Optical Characteristics 


BAND 

12 pm 

25 pm 

60 pm 

100 pm 

FILTER FUNCTIONS 
Short Wavelength 
Blocking 

MLIF* 

MLIF 

MLIF + 
Sapphire + 
ZnO powder 

MLIF + 
Sapphire + 
CaF 2 + 
KCL + 
Diamond 
Powder 

Short Wavelength 
Cuton 

MLIF 

MLIF 

Sapphire + 
MLIF 

KCL 

Long Wavelength 
Cutoff 

Long Wavelength 
Blocking 

MLIF 

BaF2+ 
MLIF + 
Si:As 

Si:Sb 

Si:Sb 

KRS-5 

KRS-5 

Ge:Ga 

Ge:Ga 

FIELD LENS 
Materials 
Anti-reflection 
coating 
Focal length 
Exit pupil diameter 

Ge 

MLIF 

6.53 mm 
1.0 mm 

Si 

MLIF 

6.59 mm 
1.0 mm 

Ge 

A/4 

parylene-C 
6.98 mm 
1.0 mm 

Ge 

A/4 

parylene-C 
8.34 mm 
1.16 mm 

DETECTORS 
Materials 
size (LxW), mm 
Electrode 
Spacing (mm) 

Si:As 

1.0x1.78 

0.64 

Si:Sb 

1.0x1.78 

0.71 

Ge:Ga 
1.5x1. 5 

1.0 

Ge:Ga 

1.25x1.25 

1.25 

OPTICAL PERFORMANCE 
Bandwidth (FWHM) 
Average Inband 
Transmission 

7.0 pm 
0.54 

11.15 pm 
0.50 

32.5 pm 
0.29 

31.5 pm 
0.34 

OUT-OF-BAND LEAKS 
Short Wavelength 
Long Wavelength 

<2 x 10' 5 
<3 xlO' 6 

<5 x 10' 5 
<2 xlO -4 

<2x 10' 4 ** 
<2 xlO' 4 

< 9 x 10' 2 ** 
<5xl0' 2 

*MLIF - multi-layer interference filter 
**See Section II.C.4 










Figure II.C.9 a) Response vs. wavelength of optical components. Solid lines show the transmission of 
filters and lenses. Dashed lines show relative detector response to constant energy input; 
b) Relative system spectral response. 


Out-of-band leaks listed in Table II.C.4 are defined as the ratio of the integrated energy longward or 
shortward of the 2% relative response wavelength to the integrated inband energy. The calculations were 
based on illumination from a 2000 K blackbody for short wavelength leaks and a 200 K blackbody for 
long wavelength leaks. The tabulated limits come from tests of the final flight focal plane except for the 
100 pm detectors which were changed shortly before launch. These measured limits differ from an ear- 
lier set of measurements of the individual components. At 100 pm these latter tests give an out-of-band 
rejection less than 1.5 x 10~ 3 for a 2000 K source, significantly lower than the upper limit in the table. 
At 60 pm these component tests indicated the presence of a spectral leak between 1.6 and 7 pm that 


11-17 




Table II.C.5 Spectral Response 


Relative 

Det. 

resp. 


Relative 

System 

resp. 


l Relative 

Det. 
resp. 

Relative 

System 

resp. 

25 pm Band 


0.33 

0.007 

0.36 

0.101 

0.41 

0.288 

0.46 

0.388 

0.49 

0.452 

0.53 

0.521 

0.56 

0.562 

0.60 

0.626 

0.63 

0.683 

0.66 

0.729 

0.68 

0.778 


100 pm Band 



could be as large as 0.02 for a 2000 K source and 0.08 for a 10,000 K source, significantly larger than the 
tabulated limit. The origin of this discrepancy is not understood. A discussion of the in-flight tests of the 
spectral response is given in Section IV.B.2, and of the possible impact of leaks on the calibration in 
VI.C. 

C.5 Electronics 

The photoconductive detector elements responded to infrared radiation by altering their electrical 
resistance. Figure Il.C.lO.a shows schematically the nominal preamplifier and bias voltage design. A 
matched pair of junction field effect transistors (JFETs) for each detector acted as a unity-gain source- 
follower amplifier, converting the high impedance output of the photodetectors to low impedance for 
transmission to the warm electronics outside the dewar (Low 1981). The JFET pairs were each 
suspended by Dacron threads inside a small 2 K light-tight box such that electrical dissipation in the 
JFETs themselves (about 200 microwatts) maintained the JFETs at temperatures of 60 to 70 K. A3 MD 
metal-film resistor cemented to the JFET acted as a heater for cold starting the amplifier during ground 
testing and several hours after launch. The JFETs formed a differential input stage of the trans- 
impedance amplifier. At low frequencies the output voltage of the trans-impedance amplifier was equal 
to the voltage difference between the gates of the JFET pair plus any offset voltage at the input of the 
operational amplifier plus the voltage drop across the feedback resistor due to the photocurrent from the 
detector. 


a) NOMINAL DESIGN 




Figure II.C. 1 0 Preamp and bias supply schematic. The JFET module is indicated by the dashed line 

the dewar boundary by the dash-dot line. 


11-19 




The detector bias was applied to one detector contact and the trans-impedance amplifier maintained 
the other contact at a constant DC voltage very near signal ground independent of photocurrent until the 
trans-impedance amplifier output saturated at about 1 1 volts. All detectors in a module had a common 
bias voltage, applied through the module frame, which are listed in Table II.C.6. An exception to this 
biasing scheme was module A in the 25 pm band. During testing, the frame of this module became 
shorted to signal ground rendering the entire module inoperative. The alternative biasing approach used 
for this module only (suggested by Dr. J. Houck) is shown in Fig. II.C.10.b. The bias voltage, with 
reversed polarity, was applied to the gate of the reference JFET.This voltage also appeared at the opera- 
tional amplifier input so the output of the trans-impedance amplifier was compensated by the same 
amount. The net effect of this modification was to increase the gain of the trans-impedance amplifier by 

a factor of 1 .2. 


Table II.C.6” 

Electrical Characteristics 

of Survey Array 


Effective Wavelength (pm) 

12 

25 

60 

100 

Nominal bias (volts) 

3.27(A) 

2.50(B) 

1.50 

0.160 

0.185 

Boosted bias (volts) 

2.50(A) 

2.00(B) 

11.00(A) 

7.00(B) 

1.00 

1.00 

Nominal/Low Gain 

7.18 

5.98 

10.8 

13.4 

Nominal/High Gain 

0.107 

0.109 

0.102 

0.100 


During laboratory testing the responsivity and noise of the detectors were found to depend on their 
history of exposure to energetic radiation, such as y-rays, and energetic electrons and protons. The 
observed responsivity change resulting from an exposure of 0.6 Rads of Co 60 , roughly equivalent in 
dosage to a passage through a deep portion of the South Atlantic Anomaly (SAA), was about a factor of 
1 2, 2, 6 and 10 for the 12, 25, 60 and 100 pm detectors. Biasing the detectors into a "breakdown" condi- 
tion resulted in a large current flowing through the detector that annealed the radiation effects. The 
hardware provided a second bias voltage level, called the "bias boost" voltage, which effected this anneal- 
ing process after exposure to SAA protons during the mission. The bias boost voltages are listed in Table 
II.C.6 and the effects of the bias boost are discussed in Section IV.A.7. Because the 12 pm detectors did 
not require annealing, the second bias voltage provided an alternative operating bias. 

The rolloff frequency of the trans-impedance amplifier was set to approximately 80 Hz by a 0. 1 pf 
shunt capacitor across the feedback resistor. The feedback resistors, Eltec model 102 metal film resistors, 
were selected from 1 x 10 l ° D room temperature elements. At 2 K their impedance was 
2.05 ±0.1 x 10‘°D and varied slightly with voltage. Figure II.C.ll shows a sample resistance vs. vol- 
tage curve as measured at 2 K. A combination of three straight lines fitted to the measured points 
defined the shape of the non-linear resistance versus voltage relationship used for data reduction (see Sec- 
tion VI.A.5). 




Figure II.C.11 Resistance vs. voltage characteristics of feedback resistor for detector 29. Measure- 
ments were made at 2 K. Solid line shows the shape of resistance vs. voltage relation 
adopted for processing. 

The rest of the telescope electronics processed the signals from the visible and infrared 
preamplifiers, transferred data to the spacecraft onboard computer for storage and subsequent transmis- 
sion to the ground station, and received commands from the spacecraft and distributed them to the tele- 
scope systems. This data-processing was split into two major elements: the analog electronics and the 
digital electronics. More details are contained in Langford et al. (1983) and Long and Langford (1983). 

The analog signal path for the infrared detectors was entirely DC-coupled. Figure II.C.12 shows a 
functional diagram of the components of the analog electronics for one infrared detector channel. In 
order to maintain a negative voltage at the output of the analog electronics, the DC offset voltage at the 
output of the trans-impedance amplifier could be changed by 17.8 mV steps using an 8 level command- 

able offset with level 4 corresponding to no change. The commandable offset levels utilized during the 
mission are listed in Table II.C.3. 

Nuclear pulse circumvention circuitry prevented sharp pulses from cosmic rays and charged particle 
hits on the detectors in the SAA and in the polar horns from contaminating the infrared data stream. 
The output of the trans-impedance amplifier was fed to an integrator and to a pole-zero amplifier which 
flattened the frequency response to 450 Hz to improve the operation of the circumvention circuit. The 
output of the pole-zero amplifier went to a differentiator and to a Bessel filter which delayed the signal by 
about 150 ps. The differentiated and integrated signals led to a comparator which opened a switch to 
prevent the unwanted, fast rise-time pulses from passing through the system. The integrator raised the 
minimum threshold to blank the unwanted spike as the DC voltage from the trans-impedance amplifier 
increased. The track/hold capacitor clamped the input to the gain amplifier to a fixed level while the 
switch was open. Further details of the design and performance of the pulse circumvention circuitry can 
be found in Emming et al. (1983) and Long and Langford ( 1 983). 

The Bessel filter boosted the trans-impedance amplifier output by a factor of two. An additional 
amplifier could increase the system gain by software commandable factors of unity (low gain), of 5 to 12 


11-21 



GAIN SELECT 
COMMAND 



Figure H.C. 1 2 Focal plane array infrared channel data flow. 












depending on the wavelength band (nominal gain), and of ten times nominal gain (high gain). All survey 
scans were made using the nominal gain except for some of the brightest regions in the Galactic plane. 
Some of these areas were rescanned using low gain (Section m.D). The overall nominal gain for each 
infrared channel and the ratios of the nominal to low and high gains for the different detector modules 
are listed in Table II.C.3. Finally, 12 dB/octave low-pass filters with cutoff frequencies of 6, 6, 3, 1.5 Hz 
for the 12, 25, 60 and 100 pm bands, respectively, limited the frequency response and reduced high fre- 
quency noise. The outputs of the low-pass filters were fed into multiplexers and then to a 16-bit analog- 
to-digital converter operating at 125 pV per data number for subsequent processing by the digital elec- 
tronics. 

The visible detector data flow was similar to the infrared data flow, except that the trans-impedance 
amplifier was AC-coupled to a MOSFET preamplifier. For both infrared and visible channels, the pole 
zero amplifier, integrator, differentiator, comparator, track/hold capacitor, switch, gain amplifier and 
low-pass filter were contained in a single miniature hybrid circuit. 

Under low background conditions the limiting noise in the analog electronics chain was the John- 
son noise of the 2 x 10 10 £2 feedback resistor. At a temperature of 2 K this noise level was roughly 1.6 
pV Hz"' 6 . 

The digital electronics processed the digitized infrared and visible detector data, collected telemetry 
information from various sensors located on the telescope, received and executed commands issued from 
the onboard computer and transmitted the formatted telemetry, infrared and visible detector data to the 
onboard computer. The infrared detectors were sampled at 16, 16, 8 and 4 Hz at 12, 25, 60 and 100 
pm, respectively. To minimize spacecraft data storage requirements, the digital electronics compressed 
each 16-bit infrared detector value to an 8-bit value representing the difference between the successive 
16-bit numbers. For details of compression scheme, see Appendix II. 1. 

During star crossings, two of the eight visual detectors were sampled at a 500 Hz rate. The onboard 
digital electronics determined the visual magnitude of the star and its crossing time, passed this informa- 
tion to the spacecraft attitude control software to update the satellite pointing and recorded the data for 
subsequent use in the attitude reconstruction. The digital electronics also measured 108 temperatures, 
voltages and pressures to monitor the health of the telescope. These and other housekeeping data were 
multiplexed, digitized and formatted for transferral to the onboard computer. 


11-23 



Authors: 

F. C. Gillett, P. Clegg, D. Rosing, G. Neugebauer, D. Langford, A. Pouw, W. Irace, and J. Houck. 
References 

Bamberg, J.P., and Zaun, N.H. 1984, S.P.I.E. Proceedings, 509, in press. 

Darnell, R.J. 1984, S.P.I.E. Proceedings , 509, in press. 

Emming, J.G., Arenz, R.F., Downey, C.H., Long, E.C., Smeins, L.G. 1983, S.P.I.E. Proceedings, 445, 
254. 

Hamed, N„ Hamed, R„ and Melugin, R. 1981, Optical Engineering, 20, 195. 

Hamed, R., Breault, R., Melugin, R. 1980, S.P.I.E. Proceedings, 257, 119. 

Langford, D„ Simmons, J„ Ozawa, T„ Long, E.C., Paris, R. 1983, S.P.I.E. Proceedings, 445, 244. 

Low, F.J., 1981, S.P.I.E. Proceedings, 280, 56. 

Pouw, A. 1983, Journal British Interplanetary Soc., 36, 17. 

Urbach, A.R. 1984, S.P.I.E. Proceedings, 509, in press. 


11-24 



APPENDIX II.l - DATA COMPRESSION 

In order to conserve space on the spacecraft digital tape recorder, the 16-bit output of the analog to 
digital converter was compressed to an an 8-bit value. The current 8-bit data output represented the 
absolute difference between the current 16-bit analog to digital value for a detector and the same 
detector’s reconstructed 16-bit value. The reconstructed value was derived from the previous 8-bit 
compressed difference value taken during the preceding data iteration. 

The compressed 8-bit value represented only the 4 to 5 most significant bits of the difference 
between the current analog to digital converter reading and the reconstructed value. The 8-bit 
compressed value consisted of the sign of the difference (positive or negative), three significant bits of the 

difference called the sigmficand, and a shift code indicating the most significant bit of the difference as 
indicated in Fig. II.Ap.l.a. 

Figure II.Ap.l.b shows the correlation between the shift code, the positions of the 3-bit significand, 
and the most significant bits of the difference. The right column indicates the total number of significant 
bits of the difference that could be stored in the 8-bit compressed form. 

Every 4 seconds the 16-bit value, called the key value, for a detector was inserted as the first word 
of a data line. The key value was provided as a means to monitor and synchronize the data compression 
with the ground reconstruction software. 


7 6 5 4 3 2 1 0 

s X X X c c c c 

* SHIFT CODE 

- SIGNIFICAND 

SIGN 


1 MSB 


SHIFT 

CODE 


SIGNIFICANT 

BITS 

0 

1 

2 

3 

4 

5 

6 

7 

8 

9 

10 11 12 13 14 15 


1 

i 

X 

X 

li 

0 

0 

0 

0 

0 

0 

0 

0 

0 0 

0 

0 

0 

0 

0 

4 

0 

1 

1* 

X 

X 

° 

0 

0 

0 

0 

0 

0 

0 

0 0 

0 

0 

0 

0 

1 

5 

0 

0 

1 


X 

X 

° 

0 

0 

0 

0 

0 

0 

0 0 

0 

0 

0 

1 

0 

6 

0 

0 

0 

1 

* 

X 

X 

° 

0 

0 

0 

0 

0 

0 0 

0 

0 

0 

1 

1 

7 

0 

0 

0 

0 


* 

X 

X 

« 

0 

0 

0 

0 

0 0 

0 

0 

I 

0 

0 

8 

0 

0 

0 

0 

0 

• 

* 

X 

X 

« 

0 

0 

0 

0 0 

0 

0 

1 

0 

1 

9 

0 

0 

0 

0 

0 

0 

1 

X 

X 

Xi 

0 

0 

0 

0 0 

0 

0 

1 

1 

0 

10 

0 

0 

0 

0 

0 

0 

0 


X 

X 

■51 

0 

0 

0 0 

0 

0 

1 

1 

1 

11 

0 

0 

0 

0 

0 

0 

0 

0 

■>1 

X 

X 

X 

0 

0 0 

0 

1 

0 

0 

0 

12 

0 

0 

0 

0 

0 

0 

0 

0 

0 

^1 

X 

X 

X 

I 0 0 

0 

1 

0 

0 

1 

13 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

T 

T 

X 

X X 

0 

1 

0 

1 

0 

15 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

1 

0 

X 

X X 

0 

1 

0 

1 

1 

15 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

1 

1, 

X X 

X 


1 

1 

0 

0 

16 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

1 

0 

X X 

X 


1 

1 

0 

1 

16 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

1 

X X 

X 


1 

1 

1 

0 

16 

D 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

0 

X X 

4 

1 1 

1 

1 

1 

16 


XXX* THREE BITS OF SIGNIFICAND 


Figure II.Ap.l 


a) The format of the compressed 8-bit data word (top); and b) the meaning of por- 
tions of the code used for data compression (bottom). 


11-25 





III. THE IRAS MISSION 


A. Requirements 

The IRAS mission was designed to observe the entire sky in the infrared and to insure that the 
resultant survey was reliable and complete down to a specifiable flux level. This requirement was 
specified numerically as a catalog of point sources over 95% of the sky, that was 98% complete in uncon- 
fused regions and 99.8% reliable. To create accurate maps of the total sky brightness, observations of 
diffuse emission also needed to be free of significant background variations due to local effects such as 
Sun and Earth shine. It was desirable to observe the same distant sky through different regions of the 
solar system so that the effects of solar system backgrounds could be separated from Galactic and other 
backgrounds. 

Many non-astronomical sources of infrared radiation were detected by IRAS, including contam- 
inants released by IRAS itself, Earth orbiting dust particles. Earth orbiting satellites and other space 
debris. Asteroids and comets were also detected. The difficulties caused by dust particles and debris can 
be highlighted by noting that a dust particle of diameter 100 pm radiating as a blackbody at a tempera- 
ture of 200 K could be seen by the 12, 25 and 60 pm detectors with a signal to noise ratio of >3 for dis- 
tances closer than about 18 km. The infrared detectors were also sensitive to impacts by cosmic rays and 
charged particles from the Earth’s trapped radiation belts. The residual effects of these impacts could not 
be entirely removed by the on-board nuclear pulse circumvention electronics (Section II.C.5) and so were 
a potential source of difficulty. 


The requirement for a reliable survey, along with the above considerations, necessitated multiple 
observations of the sky to discriminate between the fixed (distant) astronomical objects and the moving 
(nearby) objects in the solar system, and against noise sources such as particle radiation effects on the 
detectors. By reobserving the sky on a time scale of seconds the residual effects of the radiation hits and 
the fastest moving local material were recognized and rejected (seconds-confirmation); by reobserving on 
a time scale of hours slower moving but still relatively local sources were recognized (hours-confirmation); 
and by reobserving on a time scale of weeks and months, slowly moving solar system objects such as 
asteroids were recognized (weeks-confirmation). 


It should be borne in mind that the purpose of multiple coverage was to enhance the reliability of 
the point source detections, not to coadd data to enhance sensitivity. 

The demand that a source should be seen twice within seconds was fulfilled by the design of the 
focal plane. A mission design plan (McLaughlin 1984) and a survey strategy plan (Lundy 1984) to 
achieve the required, homogeneous, multiply-confirmed coverage of the whole sky were evolved subject 
to the instrumental and observational constraints. It is necessary to explain the observational constraints 
that had to be met as these are important for understanding the strategy and the reasons that, in practice, 
the sky coverage was neither completely uniform nor homogeneous. 


III-l 


! 'I'll' 


B. Constraints 

B. 1 Introduction 

IRAS was successfully launched into its planned 900 km altitude, 99” inclination Sun-synchronous 
polar orbit with a period of 103 minutes. This orbital altitude was low enough to be below most particles 
in the Earth’s trapped particle belts yet high enough that only a negligible amount of residual atmospheric 
gasses would build up on the cold mirror surfaces during the mission. With the telescope pointing radi- 
ally outwards from the Earth and perpendicular to the Sun vector, no Earth or sunlight could enter the 
telescope and all ecliptic latitudes would be swept out during one orbit while, as the line of nodes pre- 
cessed at a rate of about 1° per day to remain perpendicular to the Sun vector, all ecliptic longitudes 
would be covered in a period of six months (Fig. III.B.l). Such a simple solution would, however, have 
allowed no flexibility. The attitude control system and telescope were designed to allow pointing away 
from the local vertical within certain constraints which are described below. 

B.2 Attitude Control 

The first constraint was designed into the spacecraft attitude control system and required the tele- 
scope to point no further away from the Sun than 120”. At greater angles the fine Sun sensor could no 
longer see the Sun well enough to function. 

B.3 Solar Radiation 

The second constraint was that the sunshield design did not allow the telescope to point closer than 
60” towards the Sun without solar radiation falling onto the inside of the sunshade. These first two con- 
straints limited the celestial sphere available to IRAS at any one epoch as shown schematically in 

Fig.m.B.2. 

B.4 Earth Radiation 

The third constraint arose from prohibiting any infrared radiation from the Earth from falling upon 
the inside of the sunshield or the top of the telescope baffle system. This constraint is shown in telescope 
coordinates in Fig. III. B.3 and the corresponding area of the sky available on any one orbit is shown in 
Fig. III.B.4. The Earth radiation constraint together with the spacecraft orbital rate determined the max- 
imum time of 15.6 min. during which the satellite could point at a given fixed celestial position, if no 
other constraints interfered. Because of the varying solar declination during the mission, the joint effect 
of these constraints changed during the mission; Fig. III.B.5 shows schematically the combined con- 
straints at two epochs six months apart. 

B.5 Moon and Planets 

Infrared radiation from the Moon and the planet Jupiter was sufficiently strong to affect the perfor- 
mance of the detectors for a significant time after passage over the focal plane. Of the other planets, only 
Venus was bright enough to have this effect but it was always too close to the Sun to be observed. An 
avoidance radius of 1° from Jupiter was set within which the telescope did not point. For the moon, an 
avoidance radius of 25” was used during the first two months of the survey but was lowered to 20” after 
April 3 except between August 26 and September 9 where it was lowered to 13”; at 25” significant "moon 
glints" were introduced into the data-stream (see Section IV.C). 


1 II! 


III-2 


NORTH 

POLE 



igure m.B.l A schematic drawing of the orbital geometry. The orbital altitude, 900 km and incli- 
nation, 99°, combined with the Earth’s equatorial bulge lead to a precession of the 
p ane of the orbit about 1 ° per day. As a result, the orbit normal always pointed to- 
wards the Sun as the satellite orbited above the Earth’s terminator. By pointing the 
satellite radially away from the Earth, the cold telescope was shielded from the heat 
loads from the Sun and Earth while providing natural scanning motion across the en- 
tire sky m about six months. A sequence of hours-confirming scans on the celestial 
sphere is also shown. 


B.6 South Atlantic Anomaly 

Another constraint was the depression in the Van Allen belts known as the South Atlantic Anomaly 
(SAA). Proton hits in the detectors when passing through the SAA increased the noise to such an extent 
that it was impossible to continue observations. Data were not taken whenever the satellite entered a 
geographically fixed flux/energy contour. As a result of analysis of the effects of radiation on the detec- 
tors, the contour shown in Fig. III.B.6 was adopted. In May, it was slightly reduced in an attempt to 
minimize the adverse effects on the survey scans. Both contours are shown in Fig. III.B.6. 


III-3 


NORTH 

ECLIPTIC 

POLE 


TO SUN 



Figure III.B.2 


With IRAS at the center of the celestial sphere, the solar constraint on soto radUhOT 
and visibility to fine Sun sensors prohibited viewing closer than 60 from the Sun and 
farther than 120°, the area shown shaded in the figure. 


270 



Fi<nire TIT B 3 The Earth infrared radiation constraint is shown in telescope coordinates. The tele- 
Ftgnre I1I.B.3 ^ ^ ^ ^ ^ +x ^ ou , of , he page . The scanning dtrec- 

tion is toward -Y. The satellite position vector had to remain within the shaded re 

gion at all times. 


III-4 





EARTH'S 

NORTH 

POLE 



Figure III.B.4 


The shaded region depicts the portion of the celestial sphere not available for viewing 
during any given orbit due to the Earth infrared constraint. The normal to the plane 
of the IRAS orbit is indicated "Trajectory Pole". 


A second effect of a passage through the SAA was a long term enhancement, by as much as factors 
often, in the responsivity and noise caused by the large radiation dosage. As described in Section II.C.5, 
these large changes could be erased by increasing the bias voltage on the detectors, a technique referred to 
as "bias boosting". 

B.7 Station Passes 

No observations could be carried out during the prime pass over the ground station (Bevan et al. 
1983), as during this time (typically 10 minutes every 10-14 hours) data from the preceding 10-14 hour 
observation period were being transmitted from the on-board tape recorders to the ground and the com- 
mands for the next 10-14 hours of observations were being sent to the satellite (Mount 1983; MacDougall 
etal. 1984). 

B.8 Constant Sun Angle 

Figure III.B.7 shows the attitude-control coordinate system of the spacecraft. Although the space- 
craft had gyros for 3-axis control, it normally used only one, the z-axis gyro (see Fig. II.B.l). Control of 
the other two axes was maintained by the fine Sun sensor which ensured that the y-axis was always per- 
pendicular to the satellite-Sun vector (see Section II.B.2). The satellite scanned with a fixed cone-angle 6 
between the telescope boresight and the Sun vector. 

A consequence of the constant Sun cone angle constraint, together with the fact that the ecliptic 
longitude of the Sun constantly increases, is that it was never possible to repeat exactly the coverage 
obtained by any scan, except at the ecliptic plane, or after an interval of six months. For example, a 


III-5 



Figure III.B.5 


NORTH 

ECLIPTIC 

POLE 



NORTH 

ECLIPTIC 

POLE 


TO SUN 




TRAJECTORY 

POLE 


Basic viewing window (unshaded) on the celestial sphere for two different dates during 
the survey. 


I Hi 


ni-6 


Figure III.B.6 


Figure III.B.7 



Sample orbital tracks through the South Atlantic Anomaly (SAA) at 900 km and the 
contours used for SAA avoidance are given. The contour (A) was determined during 
in orbit checkout; the less conservative contour (B) was used after May 9 to help 
reduce the SAA affects on the survey scans. 



The cone and clock system were used to define the scan geometry. The cone angle (9) 
is the angle between the satellite-Sunline and the boresight. The clock angle (%„) is 
measured in the plane perpendicular to the Sun-line, clockwise as viewed from the 
Sun. Sometimes y = 360°— <p is also used. 


III-7 





Figure III.B.8 The curvature of scans taken at 0 * 90" is called the "banana effect". 
See Section III.B.8. 


oole-to-pole scan executed at cone angle 0 - 90" and passing through a specilied longitude in the ecliptic 
plane would be rectangular in the ecliptic coordinate system. A pole-to-pole scan through the same 
specified longitude in the ecliptic plane a. a later or earlier time (0 * 90") would be an arc who™ 
lure increased with 190-01 (Fig. III.B.8). This "banana effect" meant that overlapping scans had to be 
made at almost the same time with only small 0 differences. 

B.9 Eclipse Operations 

Towards the end of the mission, the solar declination became such that the Sun was eclipsed by the 
Earth during part of the orbit. Because the satellite could not use its Sun sensor, gyros were used to con- 
trol all three axes during this period. This resulted in a considerable loss of control accuracy and in par- 
ticular, meant that slewing maneuvers ended in unpredictable positions. Consequently, it was decided 
not to continue survey scans or other normal observations during eclipses. 


I II 


III-8 


C. Design 

C.l Basic Strategy 

The basic time interval of the operations and data acquisition was the time between station passes, 
varying between seven or eight orbits or 10 to 14 hours. During each such period, the survey strategy dis- 
cussed below was implemented in a series of commands sent to the satellite called a Satellite Observation 
Plan, hereafter denoted as SOP. There were two SOPs per day and 600 SOPs in the entire mission. 

The strategy developed to achieve the goals of multiple survey coverage on various timescales 
divided the celestial sphere into units of half overlapping "Junes" in the ecliptic coordinate system (Lundy 
1984). Lunes were defined as the area between two ecliptic meridians 30° apart (Fig. III.C.l). The lunes 
were painted" by survey scans, one after the other, as they passed through the viewing window of the 
telescope. Figure III.C.2 illustrates the shape of a lune sketched onto a plane and, with a very exag- 
gerated focal plane width, how it was covered by different survey scans each at a fixed cone angle from 
the Sun. The first scan in a lune was placed so that it crossed the ecliptic at the lower longitude boun- 
dary of the lune. Successive scans were laid down at increasing ecliptic longitudes, each one shifted over 
by 14.23 , that is by half the width of the focal plane minus a safety margin to account for the pointing 
limit cycle of the telescope. The overlap ensured that measurements of the same area of sky were 
repeated within a few orbits (for hours-confirmation) and that the "banana effect" (Section III.B.8) was 
not too severe. The criterion for hours-confirmation was that the hours-confirming scans had to be made 
within three SOPs (34-38 hours) of each other. 

Two lunes in opposite hemispheres were observed simultaneously, one on the ascending side of the 
orbit and one on the descending side. Figure III.C.3 shows a typical day’s survey coverage and the 
regions forbidden by the constraints. After a lune was filled, a second lune in the same hemisphere was 
started; it overlapped half of the first lune, ensuring that another hour’s-confirming set of scans was 
repeated after about one to two weeks, thus providing the required repetition on the time scale of weeks 
(Fig. m.C.4). 

The observing conditions were considerably worse during some orbits than during others. Orbits 
that crossed the SAA or contained a station prime pass were interrupted for significant amounts of time. 
When a long interruption occurred, the continuation of a survey scan was sometimes impossible and a 
small hole in the coverage resulted. To minimize such events only the nine (out of 14) orbits per day 
least affected by the SAA were used for the survey scans. 

The survey strategy aimed for four coverages (two sets of hours-confirming coverages) in the first six 
months and two coverages (a third set of hours-confirming coverages) in the second six months. Any 
time left over was used to recover survey observations lost because of the various constraints, to make the 
necessary calibration observations and to carry out additional, non-survey, pointed observations to attain 
higher sensitivity or spatial resolution. For the first two weeks of observations (February 10-23, 1983, 
SOPs 31-57) half circle scans, which give redundant coverage in the ecliptic polar areas (Fig. III.C.5), 
were used for increased initial coverage as insurance against an early failure of the satellite. Subsequently, 
the more efficient lune method was used for the rest of the first six months of the survey. 


III-9 


Figure III.C.l 


Figure III.C.2 


TO SUN 


NORTH 

ECLIPTIC 

POLE 



Lunes are sketched on the celestial sphere. Shaded zone is not allowed 
by Sun constraint. 


NORTH 

ECLIPTIC 

POLE 


ECLIPTIC 



SOUTH 

ECLIPTIC 

POLE 


Half-overlapping scan swaths (highly exaggerated) fill a lune with one 
hours- confirming layer 


III- 10 



ORIGINAL PAGE IS 
OF. POOR QUALTHC 



LUNE 

No. 

0 

1 

2 

3 

4 

5 

6 

7 

8 
9 

10 

11 

12 


59 ° 74 ° 89 ° 

•— 1 j 

—1 


I— 


ECLIPTIC 

LONGITUDE 


239 ° 2 . 54 ° 269 ° 


DESCENDING LUNES 


► 1 


ASCENDING LUNES 



59 ° 

I 

I 


Figure III.C.4 


The lune coverage scheme is shown for the first two sets of hours-confirming scans. 
Each line represents a single hours-confirming coverage of an entire 30° lune (except 
short lines at beginning and end which represent half-lune coverage). 


TO SUN 


NORTH 

ECLIPTIC 

POLE 



Figure III.C.5 A zone of half-circle scans is drawn on the celestial sphere. 


Ill- 12 


C.2 The Second Six Months 


The first sky survey was completed on August 26, 1983, with nearly a full second six months of 
operations expected. The survey was continued, but aimed at a coverage by only one set of hours- 
confirmed scans and using seven orbits per day. This plan covered many of the large number of small 
regions that the survey had missed for reasons explained elsewhere in this chapter, without the time con- 
suming necessity of covering each hole separately, and also improved the survey’s completeness for 
weaker sources (signal- to-noise ratio 5-10). Another important goal was the coverage of a 5° gap left in 
the first six months’ survey (Section III.D.5). This latter goal was never achieved because the helium sup- 
ply was exhausted earlier than expected. Half circles rather than lunes were used in the second six 
months because the additional coverage in the ecliptic polar regions in August and September filled in 
regions that would be inaccessible later due to eclipses and a severe Earth infrared radiation constraint in 
December and January. 

C.3 Scan Rate 

Using two gyros, the spacecraft clock-angle <po (Fig. III.B.7) was decreased at a rate of 
(3.85/cos(9O-0))' per second (which resulted in scanning the sky at a rate of 3.85' per second independent 
of 0), i.e. 10% faster than the orbital rate of the satellite of 3.5' per second. This increased rate gave 
increased pointing flexibility with respect to the constraints and, in particular, helped to reduce the effects 
of the SAA. 

C.4 Strategy during South Atlantic Anomaly Passage 

Survey observations were interrupted for up to 14 minutes during SAA passages even after selection 
of the least affected orbits. When the telescope entered the SAA during a survey scan its scanning motion 
was halted and it remained pointing at the same point on the celestial sphere until it emerged from the 
SAA when the scan was restarted with an overlap. During the SAA passage, "bias boost" was applied 
to the 60 and 100 pm detectors until about three minutes before exit. The internal reference sources 
were flashed just before entry and just after exit at the end and beginning of survey scans (Section 
in.B.3). Some SAA passages required so much time that the telescope would have violated the Earth 
infrared constraint (Section III.B.4) after SAA passage. In such cases it was necessary to leave a hole in 
the sky coverage and try to recover it with specially prepared short scans in the preceding or succeeding 
pair of SOPs. The failure to recover all of these holes contributed to the variable depth of sky coverage. 

C.5 Moon and Jupiter Avoidance Strategy 

Jupiter was avoided by stopping a scan when it reached within 1* of the planet and side stepping by 
1° for 2° of scan before coming back to resume the original scan; hence the Jupiter avoidance procedure 
left a 2° square hole on the sky which needed to be recovered later in the survey. Figure III.C.6.1-3 
clearly show the three square holes on the ecliptic at longitude approximately 250° (RA about 17 hrs). 
Each was left by one of the three sets of hours-confirming coverages. There is no residual Jupiter hole 
when the effect of all the coverage is added. 

To avoid the moon in the same way as Jupiter would have left extremely large holes which would 
have been difficult to fill. The policy adopted, therefore, was to stop the survey on the moon’s side of the 
orbit for the approximately three days needed for the moon to pass through this approximately 40° 
avoidance region. 


Ill- 13 







ORIGINAL PAGE IS 
OF POOR QUALITY 


H I! 


Figure III.C.6.1 The first hours-confirmed coverage is overlaid on a map of the sky in ecliptic coordi- 
nates. The scans converge at the ecliptic poles. Small, equally spaced holes in the 
ecliptic plane are due to the lunar avoidance stragegy and represent gaps in the sky 
coverage used to generate the extended emission images. The point source survey 
was not affected by these gaps. 






ORIGINAL PAGE IS 
OF POOR QUALITY 


3 <i n 1 1 1 v i 


5 1 Id I 1 5 3 





ECJL 



III- 16 


ORIGINAL PAGE IS 
GS BOOR QUALITY 


Figure III.C.6.3 Same as Fig. 


C.6 Strategy of Attitude and Photometric Calibration 


To refine the pointing after the slew to the beginning point of a scan, attitude calibrations were 
made using visual stars at the beginning and end of each scan. Star sightings near the middle of the scan, 
if possible, allowed any pointing drifts to be detected and corrected (Section V.B). Long scans proved 
most suitable as they offered the greatest probability of finding these calibration stars. 

To check the photometric calibration the internal reference sources were flashed at the beginning 
and end of every scan, including scans broken by the SAA. The internal reference sources were them- 
selves regularly checked against an astronomical reference source (Section VI.A). 

C.7 Realization of Survey Strategy 

The survey scans resulting from this strategy and the constraints were generated by a computerized 
scheduling program (Oord et al. 1981). As well as generating the survey scans, the program reported 
which pieces of sky were not covered due to a combination of constraints, so that scans could be gen- 
erated separately for inclusion elsewhere in the same or other SOPs (MacDougall 1984). It was not 
always possible to generate these scans within the constraints and so holes were left. 

C.8 Half-Orbit Constraint 

The scheduling program constrained survey scans to stop at the end of a half-orbit at one of the 
ecliptic poles. A side effect of this half-orbit design was that on occasions when the satellite had looked 
back in its orbit to cover a region of sky whose observation had been interrupted by a long SAA passage 
(Section III.C.4), the scan was cut short by reaching the end of its half-orbit, even though no other con- 
straints had been violated. 

C.9 Lune Constraint 

Full coverage with the lune method required that while a lune was being painted it had to contain 
the meridian at 90° from the Sun. Otherwise the ecliptic poles could not be reached and holes would 
result. The scheduling program used to generate the survey scans also required that 9 = 90° lie within the 

lune. This geometry became a hard constraint on the lune strategy, and became relevant to the decision 
that led to the 5° gap (III.D.5). 

C. 1 0 Hole Recovery Strategy 

A record was maintained of those regions of the sky that were not covered by the automatically 
generated scans due to constraint violations and as suitable opportunities arose, attempts were made to 
fill the holes (Lau and Wolff 1984; Lundy 1984). 

C. 1 1 Pre-Survev Observations 

Before commencing survey operations numerous checks were required to verify the health and 
safety of the satellite and to determine the best modes of operation. The cooled aperture cover was kept 
on the telescope for the first six days to allow sufficient time for contaminants carried up with the satellite 
to outgas and disperse so that they would not freeze on the cold optics when the cover was ejected. The 
eight days after cover ejection were used to test those aspects of the instrument that could not be tested 
with the cover on. This period was followed by a period of repeated surveying on a limited region of sky 
to verify the survey strategy and the data processing facilities (Section VIII.D; Rowan-Robinson et al.. 


Ill- 17 



Table III.C.l Mission Chronology 


(All dates are 1983 and given in GMT) 

Date 

SOP 

Event 

26 Jan 

1 

Launch 02h 17m 

26-31 Jan 

1-12 

In Orbit Checkout (Cover on) 
Outgassing of satellite 

31 Jan 

12 

Cover ejection 19h 37m 

31 Jan-8 Feb 

12-28 

In Orbit Checkout (Cover off) 

9 Feb 

29 

SAA contour A (Fig. III.B.6) usage begins 

9-10 Feb 

29-30 

Minisurvey layer 1 Hand made scans 

10 Feb 

31 

Start first two hours-confirming 
coverages using half circles. 

11-12 Feb 

33-34 

Minisurvey layer 2 Hand made scans 

13-14 Feb 

37-38 

Minisurvey layer 3 Hand made scans 

15 Feb 

41 

Minisurvey layer 4a Hand made scans 

16 Feb 

43 

Minisurvey layer 4b Hand made scans 

23 Feb 

57 

Half circle method ended 

23 Feb 

58 

Lune method started 

3 Apr 

135 

Moon avoidance radius lowered from 25 to 
20° 

9 May 

207 

SAA contour B (Fig. III.B.6) usage begins 

26 Aug 

425 

End of first two hours-confirming coverages 

26 Aug 

426 

Start third hours-confirming coverage 

26 Aug 

426 

Moon avoidance radius lowered from 20 to 
13° 

9 Sep 

454 

Moon avoidance radius raised from 13 to 
20° 

18 Nov 

593 

First eclipse. Fallback to safety mode 

21 Nov 

600 

Survey operations resumed 1 9h 40m 

22 Nov 

600 

Liquid helium ran out OOh 16m 

22 Nov 

600 

Last survey scan started 03h 34m 

23 Nov 

603 

12 pm detector baselines saturated 09h 30m 


Ill- 18 




1984). The scans of this "minisurvey" were hand-tailored for maximum efficiency in coverage. After all 
these tests had been completed and the problems they revealed had been resolved (MacDougall et al 
1984), the all-sky survey was started. The dates of these and other important events are given in Table 

D. In-Flight Modifications 

D- 1 Introduction 

Although the implementation of the survey strategy was remarkably effective, there daily arose 
regions of sky that should have been covered but for one reason or another (to be discussed below) were 
not. Attempts made to recover these scans were not always successful with a resultant loss in depth of 
sky coverage and data quality. The checks that were made at the ground station, including a search for 
moving objects such as asteroids and comets (Davies et al 1984) are described elsewhere (Thomas 1983, 
MacDougall et al. 1984). These coverage gaps were not predictable and so recovery scans had to be gen- 
erated rapidly to allow the uplink of commands during the next prime pass some 12 hours after the 
downlink of the data that revealed the difficulty. Given the time needed to process the data to reveal the 
problem, generate the recovery observation and regenerate the SOP, this was a tight schedule to 
maintain. The occasional failures are reflected in some of the remaining coverage holes. 

The coverage obtained in each of the first two hour-confirming layers is summarized in Fig 

m.C.6.1,.2. 

D.2 Polar Homs 

An unexpected source of noise was particle hits in the horns of the Van Allen belts (Fig. 
III.D.2). The effects of the horns were variable and often negligible. However, observations affected by 
increased noise were rescanned. The criterion used to require recovery scans for trapped particle radia- 

tt°^ IT"' 10 thC P ° lar h ° ms ^ was that the on - board nuclear particle circumvention electronics (Section 
II.C.5) be activated during more than an average of about 0.7% of the time in 48 sec. This duration was 

determined empirically to correspond to roughly a doubling of the noise at 60 and 100 pm. In fact, 
because of the high degree of redundancy inherent in the data processing, the point source results were 
not noticeably affected by the level of radiation hits in the polar horns. The point source processing 
used both the data suffering from high radiation hit rates and the rescans. The extended emission maps 
(Section V.G) do not use data affected by radiation hits. In many cases the recovery scans were them- 
selves affected by high radiation as the location of the satellite relative to the horns was often similar to 
that in the original scan. Losses due to the horns were greatest at the start of the mission because of a 

period of intense solar activity, which reduced considerably later on. The radiation recovery is dis- 
cussed in greater detail by Wolff (1984). 

D.3 Operations Problems 

Scans lost because of occasional problems with the satellite or ground station (see e.g. Holdaway 
1984, and MacDougall et al. 1984) were recovered as soon as possible after the events. Occasionally the 
outage was so great that it would have been impractical, within the geometrical constraints and opera- 
tional timescales, to recover the gaps with individual recovery scans and it became necessary to reset the 

survey parameters in the scheduling program (Section III.C.7) and restart the survey from where it had 
left off. 


Ill- 19 


CALENDAR DATE (1983) 




SiSSsMOONSiSi 
illl TRACK 11 


^IIIIICONSTRAI NTllIiff 


PPjupiter;: 

till TRACK i: 


SUN-SENSOR I! 
i CONSTRAINT ils 


iliDESCENDINGi|^|^l 1 ' 

ill lune track mmm? 


ASCENDING \ 
LUNE TRACK i 


| LUNE TRACK 

90° C° NE ANGLE 


1y/ -60 SSS 0 


ECLIPTIC LONGITUDE, deg 


i:- mm I une coverage is given as a function of date. The grey shaded regions are forbidden 
F,gure ni.D.l indicated on the diagram.The dashed lines running at about 45 

^“constrain, free zones is the 90* cone angle line. The survey lune 
coverages are indicted by the solid black lines. 


D.4 Saturation 

Some saturation of the electronics by bright sources occurred as had been expected. Recovery on 
low gain was initially attempted in a systematic way and was successful for regions including the Gatactic 
Center, the Orion Nebula, M17, W43, NGC6334, NGC6357, and IRC+10216. It was later decided a 
tracking and recovering saturated portions of scans was too time consuming and that it wou more 
fruitful to concentrate on ailing missing sky coverage due to other causes. Recovery of saturated scans 
was commanded only if the bright object was not one that was, from prior knowledge expected to be 
vtuy bright. Most of these sources were well known HII regions which are so bright that the saturated 
parts can be measured from ground or airborne telescopes. 


III-20 



Figure III.D.2 The SAA (> 10 protons cm- 2 sec" 1 with E > 50 MeV and Polar Homs (> 10 5 elec- 
trons cm sec ') with E > 0.5 MeV at an altitude of 900 km (Stassinopoulos 1978). 


D-5 The 5° Gan 

The scan generation software required that each of the two lunes under observation should contain 
an ecliptic meridian 90° away from the Sun. Otherwise the coverage of the ecliptic polar caps could not 
be part of the regular survey planning. Progress in completing the survey was slow enough that these 90° 
meridians slowly drifted toward the edge of the lunes. By mid May 1983 only 1° of slack was left and, 
once the June boundary was reached, either a gap in the coverage would be required or the more efficient 
lune method would have to be replaced with the half circle method. A related problem was that it was 
becoming necessary to make survey scans at large solar offset angles (9O°-0). Such scans were particu- 
larly vulnerable to interruptions due to the SAA, station passage, Earth infrared constraint, etc. Conse- 
quently, a large number of small holes occurred in the sky coverage. As an important part of the sky, the 
North Galactic Pole, was approaching the observation window it was decided to avoid falling off the 
0 - 90° line in an uncontrolled way determined by the accumulation of the inevitable, small losses of 


III-21 


coverage. A topologically simple gap was chosen that would be accessible in the second six months of the 
mission (in early December). The survey was therefore advanced by 5°, giving a 6° pad. This decision 
was based on the prediction that the liquid helium supply would last until January 1984. Unfortunately, 
the last of the liquid helium boiled away in late November, leaving the 5° gap almost completely unob- 
served except for a few scans around its periphery. 

D.6 Early Eclipse and Warm Up 

The first eclipse of the satellite occurred 4 days earlier than expected due to an incorrect modeling 
of the effects of atmospheric refraction. The satellite fell into its safety mode during which no observa- 
tions were made and the pattern of the third set of hours-confirming scans was interrupted. The coverage 
obtained before this event is shown in Figure III.C.6.3. A correction to the eclipse prediction program was 
derived and survey operations resumed after three days. Unfortunately, on the day that the survey was 
resumed the liquid helium boil-off rate dropped to zero as the last of the cryogen was expended. The 
telescope and focal plane began to warm up. By 09h 30m on November 23, 1983, 300 days after launch, 
the baselines of the 12 pm detectors saturated and the flight part of the mission was officially terminated. 

The final depth of coverage by hours-confirming pairs of survey scans that was obtained is shown 
approximately in Fig. III.D.3 and, in more detail, in Chapter XIII. 



Figure III.D.3 A diagram of the survey coverage indicates approximately the number of hours- 
confirming coverages over the entire sky. 



Authors: 

J.P. Emerson, H.J. Habing, P.E. Clegg, S. Lundy 
References 

Bevan, H.C., McPherson, P.H., Champion, R.J.R. and Reid, M.F., 1983, Journal British Interplanetary 
Soc. 36, 10. 

Davies, J.K., Green, S.F., Stewart, B.C., Meadows, AJ. and Aumann, H.H., 1984, Nature 309, 31. 

Holdaway, R., 1984, American Institute of Aeronautics and Astronautics 22nd Aerospace Meeting 
AIAA-84-0 151. 

Lau, C.O. and Wolff, D.M., 1984, American Institute of Aeronautics and Astronautics Astrodynamics 
Conference, AIAA-84-2012. 

Lundy, S.A., 1984, American Institute of Aeronautics and Astronautics 22nd Aerospace Meeting , AIAA- 
84-0 149. 

MacDougall, J.R., 1984, American Institute of Aeronautics and Astronautics 22nd Aerospace Meeting 
AIAA-84-0 148. 

MacDougall, J.R., McPherson, P.H., Mount, K.E., and Thomas, G.R., 1984, Journal British Interplane- 
tary Soc., 37, 337. 

McLaughlin, W.I., 1984, American Institute of Aeronautics and Astronautics 22nd Aerospace Meeting 
AIAA-84-0 147. 

Mount, K.E., 1983, Journal British Interplanetary Soc., 36, 34. 

Oord, E., de Pagter, P.J., van Holtz, R.C. and MacDougall, J.R., 1981, Proc. Ini. Symp. Spacecraft 
Flight Dynamics, ESA SP-160. 

Rowan-Robinson, M., el al., 1984, Ap. J. (Lett.), 278, L7. 

Stassinopoulos, E.G., 1978, NASA Technical Report X-601-79-7. 

Thomas, G.R., 1983, Journal British Interplanetary Soc., 36, 38. 
van Holtz, R.C., 1983, Journal British Interplanetary’ Soc. 36, 6. 

Wolff, D.M., 1984, American Institute of Aeronautics and Astronautics 22nd Aerospace Meeting, AIAA- 
84-0 150. 


HI-23 


I Hi 



IV. IN-FLIGHT TESTS 


During the mission, a number of in-flight tests were conducted to verify or understand those aspects of 
the performance of the instrument which could not be estimated with sufficient accuracy before the 
flight. These tests are described in this chapter. 

A. Detector/Focal Plane Performance 

A. 1 Detector Sensitivity and Responsivitv 

Although the performance of the individual detectors was quite uniform during the course of the 
mission, there was a range of sensitivities within each wavelength band. The noise equivalent flux den- 
sity (NEFD) was calculated for each operating detector to quantify this spread. Five-minute long seg- 
ments of data taken at high galactic latitude away from regions of obvious infrared cirrus were used to 
calculate the NEFD of each detector. After a baseline was subtracted from the data, a Gaussian noise 
estimator that discriminated against point sources was used to estimate the 1 a rms noise in a single data 
sample. NEFD’s from six representative SOPs were averaged together to give a single estimate of the 
noise in each detector. The NEFD for a given detector varied by less than 25 % for the data examined. 

The results of this analysis are shown in the histograms of Fig. IV.A.la-d which give the NEFD’s of 
each detector. A mean noise for the band was calculated by leaving out those few detectors that were 
significantly noisier than their siblings. It should be pointed out that since a point source contributes 


- 

— 

I 1 1 1 1 1 

12 /tm 

54 


FAILED DETECTORS: NONE 

53 


NOISY DETECTORS: 25, 26, 28 

52 

48 

MEAN NEFD: 105 ±25 mjy 

51 

30 

50 

29 


49 

27 


47 

26 


24 

23 

i 28 1 -i. ■ i r~g 


75 


100 125 150 175 200 225 250 275 


- 

— 

— 

i 1 1 I — l 

25 




FAILED DETECTORS: 17, 20 




NOISY DETECTORS: 42 



44 

MEAN NEFD: 125 ±15 mjy 



43 


45 

40 



41 

22 


46 

21 

18 


39 

19 

16 

i i m 


125 


175 200 225 


10 


SC 8 


£ 6 


i i 

38 


T 1 1 1 


37 


60 fj. m 


35 

FAILED DETECTORS: 36 


32 

NOISY DETECTORS: NONE 


31 

MEAN NEFD: 170 ±15 mjy 


15 




14 




13 

34 



10 

12 


i l_J!_ 

8 

9 

-H— -L_i_ 


75 100 125 150 


175 200 

NEFD (mjy) 


225 250 275 300 


T 


T“ 


250 

I 


275 300 


62 


59 


2 


100 

FAILED DETECTORS: NONE 
NOISY DETECTORS: NONE 
MEAN NEFD: 580 ±110 mjy 


61 

58 


57 


55 


56 


60 


1 


400 500 600 


700 800 

NEFD (mjy) 


1200 


Figure IV.A. 1 


A histogram of noise equivalent flux densities under quiescent conditions. The detec- 
tor number is indicated in each box. The absolute calibration is discussed in Chapter 
VI. The "noisy" detectors were excluded from the means. 





to three data samples, two of which have weights of 0.5, the noise in the bandwidth appropriate to a 
point source is approximately VO smaller than the values quoted in the figure. As discussed in Chapter 
V.A.3.C, however, the signal-to-noise ratio quoted throughout the data processing is based on the single 

sample noise. 

The absolute calibration procedure used to derive these numbers is described in Chapter VI. Degra- 
dations of the sensitivity due to high photon backgrounds, to electron hits on the detectors m the polar 
horns of the Van Allen belts, to cosmic rays and to proton hits in the South Atlantic Anomaly (the SAA) 

are discussed below. 

The detector responsivity and sensitivity depend on the background photon and particle environ- 
ment The responsivity as a function of time was determined by comparison with flashes from the inter- 
nal reference source which was shown to be stable to better than 2% (Section IV.D). Figure IV.A.2 
shows histograms of the distribution of the responsivity of each detector, normalized to the mean of the 
entire mission, as measured throughout the flight. The small intrinsic dispersion of the responsivities in 



Figure IV.A.2 


Histogram of the distribution of the uncorrected responsivities of the individual detec- 
tors for each wavelength band. The responsivities were found from the amplitudes of 
the response relative to flashes from the internal reference source obtained at the start 
and end of each survey scan. Responses to the latter under stable conditions had a 
dispersion of less than 2%; see Section IV .D. 



the shorter wavelength bands is evidence that the changes in the uncorrected sensitivity from scan to scan 
were not extreme. 

A.2 Detector Reliabilitv/Anomalies 

Throughout the mission the performance of the infrared detectors was very stable. Most detectors 
exhibited their pre-launch behavior. Detectors 17, 20 and 36 remained dead. Power spectra of detector 
data streams revealed that many of the detectors, especially the 12 pm and 25 pm detectors, were subject 
to low level 1 Hz electronic cross talk from the temperature sensors in the focal plane. Only for detectors 
19 and 43 did the cross talk exceed the rms noise. In addition, detector 5 (100 pm) was subject to 0.25 
Hz cross talk from an engineering data multiplexer. These three detectore had shown no excess noise 
before launch. As described in Section VI.A.4 it proved possible to either remove (in the case of detector 
5), or to greatly reduce (detectors 19 and 43) the detrimental effects with the ground software. 

Detector 26 exhibited a factor of three more noise than normal in the periods 1983 February 2-10 
(SOPs 17-33) and 1983 March 16-June 6 (SOPs 101-265) for unknown reasons. Detectors 25 and 42 
were generally a factor of two to three times noisier than other detectors in their bands. Detector 28 
showed an abnormal cross-scan response as discussed below. These detectore were sufficiently noisy to be 
declared "failed" in the processing as discussed in Section V.D.2.d. 

A. 3 Cross-scan Response 

The variations in individual detector responsivity with position across the detector were measured 
by scanning a celestial point source over closely spaced tracks across the focal plane. Four or five 
different cuts across a full width detector could be measured in this way. The results from the scans of 
the planetary nebula NGC6543 are shown in Fig. IV.A.3. 1-3.4. The individual data points have been 
normalized to the peak response for each detector and a certain amount of artistic license was used to 
draw the solid curve representing the cross-scan response. Table IV. A. 1 gives the effective solid angle of 
each detector based on the measured cross-scan and in-scan response. 

Detector 28 had abnormal cross-scan response and was declared failed for seconds-confirmation 
purposes (Section VII.D). 

A.4 Verification of Linearity 

The responsivity of infrared detectors can depend on their frequency response and on the total 
amount of infrared radiation falling on them. In-flight tests were conducted to investigate the importance 

of these effects Necessarily, the two types of tests were often coupled and the results were not always 
unambiguous. 

The effect of the total photon flux on the responsivity was measured by repeatedly observing 
asteroids as they approached the lunar limb to within 3°, making use of the out-of-field stray radiation 
from the moon. The highest background levels reached by this technique were 5, 4, 9 and 34 times the 
zodiacal background in the ecliptic plane. The test was thus overly severe in the 12 and 25 pm bands 
where the flux in the zodiacal bands represented the maximum background. The tests were adequate at 
the longer wavelengths where the background flux in the Galactic plane exceeds that in the ecliptic. The 

internal reference source was flashed immediately before and after each observation at the same back- 
ground level. 


IV-3 









Figure IV.A.3.3 Same as Figure IV.A.3.1, except for 60 pm. 





Figure IV.A.3.4 Same as Figure IV.A.3. 1, except for 100 pm. 



Table IV.A.1 Detector Data Based on NGC6543 Scans 


Band 

Det # 

Solid Angle* 
[10 — 7 sr] 

Cross-Scan** 

Dispersion 

% 

Solid Angle 
Det# 

Cross-Scan 
[10 — 7 sr] 

Dispersion 

% 


1 

14.5 

9 

55 

7.1 

11 


2 

12.7 

9 

56 

14.0 

9 


3 

13.0 

10 

57 

13.2 

8 


4 

11.53 

13 

58 

11.2 

15 

100 pm 

5 

12.0 

11 

59 

11.7 

14 


6 

12.4 

12 

60 

13.3 

11 


7 

12.6 

10 

61 

13.5 

11 





62 

10.6 

10 


8 

” 7.2 

9 

31 

2.1 

9 


9 

6.7 

9 

32 

6.4 

9 


10 

6.6 

10 

33 

5.9 

9 


11 

2.8 

9 

34 

6.5 

12 

60pm 

12 

4.3 

9 

35 

6.3 

13 


13 

6.6 

14 

36 

— 

" 


14 

6.1 

12 

37 

6.6 

11 


15 

6.2 

10 

38 

3.9 

14 


16 

3.5 

4 

39 

1.4 

7 


17 

— 

— 

40 

3.1 

7 


18 

3.6 

7 

41 

3.1 

6 


19 

2.8 

4 

42 

3.4 

6 

25pm 

20 



— 

43 

3.2 

6 


21 

2.8 

12 

44 

3.2 

6 


22 

3.1 

9 

45 

3.2 

7 





46 

2.4 

7 

■ggjgg 

23 

2.9 

7 

47 

0.77 

4 

I 

24 

3.0 

4 

48 

3.1 

6 

1 

25 

3.2 

4 

49 

2.9 

6 

■ 

26 

1.2 

9 

50 

3.0 

7 

12pm 

27 

2.0 

8 

51 

2.7 

6 


28 

3.1 

37 

52 

2.5 

7 


29 

2.5 

22 

53 

2.8 

7 


30 

2.8 

10 

54 

2.0 

8 

i — 

* Solid angle based on the measured detector cross-scan response. In-scan response based on average 


point source detector template. 

“ Cross-scan dispersion is the uncertainty in the flux assigned to a single detection due to the fact that 
the detector cross-scan response is not uniform, but the source crosses the detector with uniform pro- 
bability in the non-overlap region. 


IV-8 


1 111 



The frequency response was measured by scanning a given source at rates varying from the nominal 
scan rate (3.85 s ) to 1/16 the nominal scan rate. In addition, the internal reference source was 
observed for varying lengths of time up to 120 seconds. Finally, selected sources were observed in the 
"stare" mode with one detector being positioned on the source for up to 120 seconds; this procedure was 
never successfully executed using a 25 pm detector. 

In Fig. IV.A.4.1,.2 the results of the tests to measure the dependence of the responsivity with fre- 
quency are given for the four IRAS bands. The ratio of the responsivity at nominal survey scan speed 

to that at "DC" was adopted as 0.78, 0.82, 0.92 and 1.0 at 12, 25, 60 and 100 pm; these ratios are indi- 
cated in the figures. 



10 20 
DWELL TIME, SECONDS 


■ 1/2 SURVEY RATE CALIB 0BS aLYR DET 29 
□ STARING OBSERVATION aLYR DET 29 
• 1, 1/2, 1/4, 1/8, 1/16 SURVEY RATE CALIB 0BS 
MEAN OF ALL SOURCES DET 53 a LYR 
A REFERENCE SOURCE LONG DURATION FUSH - MEAN 
OF caOR BAND NORMALIZED AT 3.2 SEC 



10.0 20.0 
DWELL TIME, SECONDS 

• 1, 1/2, 1/4, 1/8, 1/16 SURVEY RATE CALIB 0BS 
FLORA, a LYR DET 45 

AREFERENCE SOURCE LONG DURATION FLASH-MEAN OF 
COLOR BAND NORMALIZED AT 3.2 SEC 


30.0 


Figure IV.A.4. 1 Measurements of the response vs. dwell time to measure frequency dependence of the 
detectors at 12 (left panel) and 25 (right panel) pm. The measurements were made ei- 
ther by crossing a source at scan rates less than the survey rates or by viewing long 
flashes of the internal reference source. 


The tests also show that, at the 5% level, there was, at 12 and 25 pm, no effect of the source 
strength on the dependence of the responsivity with frequency. The "mean" observations in the 
figures represent stars whose amplitudes span more than a factor of 100 in brightness and which show no 
significant departure from the curves shown. At 60 and 100 pm the situation is clearly different. Tests of 
sources up to 10 times brighter than a Lyr show the same behavior as does a Lyr. Stronger sources show 
a variety of behaviors as indicated in the figures. 

Of particular interest is the frequency dependence between observations at the survey rate and 
observations taken at half survey rate since the latter rate was used in pointed observations, some of 
which were crucial in the absolute calibration procedure. In Fig. IV.A.5 the magnitude differences 

between the observations at the survey rate and at 1/2 survey rate are given for sources of varying 
strengths. 


IV-9 






+ 1 SURVEY RATE OBS DET 14 
• 1, 1/2, 1/2, 1/8, 1/16 SURVEY RATE DET 37 
■ 1/2 SURVEY RATE CALIB OBS DET 37 
▲ 1/8 SURVEY RATE DET 14 
□ STARING OBS DET 14 
A STARING OBS DET 37 


60 jim 


■ l/2x SURVEY RATE OBS DET 14 
□ STARING OBS DET 14 




• 1 SURVEY RATE OBS DET 6 
■ 1/2 SURVEY RATE OBS DET 6 
a 1/8 SURVEY RATE OBS DET 6 
□ STARING OBS DET 6 


TOO //m 


Measurements of the response vs. dwell time to measure frequency dependence 
of the detectors at 60 (top panels) and 100 (bottom )pm. The measurements 
were made either by crossing a source at scan rates less than the survey rates or 
by viewing long flashes of the internal reference source. 


IV- 10 










Figure IV.A.5 


Measurements at survey and 1/2 survey scan speeds. The magnitude scale is defined 
in Section VT.C.2.a. The source IRC+10216 is at the extreme left in both panels. 


A.5 Baseline Stability 

The electronic baseline stability proved to be quite good on a time scale of a day, with drifts typi- 
cally less than 5% of the sky brightness toward the north ecliptic pole over this period at 12 and 25 pm 
and less than 20% at 60 and 100 pm. Throughout the mission, baseline drifts over periods longer than 
about one day were monitored by daily observations of a region of sky near the north ecliptic pole which 
was called the Total Flux Photometric Reference (TFPR). A detailed discussion of the determination of 
the brightness of this region and its time variation is given in Section VI.B.3. Additional differential 
effects between detectors which were important in the extended emission maps were removed in the de- 
striping processor described in Section V.G.6. 

A.6 Particle Radiation Effects 

In order to minimize the expected degradation of the data quality due to energetic particle radia- 
tion, the IRAS hardware incorporated a number of features including nuclear shielding, radiation hit 
deglitchers, radiation hit deadtime counters and bias boost circuits; see Section n.C. In addition, opera- 
tional procedures were developed during the in-orbit checkout phase to minimize radiation effects. These 
procedures were incorporated into the routine mission procedures (see Section III.B.6 and III.C.4). 


IV- 11 


A 6 a Detector Responsivity and Noise 

The radiation effects from the horns of the van Allen belts were basically limited to an increase, 
typically by a factor of two or less, in noise due to many small pulses not eliminated by the deghtcher ctr- 
cuitry. P assage through the SAA caused large changes to the detector responsivity particularly in the 60 
and 100 um bands. These were monitored by comparisons with flashes from the internal reference 
source after passage through the SAA. Cosmic ray hits at the rate of about one per twenty seconds on 
each detector were handled adequately by the nuclear spike deglitchers and caused little degradation to 

the data (Section VII.D). 

A h Radiation Effects on Detector Baselines 

Ionizing radiation in the IRAS environment affected the baseline stability in two ways. Residual 
noise spikes from polar horn passage and entry into the edges of the SAA were not removed from the 
data and when added into the extended emission maps resulted in elevation of the baseline and increased 
note. The radiation level was monitored by using the activation of the nuclear particle circumvention 
circuit as described in Section III.D.2, and data taken during times of high radiation levels (blanking time 
greater than 10%) were excluded from the extended emission data base. 

A.7 Effects of Bias Boost 

The bias boost (Sections II.C.5 and III.B.6) applied to the 60 and 100 pm detectors to minimize 
radiation exposure effects during deep SAA passages reduced the responsivity and noise as expected from 
the pre-flight measurements. The change in the responsivity of the 25 pm detectors was sufficiently sma 
that no bias boost in that band was regularly applied. Figure IV.A.6 shows the mean response to the 
internal reference source from the detectors of each of the four IRAS bands as function of time before 
and after two consecutive SAA crossings of more than average radiation dosage using the bias boost pro- 
cedures developed during in-orbit checkout. It can be seen that the responsivity was stable to within 5%- 
10% in the 12 and 100 pm bands, about 1 5% in the 25 pm band and better than 5% in the 60 pm band. 

The bias boost also resulted in a baseline shift significantly larger than described in Section IV.A.5 
which decayed exponentially with time. The details of this behavior were measured by special observa- 
tions and analysis; a complete description is given in Section VI.A.3. There were times when the track of 
the satellite went near the boundary of the SAA, but when no bias boost was applied. If the internal 
reference source was activated near the SAA in these cases, errors in the responsivity as large as 8%, 
which resemble baseline errors, affected all detectors in the band. These produced increased apparent 
brightness at 12, 25 and 60 pm and decreased apparent brightness at 100 pm. 

A.8 Photon Induced Responsiv ity Enhancement 

After the mission was completed, a comparison of scans which crossed the Galactic plane in an 
ascending manner with those crossing it in a descending manner showed an enhancement m the respon- 
sivity in the 100 pm detectors due to passage through the Galactic plane. 

Subsequent analysis of special calibration scans over Saturn confirmed that this effect was due to 
responsivity enhancement caused by infrared photons. During these special observations, Saturn was 
scanned across the focal plane ten times with an integrated photon dose on the 100 pm detectors ranging 
up to 12 x 10 -10 Joules «T 2 . The enhancement associated with this photon dose was defined as the ratio 


IV-12 




Figure IV.A.6 Effect of bias boost on responsivity after passage through SAA. The responsivity was 
measured by comparison with repeated flashes of the internal reference source. For 
clarity, the 25, 60 and 100 pm observations have been shifted by arbitrary amounts. 


of the amplitude of the internal reference source after the observation to the amplitude of the flash prior 
to the observation. The results are shown in Fig. IV.A.7. The 100 pm detectors 3, 6, and 7 show little 
enhancement while the rest of the 100 pm detectors show an enhancement increasing with increasing 
dose to about 13% enhancement at about 9 x 10 -10 Joules m -2 . 

Figure IV.A.8 illustrates how the response to the internal reference source decayed with time after a 
Saturn exposure. Each amplitude of the flash from the internal reference source has been normalized to 

the amplitude from the source flash that preceded the Saturn exposure. This was done for three observa- 
tions following a Saturn observation. Care was taken to ensure no bias boost occurred between these 
observations. Of course, there is no guarantee that the response of the detectors to the internal reference 
source was not affected by photon exposure within these post-Satum observations. 

Figure IV. A. 9 shows the integrated 100 pm photon dose due to the Galactic plane emission as a 
function of ecliptic longitude for a nominal scan inclination of 45° with respect to the Galactic plane. 
The integrated photon dose is generally in the range of dosages encountered in the Saturn observations 
with peak dosage near the Galactic center being approximately twice the peak Saturn dosage. No 
significant enhancement was observed for the 12, 25 or 60 pm detectors in the Saturn observations. 


IV- 13 



Figure IV.A.7 


Observations of photon induced responsivity enhancement in the 100 pm detectors 
after scans of Saturn. The response for the three detectors #3, 6 and 7 are plotted 
separately; the responses of the other 100 pm detectors are averaged. 



TIME (sec) 


Observations of the decay of photon induced responsivity enhancement after scans 
over Saturn. The response was measured using flashes of the internal reference source 
and is normalized to the flash amplitude which just preceded the scan over Saturn. 
The Saturn crossing occurred between 0 and 300 seconds. 


1 I 


IV- 14 






Figure IV.A.9 Typical photon dosages in the 100 pm detectors for 45° crossings of the Galactic 
plane. The dosages were calculated from the mean surface brightness within 20° of 
the Galactic plane in 0.5° intervals. 


The statistical analysis of the elfect of photon induced responsivity enhancement cause by passage 
over the Galactic plane is discussed in Section VI.B.4.C. 

A.9 Fee dback Resistor Nonlinearity Analysis 

The impedance of the detector feedback resistors, nominally 2 x IO ,0 ii, decreases with increasing 
voltage as shown in Fig. II.C.2. The shape of the resistance versus voltage relation for the feedback resis- 
tors was verified by (a) comparing fluxes for stars measured at 12 and 25 pm with ground-based observa- 

rorlim°i< Un c 8a ’ 1984 ’ RlCke Ct aL ’ 1984) 3nd (b) Special ,0W gain calibration observations of 
IRC+10216. Since IRC +10216 is known to be variable, the latter were compared with nearly simul- 
taneous 10 pm ground-based observations (Joyce, 1984). The results for the 12 and 25 pm bands are 
shown in Fig. IV. A. 10. In this test the 25 pm comparison value for IRC+10216 was estimated from the 
published spectral distribution (Campbell « al, 1976) normalized to the 12 pm measurement. Detector- 
to detector differences in the Rf shape were assessed by comparison of detector response ratios There 
appears to be no significant deviation from the adopted R f vs. voltage curve over the range examined 

Furthermore, there appeal, to be no source strength dependent term larger than a few percent in the 12 
and 25 Jim bands. 

Sl ”" ar teStS f0r Ihe 60 and 100 bands were not carried out duc 10 uncertainty in the 60 and 
100 tint flux from IRC+102,6 at the time of the IRAS observation and the lack of bright, pointlike iso- 


IV-15 



(GND - 12 ^m) ( GND " 25/xm) 



Figure IV. A. 10 


Comoarisons of 12 and 25 pm IRAS pointed observations with ground based obser- 
STover a wide range of magnitudes (Section VI.C.2.a) to check nonhneanty of 
the feedback^ resistor The ground based observations of standard stars are from 
ttek Tf al (1984) (open circles) and from Tokunaga (1984) (closed circks). The 
offsets reflect a zero point difference of 0.02 mag used by Tokunaga and the fact that 
both ground based systems adopt zero c £ lor ^ iffere ^^ 

whereas a color difference of 0.03 mag has been adopted for the IRAS cahbration. 
See Section VI.C. The ground observations of IRC +10216 are by Joyce (198 , 
Section IV.A.9). 


lated sources well observed with other telescopes which were also observed with IRAS calibration observe- 
tions. 

B. Spectral Passband Verification 

The product of the transmission efficiency of the optical filter stack and the detector efficency 
defines the effective wavelength passbands and the nominal out-of-band rejection performan 
detectors This is described in Section II.C.5 and Tables 11.C.4 and H.C.5. The detector passbands were 
verified before launch at the component level, but not at the system level. In-flight tests were perforate 
to test the consistency of the passbands and out-of-band rejection within each band and to verify the 
There was no dime! in-orbi, veHfication of the effective transmission charactenst.es 

of the detector/filter combination. 


IV- 16 


B.l. Verification of the Relative Consistency 

The consistency of the transmission characteristics within each band was checked by a two step 
procedure. First, all detectors were carefully calibrated relative to NGC 6543, which, for the purposes of 
this test, may be considered as a 150 K blackbody with spectral lines. The second step was to measure 

with each detector a number of sources with much hotter or colder spectral flux density distributions 
than NGC6543. 

Numerically, the approach was to measure, in each wavelength band, the flux S tJ from test source i 
on detector j. The average of Sjj over the detectors in the wavelength band, (S^, then gives the best 
estimate of the flux from the source in that wavelength band. The statistical distribution of the ratio 

) (IV.B.l) 

has, by definition, a mean of one if all detectors respond equally to the test source and NGC6543. The 
test sources (Table IV.B.l) were four stars and two cold sources, Neptune and the galaxy UGC1 1348. 

The stars probe the short wavelength rejection and the short wavelength transmission cuton of the 
detectors in all bands. The two 50 K sources probe the longwave cutoff and longwave rejection in the 1 2 
and 25 pm bands and the inband transmission of the 60 and 100 pm bands. Significant changes in the 
band pass characteristics of an individual detector relative to the band average transmission would be 
detectable as a statistically abnormal member of the /?, , data set 

The results of this test indicate that deviations in the flux estimates from the mean are typically less 
than 10%, with a worst case value of 16%. Table IV.B.2 gives the typical and worst case value of the 
absolute deviation of Rjj from unity in each band for the stars and the cold sources. In spite of their 
wide temperature range, there was no difference between the hottest and the coolest star. 

Verification of the Nominal Inband/Out-of-band Transmission 

The nominal response of the filter/detector combination and results of preflight tests at the com- 
ponent level are summarized in Section II.C.4. In-flight verification was required for two reasons. First, it 
was not possible to test fully the detector/filter combinations before the flight even at the component 
level. This applies in particular to the spectral response of the 100 pm detectors. Furthermore, a number 
of problems could have resulted in a degradation during system integration or during the flight. 

Potential problems with the spectral response can be divided into two categories, those that 
represent systematic deviations from the nominal characteristics of all detectors in a band and those that 
result in a random scattering in one or several detectors in a band from the band average. Problems 
likely to result in random scattering are cracking or delamination of the filters, photon leaks around filter 
mounts etc. Extensive in-flight tests, described in the previous section, indicate that the detectors had 
rather similar pass-bands, thus eliminating random, but not necessarily systematic, effects. 

An attempt was made to address the problem of systematic deviations of the nominal inband and 
out-of-band transmission of the four bands in a semi-quantitative way, by using observations of stars with 
widely different temperatures, asteroids and the planet Uranus to probe different portions of the 
passband. 


IV- 17 


Table IV.B.l 

Out-of-Band Rejection Test Sources 



Hot test source 

Cold test source 

a Vir, 

B1 IV 

Te-24000 K 

Neptune Te“55 K 

a CMa, 

A1 V 

Te- 10000 K 


a Car, 

F2 II 

Te-6900 K 

UGC 11348 Te— 50 K 

P Peg, 

M2.5 

Te-2800 K 


1 




— 1 


Table IV.B.2 

Out-of-Band Rejection Test Results 




Wu 

- Dl 


Band 


Stars 

Planet/Galaxy 


pm 

Typical 

Worst Case 

Typical Worst Case 


12 

0.02 

0.09 

* * 


25 

0.02 

0.07 

0.05 0.14 


60 

0.03 

0.09 

0.06 0.12 


100 

0.10# 

0.13 # 

0.06 0.16 


l — 

* not detected 

# 

low signal-to-noise ratio 


The following hardware related issues affect the various bands: 

a. 12 pm band: Rejection of 15-23 pm depends on an interference filter. A weak upper limit on the 

mean transmission in this region is based on the non-detection of Uranus in the 12 pm band. 

b. 60 pm band: Rejection shortward of 27 pm depends on a combination of blocking filters and the 

Ge:Ga detector response. Very high short wave blocking is required in the case of stars. Based on the 
semi-quantitative analysis of observations of hot stars, cool stars and asteroids, the shortwave blocking 
is consistent with the design specifications given in Section II.C.4. 

c. 100 pm band: Rejection shortward of 50 pm depends on a combination of blocking filters and the 

Ge:Ga detector response. The comments in b above apply. The longwave cutoff is determined by the 


IV- 18 




detector cutoff at approximately 120 pm. The detector response, including this cutoff, was not meas- 
ured on the actual flight detectors and the cutoff may vary a few microns from one detector to 
another. 


C. Optical Performance 

Diffraction and scattering of the infrared radiation from bright sources from optical surfaces or tele- 
scope structures result in potential artifacts in the IRAS data. Early in the mission an effort was made to 
assess the magnitude of such effects using specially designed observations of Jupiter and Saturn. Observa- 
tional procedures (discussed in Chapter III) and software (Section V.D.2.c) were used to eliminate most 
of the artifacts from routine observations. 



Optical cross talk is here defined to mean the detection of the flux from a source on a detector that 
is inconsistent with the reconstructed position of the image of that detector and the position of the source 

on the sky. This is distinct from electronic cross talk in that the detected signal is in fact due to infrared 
radiation incident on the detector. 

The fact that the secondary mirror is supported by three spider arms resulted in a diffraction pattern 
consisting of an Airy disk and six diffraction spikes equally spaced on a circle. Two of the diffraction 
spikes are aligned with the scan direction. The diffraction spikes constitute a minor component of the 

point source diffraction pattern, but the point source detection algorithm (see Section V.C) is quite sensi- 
tive to the diffraction spikes. 



Bright sources not directly imaged onto the focal plane could produce apparent infrared signals 
which, if care were not taken, could be confused with real infrared sources. The attenuation of such 
signals as a function of angular distance from the telescope boresight is defined as the out-of-field rejec- 
tion ratio, with the nominal specifications given in Section II.C.3. A number of tests were carried out to 
verify the nominal out-of-field performance and to check for unexpected glints. These tests confirmed 
the anticipated need for operational procedures not to make routine observations closer than 1° from 
Jupiter, 20° from the moon, 60° from the Sun and 88° from the Earth horizon. Estimates of the in-orbit 
out-of-field rejection are included in the discussion in Section II.C.3. 


Cross talk from Jupiter appeared to be associated with diffraction from the secondary mirror sup- 
port structure. The amplitude of the spikes and associated diffuse cross talk components was generally 
less than that predicted from simple telescope models by about a factor of three. 


The moon tests revealed significant "glints", in addition to the anticipated amount of diffuse out-of- 
field radiation. Glints are well defined regions of about 1/2° extent, where the telescope out-of-field per- 
formance is significantly reduced. The glints are of two types, one of which propagates across the focal 
plane in the survey scan direction affecting all detectors along the scan and the other which affects only 
specific detectors around the edge of the array. Figure IV.C.l shows a map of all presently known 

moon glints out to 30 from the moon. It should be noted that routine observations were carried out as 
close as 20° from the moon. 


IV- 19 



Figure IV.C.l 


A polar plot showing where glints from the moon were delected by the ^numbered 
detectors The boresight of the telescope is at the center. The moon is at the cone 
and azimuth angle plotted with the tracks. The angle 6 gives the angle to the sun. 


C 3 Out of Field Rejection M onitoring: 

' ' Due to concern about possible slow degradation of the ont-of- held performance during the mission. 
Ore out-of-field performance was monitored throughout the mission using two observattonal procedures. 

The first was a test using the moon as a bright source. Approximately once per ”° nth ’ 
moon about 90‘ from the Sun. the spacecraft made 60* long scans which passed within 20 the . 
These scans were compared with reference scans over (roughly) the same stnp of the sky preceding (or 
r U Twi"lunm 1 b, one day, such that the moon was mughly 32- from the boresight a, the 
closes, approach. The difference in the flux for the same point on die sky measured in the two scans was 
"be due to out-of-field radiation from the moon. The results of this monitoring program are 
shown in Fig. IV.C.2. No significant changes in the out-of-field performance were observed during the 
ZZ of Z mission and the deduced mean value of die ou,-of-fie.d tadiadon from the moon a, 20 
from the boresight is in good agreement with expected performance. 

The second approach was to attempt to monitor and/or detect out-of-lield radiation from the Earth, 
due to radiation from the Earth scadered off the inner surface of die sunshade, followed by d.ffmct.on 
and/or scadering through the baffle structures. The technique employed was to observe several region 


I I 


IV-20 



£ 

Ll_ 

I— 

z 

< 

LU 

5 


o 

Csl 

y— 

< 

LU 

u 

z 

< 

Q 

< 

O' 

ex' 


LT) 



I 


o 

ex' 

< 

z 

3 


0.3 

— 

i 


T 1 1 

12/im BAND 

i 

— 

0.2 
0. 1 

— 

1 

(X) 

X x 
(X) 

! 

— 

0.3 

— 

I 


1 1 1 

25/im BAND 




0.2 

— 





— 

0. 1 

— — 

X-#- 

1 

(X) 

X 

,(X) ^ 

tt — 

1 


0.3 

— 

t 


1 1 1 

60 BAND 

1 

— 

0.2 
0. 1 


X 

X 

1 

X 



i 1 1 

(X) 

1 


0.3 

1 

1 


f 1 1 

100 fj.m BAND 

1 

(X) 

1 

— 

0.2 

0.1 

h- 

1 

X 

X 

1 

■~5T" 

! _ 
i 

*! 

i 

k 


( 

) 

100 

200 300 400 

500 

600 


SOP 


Figure IV.C.2 Monitoring of the out-of-field rejection by observations of the moon 20° off axis. The 
observations are normalized to values of the TFPR as are the calculations in Fig. 
II.C.5. The dashed lines show the mean values which are consistent with the pre- 
launch calculations shown in Fig. II.C.5. The bracketed points are of low weight. 
The observations in SOPs 170 and 225 required large corrections for responsivity 
changes in the 25 Jim band. 


the sky, generally twice, when the orientation of the sky was such that the region could be observed while 
no Earth radiation was incident on the inner surface of the sunshade. The observation of the same 
region was then repeated, typically about a month later, when the orientation of the sky was such that the 
inner surface of the sunshade was as nearly fully illuminated as possible without violating the Earth limit 
pointing constraint. The difference in observed flux, after correcting for zodiacal emission differences and 
changes in the baseline attributable to electronics drift is presumably due to out-of-field radiation from 
the Earth. The analysis is limited by the accuracy of the zodiacal emission model and baseline uncertain- 
ties. The upper limits from the tests represented a background less than 0. 1 the flux from the TFPR (see 
Section IV.A.5), but the tests were not sensitive enough to confirm pre-launch calculations of the out-of- 
field rejection (see Fig. II.C.5). 

The extended emission data reduction subsystem further controlled the effects of out-of-field radia- 
tion by rejecting data taken with larger avoidance cones around bright objects. Tests early in the mission 
indicated that the moon was the only object bright enough to warrant further avoidance and any data 
taken within 30° of the moon was ignored. A more detailed discussion of avoidance angles is given in 
Section III.C.5. 


IV-21 




(IRS/N6543 - MEAN) /MEAN 


D. Interna] Reference Source Stability 

During the course of the entire mission, the stability of the internal reference source was monitored 
at least daily relative to the planetary nebula NGC6543. The results of this monitoring program at 25 
pm for a large fraction of the mission are shown in Fig. IV.D.l. Let R i<t be, for detector i and time t, 
the ratio of the amplitudes from a flash of the internal reference source and from a scan over NGC6543. 
If (/?,,,} is the mean ratio averaged over the duration of the mission, the quantity 

is a direct measure of the stability of the internal reference source since the infrared output from 
NGC6543 can be presumed to be stable. The quantity Y i>t is shown as the ordinate in Fig. IV.D. 1 for all 
full-size detectors in the 25 pm band as a function of time from the start to the end of the mission. This 
wavelength band was selected for presentation since the signal-to-noise ratios for both the stimulator 
flashes and NGC6543 were the highest in this band. The dispersion in the measurements at 25 pm is 
1 .6%, consistent with the dispersion due to the measurement uncertainties. The measurements in the 1 2, 
60 and 100 pm bands give dispersions of 2.6, 1.7 and 2.7% respectively. The stability of the internal 
reference source is thus better than 2% throughout the mission. 


0.20 


0.00 


- 0.20 


25pm BAND 


; Lf ' 

j * ' * * * . * * * * 




JL 


JL 


_L 


-L 


_L 


_L 


_L 


1.00 60.90 120.80 180.70 240.60 300.50 360.40 420.30 480.20 540.10 600.00 

SOP 


Figure IV.D. 1 The response of the individual detectors to the internal reference source compared to 

that from the planetary nebula NGC6543 at 25 pm. The results from all the 25 pm 
detectors were combined after normalization to the mean response over the duration 
of the mission. 


IV-22 


1 I! 



Authors: 


H. H. Aumann, T. N. Gautier, F. C. Gillett, P. Hacking, G. Neugebauer, F. Olnon, S. Wheelock 


References 

Campbell, M. F„ Elias, J. H„ Gezari, D. Y., Harvey, P. M„ Hoffmann, W. F„ Hudson, H. S 
gebauer, G., Soifer, B. T., Werner, M. W., and Westbrook, W. E. 1976, Ap. J. 208, 396. 

Joyce, R. 1984, private communication. 

Rieke, G. H„ Lebofsky, M. and Low, F. J. 1984, preprint. 

Tokunaga, A. 1984, A.J. 89, 172. 


, Neu- 




IV-23 



V. DATA REDUCTION 


This chapter describes the transformation of the raw data into its final forms. Section V.A provides 
a general overview of the data reduction, while subsequent sections provide a more detailed description. 
The relative and absolute calibration of the data is discussed separately in Chapter VI. The results of the 
analysis of the data processing are given in Chapter VII. An overall description of the survey, including 
sky coverage, completeness and reliability, is given in Chapter VIII. 

A. Overview 
A.l General 

The Science Data Analysis System (SDAS) at JPL received the scientific and housekeeping data 
taken by the satellite from the ground operations facility at Rutherford Appleton Laboratory in Chilton, 
England. The data were obtained and processed in units of a Satellite Operations Plan (SOP) containing 
10-14 hours of data (Section III.C. 1). A single SOP was further divided into 10-60 observations or scans, 
each of 3-50 minutes duration. Survey observations obtained the data for the all-sky survey, while cali- 
bration observations were made to determine, monitor and verify the survey photometry. 

The survey strategy and observation constraints produced scans that were generally along meridians 
of ecliptic longitude. At the ecliptic plane, the scans followed meridians exactly, while at higher latitudes, 
deviations from meridians could become pronounced. The rectangular aspect of the infrared detectors 
(Section II.C.4.) combined with the survey strategy means that there are important differences between 
quantities, such as source size and positional uncertainties, measured in the in-scan direction (roughly 
ecliptic latitude) and those in the cross-scan direction (roughly ecliptic longitude). In extreme cases, 
whether a source is found in the point or small extended source catalogs could depend on whether it was 
extended in the in-scan direction. 

A.2 IRAS Catalogs and Atlases 

The results of the survey are to be found in either the catalogs of point sources or small extended 
sources, or in the catalog of low-resolution spectra or in the atlas of sky brightness images. Where the 
data for a given astronomical object are located depends primarily on its angular size. Sources with 
angular extents less than approximately 0.5, 0.5, 1.0, 2.0’ (in the in-scan direction) at 12, 25, 60 and 100 
pm, respectively, are listed in the point source catalog. As discussed in Chapter IX, low-resolution spec- 
tra in the 8-22 pm region are available for objects brighter than about 10 Jy at 12 or 25 pm. Sources 
with angular extents larger than the above values but less than about 8’ may be found in the catalog of 
small extended sources. Information about sources larger than about 8’ in-scan will be found only in the 
atlas of sky brightness images. The images have a pixel size of 2' and an angular resolution of 4-6'. 

A. 3 Processing Summary 

The following processing steps were applied to the data: 

1. Data reconstruction (Section V.A), 

2. Pointing reconstruction (Section V.B), 

3. Calibration (Chapter VI), 


V-l 


4. Source detection (Section V.C), 

5. Point source confirmation (Section V.D), 

6. Low-resolution spectra extraction (Chapter IX), 

7. Small extended source confirmation (Section V.E), 

8. Identification of possible asteroids and comets (Section V.F), 

9. Data compression for extended source images (Section V.G). 

The point source catalog was generated after all of the data had been processed to create the inter- 
mediate Working Survey Data Base (WSDB) which contained all sources that were hours- or weeks- 
confirmed. The following steps were then applied to sources in the WSDB (Section V.H): 

1 . Final calibration, 

2. Source clean-up and neighbor tagging, 

3. Average flux computation, calculation of flux uncertainty and source variability analysis, 

4. High source density region clean-up, 

5. Source selection, 

6. Association and classification of low-resolution spectra (Chapter IX). 

7. Positional associations with other astronomical catalogs, 

8. Cirrus flagging, 

9. Transformation of coordinates, assignment of source names and calculation of positional 
uncertainty ellipses. 

To generate the small extended source catalog, the following steps were applied to a file containing 
hours-confirmed small extended sources (Section V.E): 

1. Cluster analysis processing, 

2. Weeks- and months-confirmation, 

3. Band-merging, 

4. Final calibration, 

5. Positional associations with other astronomical catalogs. 

To generate the atlas of all-sky images, the following steps were applied to the compressed data 
(Section V.G): 

1. Quality checking and data selection, 

2. Regridding into images in equatorial coordinates, 

3. Final calibration. 

The remainder of this overview briefly describes these basic processing steps, except for calibration 
which is discussed in Chapter VI. A much more detailed discussion of each topic is given in the subse- 
quent sections of this chapter. 


V-2 



A.3.a Data Reconstruction 


The detector data were reconstructed from the telemetry data as discussed in Appendix II. 1 . Any 
one-second frame of detector data containing more than one parity error, indicative of a transmission 
error from the satellite station to the ground, or from Chilton to JPL, was discarded. Frames next to a 
data outage with even one parity error were also discarded. The frequency of such errors was quite low, 
with typically fewer than five seconds of data being rejected from each SOP. Each second of data was 
tagged with its time of observation in seconds since 1981, January 1.0 UT. 

A.3.b Pointing Reconstruction 

Data from Sun sensors, gyros, and observations of bright, m v < 7, visual stars were combined to 
determine a continuous record of the position in the sky that a reference point in the focal plane was 
observing. Infrared observations of stars were not used in the reconstruction process, but were used to 
verify the accuracy of the pointing. Most survey observations contained two or more visual star observa- 
tions, so that in-scan pointing errors were typically less than 5" and cross-scan pointing errors less than 
10 . However, a few survey observations had none or only one visual star observation with resultant in- 
scan errors as large as 1-2'. 

A,3.c Source Detection 

For point sources, a zero-sum square-wave filter was used to search the calibrated data stream from 
each detector for peaks larger than roughly 2.5 times the rms noise. The noise at a given time was 
approximated by a filtered running average of the values of all preceding square-wave peaks. The noise 
estimator was a one-sided estimate of the noise which was in error when the noise changed rapidly. For 
each candidate peak, an 1 1 -point template was fitted to the data, giving the amplitude and crossing time 
of the peak and a correlation coefficient. Those peaks with an amplitude above 3.0 times the noise and a 
correlation coefficient greater than 0.87 were passed to the confirmation processor as valid detections. 
Since the noise refers to a single data point and a point source contributes to three successive data points 
(with weights 0.5, 1.0 and 0.5), the true signal-to-noise threshold of the detection process was roughly 3.7, 
which is significantly below the local signal-to-noise cutoff of about 5.7, implied by a correlation 
coefficient cutoff of 0.87. The noise on a single data point was quoted throughout the processing. 

For small extended sources a simpler algorithm was used to detect potential sources. Zero-sum 
square-wave filters of various angular extents were applied to the data from each detector. The filter sizes 
were 1, 2, 4, 8, and 16 at 12 and 25 pm, 2, 4, 8 and 16' at 60 pm, and 4, 8 and 16' at 100 pm. The 
data streams filtered with the smallest (P at 12 and 25 pm, 2’ at 60 pm and 4’ at 100 pm) and largest 
(16) spatial filters were used to discriminate against point sources and extended sources larger than 8’. 
Detections with a larger flux in one of the extended source filters than in the point source filter were con- 
sidered as possible extended sources and passed to the corresponding confirmation processor, indepen- 
dently of whether a point source detection occurred. 

A.3.d Point Source Confirmation 

The aim of point source confirmation was to glean from the hundreds of thousands of detections 
per day the properties of those sources, and of only those sources, which appeared as inertially fixed on 
the sky. Detections of dust, space debris and asteroids had to be rigorously excluded, yet not at the 


V-3 


expense of celestial objects, if the complementary goals of completeness and reliability were to be 
saisflej T he existence of a point source had to be confirmed on timescales of seconds, hours and weeks 
for the object to be included in the catalog. 

The layout of the focal plane ensured that any stationary source of infrared radiation would cross at 
least two detectors in each wavelength band as the satellite scanned the sky (Fig. II.C.6). The seconds- 
confirmation processor, which treated each wavelength band independently, examined each detection in 
turn in an attempt to find a detection on one of the other detectors in the focal plane. The other detec- 
tion had to occur on one of the two detectors that could be hit by a true source traversing the focal plane, 
at the correct displacement in time, and agree in flux with the first detection to within roughly a factor of 
two. Detections satisfying these criteria were merged to become seconds-confirmed detections. A success- 
ful seconds-confirmation resulted in a refinement of both the flux and position of the object. 

After all of the detections in the four bands for a given survey scan were examined, an attempt was 
made to combine observations in the different bands to produce a seconds-confirmed band-merged 
source. Detections in the different bands had to satisfy tests based on in-scan timing and cross-scan posi- 
tion. The order in which bands were taken in the search for merger candidates proved to be important 
because of the effects of the 100 pm "cirrus". As discussed in Section V.D.3.b, a priority was chosen to 
minimize these deleterious effects. Bands that were not filled by successful mergers were given upper lim- 
its by determining which detectors in those bands were crossed by the source and then taking three times 
the noise on the detector with the lowest noise. The source position was refined during band-merging. 

A file containing roughly 32000 previously known sources including SAO stars, asteroids with well 
determined orbits, and objects from the Two Micron and AFGL Surveys (Neugebauer and Leighton 
1969; Price and Walker 1976) was used to predict positions and fluxes for routine monitoring of the pro- 
cessing. The seconds-confirmed band-merged sources were compared to the predictions, and 
identification numbers were assigned to successful matches. 

The next level of confirmation, hours-confirmation, was run on groups of three successive SOPs in 
at least one band. Every source in the first SOP received its turn to find candidates from successive scans 
in any of the three SOPs. Candidates had to pass a position and flux test (again, roughly a factor of two) 
in at least one band to become an hours-confirmed source. An hours-confirmed source had to have 
either two seconds-confirmed detections, one seconds-confirmed sighting plus one non-seconds-confirmed 
sighting or two non-seconds-confirmed sightings, each with the alibi of a failed detector. A successful 
hours-confirmation resulted in a refined position and a refined flux for each band where possible. If each 
candidate source had only upper limits in a given band, the lowest upper limit was adopted. 

Finally, the third level of confirmation, weeks-confirmation, was attempted for each hours- 
confirmed source prior to its entry into the WSDB by examining the WSDB for any other hours or previ- 
ously weeks-confirmed sources that could pass a position test. Successful matches resulted in a refined 
position for the source, but not a refined flux; individual hours-confirmed fluxes were retained for each 
source It is important to note that weeks-confirmation represents the first and only time interval for 
which at least rough flux constancy was not required. Thus, the catalog is weakly biased against celestial 
sources more variable than a factor of two on time scales shorter than a few days. 


A.3.e Small Extended Source Confirmation 


The objective of the processing of small extended sources was to detect stationary sources which 
were resolved by the telescope, up to 8' in extent. Because of the simple detection algorithm described 
above, it was not intended that the processing produce intensity or flux measurements more accurate 
than about 50%. Nor were stringent goals of either completeness or reliability set. 

Potential small extended source detections were identified in the raw data stream by their response 
to square-wave filters (Section V.A.3.c). Each detection was tested to discriminate against point sources 
and sources larger than 8'. After individual source detections were identified, they were tested for 
seconds-confirmation; those that failed were discarded. Detections which were seconds-confirmed, which 
were from the same hours-confirming survey coverage, and which were sufficiently close together on the 
sky, were combined to produce a single source. These small extended sources became the input for all 
subsequent processing. No attempt was made to combine observations at different wavelengths until the 
last stages of the processing. 

Structure on the scales of 2 to 10 f within much larger extended sources gave rise to clusters of small 
extended sources. Such clusters were eliminated so that only those sources whose sizes were compatible 
with the square-wave filter used for detection were accepted. Sources were then tested for repeatability in 
position and flux on time scales of weeks and months. Any source which was not weeks-confirmed was 
discarded. The final stage of the processing was to combine those weeks-confirmed sources observed at 
different wavelengths which were sufficiently close together on the sky, and to merge them into a multi- 
band source. 

A.3.f Asteroids and Comets 

Asteroids are a population of bright infrared sources, particularly at 12 and 25 pm. The hours- 
and weeks-confirmation strategy was developed to discriminate against these moving sources. Positions 
of known asteroids were calculated and associated with those of hours-confirmed point sources. In this 
manner the known asteroids were identified and used as tracers to evaluate the effectiveness of the 
confirmation filters. No known asteroids satisfied the stringent position coincidence requirements to 
become final catalog sources. As discussed in Section VII.F, confusion with asteroids (known and previ- 
ously unknown) was the cause of about 100 unreliable sources and fluxes within 25° of the ecliptic plane. 

To provide data for the study of the properties of known and newly discovered asteroids, all sources 
with infrared colors typical of solar system objects were written to auxiliary files at both seconds- and 
hours-confirmation. The emphasis was on completeness. The results of analyzing these data are not 
included in any of the IRAS catalogs released in 1984. 

A.3.g Extended Emission Processing 

The goal of the extended emission processing was to produce moderate resolution maps of the total 
infrared emission over the whole sky at 12, 25, 60 and 100 pm. Special procedures discussed in 
Chapter VI were used to calibrate the absolute E>C level of the data from each detector. These data were 
smoothed and resampled (compressed) in the time domain to produce two samples per second. At the 
nominal scan rate of 3.85 s 1 this sampling closely matched the 2 f spacing of detectors across the focal 
plane. The focal plane geometry and the reconstructed pointing information were used to locate the 


V-5 


position on the sky of each data sample for projection into the 212 images, each 16.5 x 16.5°, that cover 
the whole sky with 2x2 pixels. 

Data suspected of being either noisy due to particle radiation events or contaminated by scattered 
tight from bright celestial sources were automatically discarded. Human inspection of each image 
removed obviously bad data that occurred, for example, because of the presence of near field objects 
which flooded the focal plane with radiation. Otherwise, nothing was done to remove data which were 
inconsistent from one measurement to the next. Thus, many non-confirming sources appear in the 
maps. Each of the three sky coverages of the survey was made into a separate set of images so that spuri- 
ous, moving or variable sources could be detected by comparison or "blinking" of the coverages. A fourth 
set of images was made of the limited region of sky known as the mini-survey (Section III.C). 

Differences in the path length of the tine of sight through the interplanetary dust cloud caused 
significant variability of the zodiacal emission background from one survey coverage to the next. The 
Zodiacal History File consists of time ordered 0.5 x 0.5" averages of the intensity data at each wavelength 
and associated pointing information. This information was designed to simplify modeling and removal 

of the zodiacal component. 

A set of maps with 2' pixels of the Galactic plane for latitudes within 10" of the plane and two 
maps of the full sky with 0.5" resolution, one centered on the Galactic center and one centered on the 
Galactic anti-center, were also produced. 

A-3-h Final Processing Steps 

After the data from the entire mission had been processed, several programs operated on the WSDB 
to produce the final entries for the point source catalog. For example, these programs applied final cali- 
bration corrections, deleted unreliable sources in regions of high source density, and searched for associa- 
tions in other astronomical catalogs. 

All of the objects found in the point source catalog were selected from the WSDB according to cri- 
teria that depended on the density of sources in a 1 sq. deg area containing the source (see V.H.6 for a 
detailed discussion). Outside of high-density areas sources were required to have at least one band that 
was hours-confirmed at least twice and at least one of those sightings had to have a valid seconds- 
confirmed detection without the alibi of a failed detector. For objects detected in two or more bands, this 
rule was relaxed somewhat. Inside high-density areas, sources had to satisfy much more stringent cntena: 
they had to be perfectly hours-confirmed in one band at least twice, have no brighter or confusing close 
neighbors, have no fluxes more discrepant than a factor of three and have satisfied a detection correlation 
coefficient threshold greater than 0.97. 


B. Pointing Reconstruction 

The term "pointing reconstruction" as used here refers to the process of reconstructing the pointing 
direction of a fiducial point in the focal plane, as well as the twist angle of the telescope about the refer- 
ence point, referred to as the boresight. Since the accuracy with which the pointing could be recon- 
structed varied as a function of time, uncertainty histories were also specified. Use was made of all avail- 
able fine attitude sensor data to maximize the accuracy of the reconstructed pointing during the various 


V-6 


attitude modes. A smoothed estimate of the pointing was obtained using a recursive form of an extended 
Kalman filter. A more detailed discussion of the pointing reconstruction than is given here is provided by 
McCallon and Kopan (1985). 

The pointing reconstruction was based on the output of the fine attitude sensors which include eight 
slit-type visible star sensors, a three-degree-of-freedom gyro package and a dual-axis fine Sun sensor. The 
star sensors were mounted on the periphery of the focal plane, four slits being normal to the scan direc- 
tion, and the remaining four slits skewed at angles of approximately ±40° (see Fig. II.C.6). The slits were 
used in pairs, one normal and one skewed, to provide two axes of boresight information at each visual 
star sighting, hereafter called a fine attitude calibration, or FAC. The gyros and fine Sun sensors were 
mounted together on the outside of the spacecraft. There were, of course, misalignments between the 
gyros and fine Sun sensors due to limits on the accuracy with which they could be positioned. Unfor- 
tunately, the misalignments between the gyro-Sun sensor package and the telescope were found to vary 
with time (mainly in cross-scan) due to bending of the telescope mount with respect to the spacecraft due 
to temperature gradients. 

Pointing reconstruction was accomplished by integrating the gyro outputs to provide an initial esti- 
mate of the attitude history. Concurrent processing of frequent Sun sensor measurements and relatively 
infrequent FACs was performed with an extended Kalman filter to estimate model parameter errors and 
refine the attitude history. The model parameters estimated include initial-attitude errors, spacecraft- 
telescope misalignment errors, gyro drift, gyro scale factor errors, and gyro alignment errors. Gyro drift 
was modelled as a constant rate plus a random Gauss-Markov process. Process noise was added to all 
model parameters with the exception of the initial attitude parameters to allow for modelling errors and 
slow changes with time. 

The recursive Kalman filter was run first forward in time and then backward over a scan and then 
the two estimates were combined to provide a smoothed estimate sampled every second. The Kalman 
filter was reinitialized with the updated model parameters, and the process was repeated for the next scan 
using any available boundary conditions. The scan was chosen as the basic processing block because it 
normally contained at least two FACs and was of manageable length. A special file of boundary condi- 
tions was maintained for each SOP, and processing of observations with fewer than two FACs was 
deferred until the multi-FAC observations were finished in order to pick up final boundary conditions for 
the others. 

This approach to the pointing reconstruction worked quite well as shown by the statistics on known 
star position matches (see Section VII. C). One danger which was carefully guarded against was giving 
bad data to the filter. Considerable effort was put into algorithms to perform automated consistency 
checking and to provide measurement rejection capability. There were problems associated with each 
type of attitude sensor. 

The fine Sun sensor hardware had severe spiking problems from the beginning of the mission. This 
had the greatest effect on spacecraft control, causing several fallbacks to the safe attitude control 
mode. The spikes were thrown out in the pointing reconstruction algorithm by consistency checks. Spikes 
occasionally caused a brief cross-scan excursion in the actual pointing. 


V-7 


Another problem associated with the fine Sun sensors was less dramatic but potentially of greater 
consequence for the pointing reconstruction. The y-axis fine Sun sensor characteristics varied with time 
over the mission. Thus the transfer function used to convert from the integer output of that Sun sensor 
to the desired cross-scan angle should have been varied slowly over the course of the mission. This prob- 
lem was only partially alleviated by the fact that the spacecraft-telescope misalignment angle about the y- 
axis was being continually re-estimated. The pointing reconstruction used only two transfer functions for 
the y-axis fine sun sensor over the course of the mission. Ideally, a separate transfer function would have 
been used for every 75 to 100 SOPs. 

The gyros also had their difficulties. Early in the mission it was apparent that the gyros (espe- 
cially one of the z-axis gyros, denoted Gyro ZA) were noisier than expected, and that their characteristics, 
especially those of the x-axis gyro, varied with time. Gyro ZA was used by the on-board control system 
for in-scan control. Unfortunately the control system was not designed for a gyro as noisy as ZA, and 
this resulted in the occasional occurrence of a phenomenon referred to as the limit-cycle-burst problem. 
During a burst, the limit cycle amplitude and frequency increased by a factor of three or four; amplitudes 
of 13" were observed. This problem occurred much more frequently after the start of the third hours- 
confirming coverage and was finally solved by switching to Gyro ZB for control. The limit cycle bursts 
were, however, reconstructed without apparent difficulty. The long-term changes in gyro characteristics 
were tracked by the Kalman filter, but process noise had to be increased to give the filter enough freedom 
to follow the changes. There was also some indication of even faster changes suggestive of thermally 
induced misalignments between the gyros and the fine Sun sensors. These problems worsened as the mis- 
sion progressed. The x-axis gyro was dropped from use in reconstruction beginning with SOP 187 and 
the y-axis gyro was dropped beginning with SOP 316. 

The gyros were far more susceptible to the Earth’s magnetic field than expected. This problem was 
discovered early in the mission, and a software compensation algorithm was developed for pointing 
reconstruction. 

Very early in the processing it was noted that errors in the reconstructed position of known infrared 
sources were sometimes larger near certain FACs rather than smaller as expected. It appeared that the 
FAC was degrading the solution. These cases became generally known as "biased FACs". Most biased 
FACs had errors on the order of 5" to 10", but some were much worse, up to 6'. It was determined that 
the worst biased FACs were the result of having selected stars for FACs which had other bright stars close 
enough to interfere with the star sensor observations. Occasionally, biased FACs were the result of cata- 
log errors. An algorithm was developed to identify the FACs with disturbing stars, and in reprocessing 
SOPs 29 through 446 these FACs were either given larger uncertainties or deleted. The worst biased 
FACs (errors greater than 30") in the 446 to 600 SOP range were corrected in the same manner. 

Another problem was the thermal misalignment about the y-axis (cross-scan) between the gyro/Sun 
sensor package and the telescope. This caused particular difficulties during the third sky coverage because 
of large and frequent changes in solar aspect angle. 

Lack of a time-variant y-axis fine Sun sensor transfer function had the greatest effect on cross-scan 
reconstruction for observations with slews between the FAC and the survey scan. In this situation the 


V-8 



Kalman filter was unable to compensate for the y-axis fine Sun sensor errors by adjusting the telescope 
misalignment about the y-axis as it would normally do. In the worst-case, cross-scan slews on the order 
of one to two degrees resulted in reconstruction errors as large as 25". The actual magnitude of the error 
varied greatly depending on the exact start and end points of the slew and was not necessarily greater for 
longer slews. The only survey scans affected by this problem were those with out-of-scan FACs. 

Lack of the improved y-axis thermal misalignment model was significant only for observations 
which were both without a FAC and without a boundary condition on the same end. In the worst-case 
situation, cross-scan reconstruction errors as large as 30" were possible for a one-FAC observation; this 
could happen on full length observation following a maximum cross-scan slew with the worst possible 
placement of the FAC. Such large slews were executed only after the start of the third survey coverage 
(SOPs 426 to 600). For SOPs 29 to 425, peak errors due to this problem were less than half as large. 

C. Source Detection 

The calibrated raw data for each of the 59 operating detectors were examined for point sources and 
small extended sources. The detection of the latter is described in Section V.E.l. For each observa- 
tion the accepted point source detections were passed, with detector number, time of detection (and 
uncertainty), flux (and uncertainty), signal-to-noise ratio (SNR), and the correlation coefficient with the 
point source template (CC, see below), to the seconds-confirmation processor (Section V.D.2). A noise 
history was also created for each detector. If a detection occurred in a one-second period in which the 
analog-to-digital converter was saturated, then the detection was flagged. 

C. 1 Square Wave Filter 

The first step in the detection process was to search for potential sources by applying a narrow 
bandpass digital filter to the detector data streams. This filter consisted of an eight-point zero-sum 
square-wave function. The effect of the filter was to subtract the first two and last two points from the 
sum of the middle four points; more formally, for a sequence of data points x, (Fig. V.C.la), the ampli- 
tude of the square-wave of x at the point i is defined as: 

E(x,i) - —Xj ~ X, + 1 + X I+2 + *,+3 + x,+4 + x, +5 - x 1+6 - x, +7 (V.C. 1) 

This square-wave filter was applied at each point in the data stream, and a search was made for 
positive square-wave excursions between zero crossings, defined as a pair of data points (ij) such that (see 
Fig. V.C. lb) 

E(x,i) > 0; 

E(x,k) > 0 for / < k < j; 

E(xJ+ 1) < 0; 

and for some n, E(x,n) < 0 

and E(x,n), E(x,i— 1) < 0. 

That is, find the values n, i, j such that: n with E(x,n) < 0; the first / > n with E(x,i) > 0; and 
the first j > i with E(xJ+ 1) < 0. The positive excursion (ij) has a peak at the first p with i < p < j, 
such that E(x,p) is maximal among E(x , /),..., E(xJ). Peaks with square-wave amplitudes, E{n,p), 
greater than 2.5 times the noise N x were passed on as candidates for point sources. 


V-9 


(a) 


+ 1 


0 


-1 



i 


i+2 i + 4 

• • • • 

i + 1 i + 3 


i +6 

• • 

i + 5 


i + 7 


i + 8 

i + 9 


(b) 


E(x, i) 



Figure V C 1 a) An eight-point zero-sum, square-wave filter was applied to the data streams (top 

panel); b) The detection processor looked for positive square wave peaks between 
zero crossings in the filtered data stream (bottom panel). 


C.2 Noise Estimator 

The noise N x for a data stream was defined as the median of all E(x,p) for square-wave peaks p. 
Such positive square-wave excursions occurred about once every 6 samples. It was found from prelaunch 
simulations and from analysis of in-flight data that this median noise estimator gave a reasonable 
representation of the rms noise, in the sense that 


®rms 


1.2 N x 


(V.C.2) 


The enormous volume of data meant that determination of a running estimate of the rms noise 
would have involved a prohibitive computational run time. 

The initial value of N x was the median of the first 50 square-wave peaks. N x was then updated at 
every square-wave peak E(x,p) as follows: 


if E(x,p) < N x then reduce N x by the factor X (< 1); otherwise, increase N x by the factor 1A. 


V-10 



The parameter X controlled the stability of the noise estimator. As X approached 1, the noise esti- 
mator became very stable, but it also lagged behind any change in the noise by about 5/(1 - X) samples. 
In tuning the value of X, great importance was attached to achieving a stable noise estimate at high 
Galactic latitudes, where the noise was mainly due to detector noise, and a value X - 0.95 was set at 12, 
25 and 60 pm, and 0.90 at 100 pm. This meant that the noise estimate lagged by about 25', 25', 50', 50' 
at 12, 25, 60 and 100 pm. Regions with steep gradients in the density of point sources, such as the 
Galactic plane, had large gradients in the noise amplitude. Hence the noise was underestimated as the 
plane was approached and overestimated after it was passed (Sections V.C.7, VTII.D.6). This error was 
very large, and since sources were thresholded partly on signal-to-noise ratio (Section V.C.4, below), the 
effective threshold was raised to very large values after passing the Galactic plane, resulting in a shadow 
zone in which few sources were accepted. To keep the extent of the 100 pm shadow zone no larger than 
that at 60 pm, X was set to 0.90 at 100 pm compared with 0.95 at 60 pm. However, this adversely 
affected the stability of the 100 pm noise estimate in the presence of cirrus at higher Galactic latitudes, 
resulting in the rejection of some detections that should have been accepted and hence a reduction in 
completeness of the catalog at 100 pm (Section Vffl.D). 

The noise estimate was maintained in a noise history file for each detector after multiplication by 
the factor to convert it to an estimate of the rms noise on a single sample. To compress the size of this 
file, an entry was made only if linear extrapolation of the previous two entries would lead to an error 
greater than 35%. 

C.3 Timing Estimate 

The time of the square- wave peak at E(x,p) was estimated from the maximum of the parabola 
passing through the three points (p - 1 ,E{x,p - 1)), (p,E(x,p)), (p + 1 ,E(x,p + 1)). The delay between 
a source in the unfiltered data and its peak in the square-wave function was subtracted from the estimate 
to give the detection time. A small offset to account for electronic delay and the sampling time of the 
detector was included. The timing uncertainty was taken from a look-up table as a function of the values 
of signal-to-noise ratio and correlation coefficient for the source. 

C.4 Correlation with Point Source Template 

The heart of the point source detection processor was the comparison of the data for candidate 
sources selected by the square-wave filter with the profile, or template, expected for an ideal point source. 
For this purpose the 1 1 samples centered on the candidate detection time, y,, i — 0 ... 1 0 , were compared 
with the appropriately shifted template /?, superimposed on a linear baseline. The amplitude A of the 
detection was determined from fitting the 1 1 data values y, to the function 

ARj + iM + B, /-0...10, (V.C.3) 

where B is the baseline height and M is its slope. A , M , and B were determined by the method of least 
squares, i.e., by minimizing 


U - I {ARi + iM + B - y,) 2 


where I stands for [X 


10 

I ] 


i-0 


V-ll 


Thus, 


r 

ZRtRi Zy,Ri 

ZiR 

/ 


I 

LRi 

Zy, 

55 



V 

HRj Ziy t 

385 


T 

1 1 RjRi I Ri 

ZyjR, 



1 

ZRj 

11 

Zyi 



V 

UR 

,■ 55 

Ziyi 



J 

1 Zy t Ri I Ri 

ZiRi 



1 

Zy t 

11 

55 



V 

Ziyi 

55 

385 





ZR t Ri 

Z Ri 

ZiRi 1 

where V — 

ZR, 

11 


55 



ZiRj 

55 

385 | 


(V.C.4) 


The correlation coefficient of y, with R, is given by 

CC - lZiR,/(LZ? x I/??)* (V.C.5) 


where Z, — y, ~ ( iM + B). 

A candidate detection was accepted only if 

(i ) CC > 0.87 

and (V.C.6) 

(ii) SNR - A/l.2N x > 3 , 

where the factor 1.2 converts the median noise estimate to an rms noise estimate (see Section V.C.2). 

The total rms uncertainty in amplitude, A, over the 1 1 data samples can be shown to be 

— - J2. CC(1 - CC 2 r' h (V.C.7) 


Thus the correlation coefficient is a measure of the local signal-to- noise ratio and a threshold of 
0.87 corresponds to a signal-to-noise ratio of about 2.5. In regions where the noise was roughly indepen- 
dent of time, the main thresholding was therefore provided by the signal-to-noise ratio. The square-wave 
filter threshold (Section V.C.l above) was set low so that as few acceptable detections as possible were 
rejected, within the constraints of the available computer time. It should be noted that a low correlation 
coefficient for a bright point-source is probably an indication that the source is slightly extended. In 
regions of high source density (see Section V.H.6), where extended structure is a considerable problem, 
the correlation coefficient threshold was increased to 0.97. 


V-12 



C.5 Determination of Templates 


The templates for each wavelength band were stored with a sampling frequency 64 times that of the 
survey data. The candidate detection time in sampling was determined, rounded off to the nearest 1 /64th 
of a sample and the appropriate 1 1-point template selected by taking every 64th point from the template 
array. 

Immediately after launch, predicted detector responses to an ideal point source were used. Compo- 
site templates were constructed for each detector using sources detected with high correlation coefficient 
and signal-to-noise ratio. The a priori templates were replaced with the composite templates and the 
analysis repeated, using 12 hours worth of data. Convergence was achieved after only a few iterations. 
Figure V.C.2 shows representative point source templates for one detector in each wavelength band. 
Since no evidence for detector-to-detector variation within a band was found, the results for all the detec- 
tors in each band were averaged together to produce the final composite templates. 

C.6 Low Sienal-to-Noise Detections 

A secondary class of detections called low signal-to-noise detections was defined as those with 
signal-to-noise ratios between 3 and the threshold required for a valid detection. Because the threshold 
for valid detections was itself set at 3, no low signal-to-noise detections should have been generated. 
However, due to the round-off errors in the computation of the signal-to-noise ratio, a few were 
created. These were only used to provide upper limits for sources confirmed in other bands. 

C. 7 Source Shadowing 

When two sources crossed the same detector within 6 samples of each other, i.e., within 1.4 , 1.4 , 
2.9', 5.9' of each other in the scan direction at 12, 25, 60 and 100 pm, respectively, the detection of one 
or both of the sources may have been inhibited. Generally, the brighter source was detected without 
mishap, but the fainter source may have had its baseline so modified by the brighter source that it 
failed to be detected at all. This is the phenomenon of source "shadowing". A source may have been 
shadowed in a longer wavelength band but detected perfectly at shorter wavelengths. To warn of the pos- 
sibility of this effect, sources were tagged at a later stage in the processing (see Section V.H.3) if they had 
near neighbors. The fluxes of such flagged sources should be regarded with caution. No significance 
should be attached to the absence of a detected flux in a shadowed band. The completeness figures 
given in Chapter VIII do not apply to the shadow zone around a source. 

D. Point Source Confirmation 

D. 1 Processing Overview 

The arrangement of the detectors in the focal plane of the telescope and the survey strategy permit- 
ted reobservation of inertially fixed point sources after time intervals of several seconds, several hours, 
and several weeks. The confirmation process consisted of examining those multiple observations and 
identifying which plausibly belonged to the same object. Once this identification was made for a given 
source, refinement of the parameters describing that source was performed by combining the observations 
into a single improved description. 


V-13 



(a) 


(b) 



- 1.2 - 0.9 - 0.6 0.3 0.0 0.3 0.6 0.9 1.2 - 1.2 - 0.9 - 0.6 - 0.3 0.0 0.3 0.6 0.9 1.2 


(c) (d) 



IN-SCAN POSITION IARC MIN) IN-SCAN POSITION (ARC MINI 

Figure V.C.2.a-d Detections found by the square-wave filter were compared with the response of the 
telescope-detector-electronics combination to a true point source. Representative 
point source templates are shown for one detector in each wavelength band. 

When comparing two sightings, positional agreement was always a part of the decision whether to 
accept them as the same object. Over the seconds and hours intervals, photometric agreement was 
required as well. Because the cross-scan position was tested only by requiring that a real object be sighted 
by a compatible pair (or triplet) of detectors, the decision problem at the seconds-confirmation level 
involved only in-scan position agreement, and the tests could be based on Gaussian error models. At 
the hours and weeks level, the position error had to be modeled as a non-Gaussian random variable 
because of the uniform uncertainty due to the cross-scan extents of the detector slots. 

The hours-and weeks-confirmation decision was based on the correlation of the probability density 
functions which describe the two-dimensional position information. The parameter which was required 
to be above a certain threshold to confirm two sightings was the cross-covariance of these density func- 
tions, evaluated at the separation of the nominal positions. The formalism is discussed and derived by 
Fowler and Rolfe ( 1 982). There were several virtues in this approach, among which the most important 
was its freedom from the Gaussian approximation. The position error due to the detector slots was uni- 

V-14 


in: 


III muiiii I H iiHMMUiHillnilllBia MHUII MM IIIIUN Mil 1IIIIINM ■! m*\m\ Ml I ■Milfl «. |||i MMII Mil INMMM1IMI III lllf lliW| RIMIMMHI *1 INI | INS! || W»ll M III II HIM Hill Hi billlllll MV I ill II 1 1 III I 111 II III IHVIIIMHI N Hill HH III II I Hi Hill | ll|| l||lim|l»ll|llll| (IMIVIlii 


ORIGINAL PAGE IS 
OF POOR quality; 


formly distributed, making any Gaussian algorithm unacceptable. Another aspect of the decision algo- 
rithm was that the fraction of true cases accepted increased as the size of the position uncertainties 
decreased. Because the threshold could not easily be set to accept a specific fraction of all cases, it was set 
during simulation tests by opening up the the threshold until just before the acceptance of false events 
became significant. 

Figure V.D.la depicts two position probability density functions in a typical case. Each density 
function has a broad, flat ridge which shows the uniform contribution of the net intersection of the 
seconds-confirmed detector slots’ cross-scan domains. The orthogonal direction is that of the scanning 
motion, and the corresponding position errors were found to be well modeled as Gaussian. In the figure, 
the agreement between the two position estimates is about as good as it could be. The large difference 
shown between the scan directions occurred only at the weeks-confirmation level. At hours-confirmation, 
the major axes of the density functions were approximately parallel. Refinement of position leads to a 
new probability density function which is more centrally concentrated, as shown in Fig. V.D.lb for the 
current example. 

The confirmation decisions are summarized in Table V.D.l which gives, for each step in the 
confirmation process, the type of position test (Gaussian or two-dimensional), the threshold used, the flux 
agreement required and the net effect of these criteria on real sources. 

Although the confirmation decision and parameter refinement lay at the core of the point source 
confirmation processing, many peripheral issues also had to handled. These aspects are discussed at a 
level of detail which attempts to be concise while not leaving an inordinate number of questions 
unanswered. 



Figure V.D.l a) As described in the text, the position and associated uncertainty of each source is 
represented by a probability distribution function consisting of Gaussian and uniform 
components. Shown here (left) are the distribution functions for two sightings of a 
single source, b) The position of the source resulting from the merging of the two 
sightings shown in V.D. 1 .a is described by the new probability distribution function 
shown here (right). 


V-15 


Table V.D.l Confirmation Summary 


Confirmation 

Level 

In-Scan 

Limit 

X-Scan 

Limit 

2-D 

Limit 

Flux Ratio 
Limit 

Comment 

Seconds 

4.4 

Sigma 

Compat. 

Det. 

- 

2 

1 Real Match 
per 50,000 
Rejected 

Band 

Merge 

4.2 

Sigma 

Compat. 

Det. 

- 

- 

1 Real Match 
per 40,000 
Rejected 

Hours 

- 

- 

1x10 _4 sr -1 

2 

1 Real Match 
per 20,000 
Rejected 

Months 

- 

- 

1x10 _4 sr -1 

- 

1 Real Match 
per 20,000 
Rejected 


D.2 Overview of Seconds-Confirmation 

Point sources were required to be observed more than once within a few seconds by redundant 
detectors in the focal plane to eliminate false alarms caused by various noise processes. Exceptions to 
this rule were granted only on the basis of known problems in a redundant detector, either an outright 
failure or a significantly degraded sensitivity compared to other detectors in the same band. Because of 
measurement errors in the photometry and timing, it was not always trivial to identify which detections 
were in fact observations of the same inertially fixed point source. Noise, radiation hits and confusion by 
more than one source aggravated the problem. 

The input data to the seconds-confirmation processor are described in Table V.D.2. When several 
detections in a band were accepted as applicable to a single point source, the information contained in 
the multiple observations was combined to obtain a single refined description of the object. 

The seconds-confirmation process first dealt with detections in each band separately and later tried 
to combine observations from all bands. Comparisons of IRAS observations with known infrared sources 
were then made. Along the way, considerable statistical analysis was performed to calibrate the focal 
plane geometry as mapped onto the sky through the optics and to verify the quality of the position recon- 
struction. 

D.2.a Band Seconds-Confirmation 

In-band seconds-confirmation was done in the position domain rather than the time domain to 
account for variations in the scan rate. The design of the focal plane permitted legitimate sources to pro- 
duce pairs or triplets of detections, and these modes were handled separately. When neither mode 
appeared correct, confusion processing was required. Although this salvaged useful detections, it could 
also let false alarms leak through. The confirmation process, however, may be considered a serial filter 



arrangement, and so a few false alarms slipping through one stage are very unlikely to penetrate the entire 
obstacle course. 

After accepting a confirmed point source, the flux and position estimates were refined. The multi- 
ple observations were used in statistical computations aimed at gathering information about the instru- 
ment and the a priori statistics. 

D.2.b Position Reconstruction 

Detections in a single band were processed in time order. Each detection had its position in the 
Sun-referenced coordinate system computed. At the same time the photometric uncertainties, modeled 
as a white-noise Gaussian random variable in the logarithm of the flux and a constant but unknown cali- 
bration scale factor error in the flux, were obtained from lookup tables as functions of signal-to-noise 
ratio and correlation coefficient. 

The position reconstruction was performed by searching the pointing history for records with time 
tags bracketing that of the detection. The matrices for transforming from the focal plane coordinates to 
the Sun-referenced system and to the 1 950.0 mean ecliptic system were obtained by cubic spline interpo- 
lation. Each detector was characterized by a unique unit vector in the focal plane coordinate 
system. The transformation matrices yielded the polar, azimuthal, and twist angles defining the posi- 
tion of the detector slot on the sky in the Sun-referenced and 1950.0 mean ecliptic systems. The cross- 
scan half-width of the slot defined the uniform error component, denoted L z , in that direction. The 
cross-scan limit cycle reconstruction uncertainty was modeled as a Gaussian random error, a z . These 
and the in-scan errors discussed below were the only significant uncertainties involved in the 
confirmation decision process. After all confirmation and refinement processing was completed, the abso- 
lute position angle errors, which were estimated by the pointing reconstruction program but which cancel 
out of the confirmation decision, were taken into account. 

The uncertainties in the scan direction which affected the confirmation decision were a time- 
interpolation error and the detection timing uncertainty. The latter yielded an angular uncertainty when 


Table V.D.2 Input Data for In-Band Seconds Confirmation 

Detections: 

Flux and uncertainty, time and uncertainty, signal-to-noise ratio, 
detection correlation coefficient, and detector number. 

Pointing: 

Telescope boresight angles in the Sun-referenced system and 
uncertainties, scan rate, time tag, sines and cosines of the angles 
in both the Sun-referenced and 1950.0 mean ecliptic systems, 
and rates of these angles, all sampled at one-second intervals. 

Detectors: 

Status (on/off) and geometrical models in image space. 

Optics: 

Point-spread function model. 

Error Tables: 

Model of pointing reconstruction error due to time interpola- 
tions; photometric error as a function of signal-to-noise ratio and 
correlation coefficient for each band. 

Thresholds: 

Confirmation acceptance, search windows, etc. 


V-17 





multiplied by the scan rate. These were taken to be independent Gaussian errors with a combined effect 
of a zero-mean Gaussian random error, denoted o Y , with a variance equal to to the sum of the two error 
variances. 

D.2.c Optical Crosstalk Removal 

Each detection was checked to see whether it was bright enough to cause spurious sources due to 
optical crosstalk on adjacent detectors (see Section VII.E.5). This decision was made on the basis of the 
bright source’s signal-to-noise ratio, since the detection process triggered on that parameter, not flux level. 
For example, a source which might elicit optical crosstalk detections in quiet sky might not do so in the 
Galactic plane, where the higher noise level could mask the crosstalk. 

A lookup table for each band was used to determine the cross-scan distance over which a source 
might cause optical crosstalk detections. This was a function of the signal-to-noise ratio of the source, 
and was based on the prelaunch evaluation of the optical point-spread function. The signal-to-noise 
ratio thresholds for closer examination were 5000, 1200, 300, and 200 in the 12, 25, 60, and 100 pm 
bands, respectively. When a detection was above this threshold, the cross-scan distance to search for 
other detections caused by crosstalk was obtained from the lookup table. The in-scan search distance was 
constant, with values of 14", 14", 29", and 58" in the 12, 25, 60, and 100 pm bands, respectively. Any 
detections in the window which were fainter than the one being processed were deleted from further 
consideration. The deletions were performed only for the module containing the detector which yielded 
the bright source, in order not to eliminate its confirmation partner. 

No attempt was made to identify crosstalk caused by the secondary mirror support spider. The 
characteristics of the bright source and those deleted were subsequently used in the small extended source 
processing. An analysis of the effects of bright sources is given in Section VII.E.5. 

D.2.d In-Band Seconds-Confirmation Decision 

The oldest detection being considered at any time was processed for seconds-confirmation by 
searching the rest of the buffer for detections which had not yet been used and which occurred on compa- 
tible detectors. This oldest detection was called the drop-dead detection, because it had to confirm with 
another detection in the buffer or be rejected. This terminology runs through band merging, hours- 
confirmation, and weeks-confirmation, as well as the similar phases of small extended source processing. 
Coarse windows in time and in-scan position were used to isolate possible candidates. When such were 
found, the fine position test was applied. This consisted of an no test on the absolute in-scan position 
discrepancy between the drop-dead and the candidate, in units of the standard deviation of the 
discrepancy random variable. Assuming independent errors, this is the square root of the sum of the in- 
scan error variances of the two detections. This test was used primarily for its computational speed and 
for the fact that it permitted one to set the threshold by selecting the fraction of all true events which one 
was willing to sacrifice in order to deter false events (Table V.D.l). 

When the coarse in-scan window was exhausted, the search for candidates was terminated. At that 
point, the number of candidates which passed the fine position test with the drop-dead determined the 
next step. If only one was found, the double-detection mode was processed. If two were found, the 
triple-detection mode was examined. If more than two candidates satisfied the position agreement 
requirement, then confusion processing was invoked. Each of these three cases is discussed below. 


Ill 


V-18 



If no acceptable candidates were found, or if none remained after the additional testing discussed 
below, then the drop-dead detection was rejected. This means that it was considered either non-seconds- 
confirmed (NSC) if all of the compatible redundant detectors were operational, or non-seconds-confirmed 
due to a failed detector (NSCF) if the alibi of a dead or degraded detector was applicable. 

D.2.e Double-Detection Mode 

The double-detection mode was the most common mode of seconds-confirmation. A combined 
test on flux and position was applied which consisted of a two-degree-of-freedom % 2 test. The no position 
discrepancy contributed one of the degrees of freedom, and the discrepancy in the log of the two fluxes in 
units of the log-flux standard deviation provided the other. A threshold value corresponding to sacrificing 
one real event out of every hundred thousand was used for this test for unconfused cases, and one 
corresponding to sacrificing one real source in a million was used for cases which had been processed for 
confusion before arriving at this point. If this test failed, the drop-dead fell back to rejection processing as 
described above; otherwise parameter refinement was performed for the flux and position (see below), and 
the next detection was processed. 

D.2.f Triple-Detection Mode (Edge Detections) 

Some scan paths available to inertially fixed point source images traversed three detector slots in the 
same wavelength band, as portrayed in Fig. V.D.2a. Only some combinations of detector slots were con- 
sistent with this possibility. When three detections were found to be confirmable on the basis of position, 
the slots involved were tested for compatibility with this triple-detection mode. If they were consistent, 
then the source position could be localized to within the small region of cross-scan overlap, and a small 
uniform uncertainty in the cross-scan position of the source resulted. Many sources in the catalog have 
at least one edge detection, resulting in correspondingly small cross-scan uncertainties. 


C 


★ 


] 


★ 


] 


★ 




(a) TRIPLE-DETECTION MODE 


(b) INCONSISTENT SLOT 
COMBINATION 


(c) CONFUSION ON 
CONSISTENT SLOT 
COMBINATION 


Figure V.D.2a-c The confusion processing described in the text attempted to deal with the various 
combinations of possibly multiple sources seen on more than one detector. 


V-19 



If the slots were not consistent with an edge detection, then the situation was diagnosed as confused 
(see below). An example of three detections which were not consistent with the triple-detection mode is 
shown in Fig. V.D.2b. Whenever a legitimate triple-detection pattern occurred, the possibility that two or 
more point source images were confused on the central detector had to be considered. Figure V,D.2c 
shows two point-source image paths which are close together in the scan direction but separated in the 
cross-scan direction, so that the central detection is a confused response to the two. The additional 
cross-scan coverage of the triple-detection mode increased the probability that such confusion would 
occur. The only thing that could be done to avoid accepting this case as a single point source was to 
require the edge-overlap detections to be of roughly comparable brightness, so that if they were actually 
two different sources, one would expect the central detector to see about twice as much flux. This would 
tend to cause the flux test to be failed, so that the triple-detection mode would be rejected, and confusion 
processing would be invoked. A faint confused source might not upset the flux test, and so relative faint- 
ness on an edge-overlap detector itself was made to be a reason for disqualification of the triple-detection 
mode. On the other hand, if only one source were involved but it barely grazed one of the edge-overlap 
slots, there would be no great loss in disallowing it, as it contained little information anyway, although 
the detection could cause confusion again at hours-confirmation. 

These considerations led to requiring the ratio of fainter to brighter edge-overlap detection to be 
greater than 0.01 in order to continue to entertain the possibility of a triple detection mode. Failure to 
pass this test caused the situation to be processed for confusion processing. 

Next a flux test was applied to the brighter of the two edge-overlap detections and the full-hit 
detection. This was an na test on the log-flux discrepancy and had a threshold of 4.4 a (i.e., reject one 
real source per hundred thousand). If this test failed, confusion processing was applied; otherwise param- 
eter refinement was performed for position and flux. The fainter edge-overlap detection was not used for 
flux refinement, and the brighter was used only if the ratio of the fainter flux to the full-hit flux was less 
than 0.1, as this indicated that the brighter edge-overlap detection’s slot should have intercepted essen- 
tially all of the flux. To reduce the effects of spurious weak detections, the fainter edge-overlap detection 
could contribute to the position refinement only by having a signal-to-noise ratio above 5 and a detection 
correlation coefficient above 0.96. If it qualified, then the pairwise position refinement described below 
was performed first for pair consisting of it and the full-hit detection, so that the slot-extension logic (see 
below) could be activated for this pair. The brighter edge-overlap detection contributed position and flux 
information according to the same rules as the double-detection mode; these are discussed in the parame- 
ter refinement sections below. 

D,2.g In-Band Seconds-Confirmation Confusion Processing 

The drop-dead detection and all candidates which had passed the fine position test were immedi- 
ately marked as confused upon initiation of confusion processing. Status bits set here and in correspond- 
ing subsequent analyses were carried in a status word, hereafter denoted CSTAT, throughout all remain- 
ing processing stages (see Table V.D.3a,b). 

Several attempts were made to unravel confused situations. The first stage involved discarding any 
detections with signal-to-noise ratios below 5 or detection correlation coefficients below 0.92. Such detec- 
tions were never used in any future confirmation processing. If this dropped the number of candidates to 


V-20 


zero or one, then the drop-dead was processed for rejection or for the double-detection mode, respec- 
tively. Otherwise confusion processing continued by attempting to identify pairs of candidates which 
would confirm each other without the help of the drop-dead detection and without confusion from other 
detections in the buffer. Any such detection pairs were set aside from the current situation and left for 
the time when one of them would become a drop-dead. Finding such pairs reduced the number of can- 
didates involved in the confused situation being processed, and after all such possible reduction had been 
achieved, the number of remaining candidates was checked as described above. 

If the situation remained confused, then processing continued by casting aside the detections whose 
slots had an edge-overlap relationship with the drop-dead’s slot, if any, unless the detector status informa- 
tion indicated that these were allowable for pairwise confirmation with the drop-dead. Again the branch- 
ing possibilities were checked. In this case, and in all remaining cases in this section, discarded candi- 
dates remained eligible for future confirmations on subsequent processing. 


Bit No. 

Table V.D.3a Confusion Status (CSTAT) Bit Assignments 

Meaning (Applies if Bit - 1) 


0 

Either triple-detection mode or possible 
confusion in seconds-confirmation 


1 

Triple-detection mode accepted 


2 

Confusion diagnosed and cleanup attempted 
in seconds-confirmation 


3 

Confusion diagnosed in band-merging, only 
best match retained 


4 

Confusion diagnosed in hours-confirmation, 
only best match retained 


5 

In 12 pm Band Byte: Not Used 
In 25 pm Band Byte: Not Used 
In 60 pm Band Byte: Detector number arrays 
(all bands) are out of 
time order 

In 100 pm Band Byte: More then 3 sightings 

were hours-confirmed 


6 

ADC saturation occurred for at least one 
detection in this band 


7 

Not used 



CSTAT 

lable V.D.3b. Common CSTAT Values (Decimal) 

Meaning 

3 

Triple-detection mode 

5 

Possibly confused 

7 

Triple-detection mode after confusion 


cleanup; usually not reliable 

8 

Possibly confused in band-merging 

13 

Possibly confused in both seconds- 


confirmation and band-merging 

16 

Possibly confused in hours-confirmation 


V-21 



If further candidate reduction was needed, the candidates were required to pass a stricter position 
test, similar to the one discussed above, but with a threshold value such that one real detection pair out 
of one hundred would be expected to fail this test. Again the branches considered above were checked, 
except this time the case of two remaining candidates was included, with a branch to the triple-detection 
mode processing if applicable, provided that a branch from that mode to confusion processing had not 
already been executed. 

If more than two candidates remained at this point (or two which had been found unacceptable for 
the triple-detection mode), then that combination whose summed fluxes came closest to those of the 
drop-dead were retained, and the others were discarded. If this left no candidates, the drop-dead was 
rejected; if only one candidate remained, the double-detection mode processing was performed. Other- 
wise, a flux test between the summed candidate fluxes and the drop-dead flux was performed which was 
similar to the triple-detection mode flux test, except that the threshold value used was chosen so that one 
real case out of a million would be rejected. If this test was passed, then all remaining candidates were 
confirmed with the drop-dead, which was considered to be a simultaneous confused observation of all of 
them. This type of salvaging was possible only when the drop-dead was the confused sighting (rather 
than one of the candidates). In this case, no flux refinement was performed; the individual candidate 
fluxes were retained, but position refinement was performed as described below. 

D,2.h In-Band Seconds-Confirmation Position Refinement 

Position refinement was performed for detection pairs as follows. The in-scan position angle and its 
uncertainty were obtained from those of the two detections by applying Gaussian estimation. The cross- 
scan position angle and its uncertainty were obtained by computing the nominal intersection of the slot 
extents as mapped onto the sky (i.e., including spacecraft limit cycle motion), and then including a slot 
extension" safety term if the ratio of the fainter detection to the brighter was below 0.8 in the 12, 25, and 
60 pm bands, and in all cases for the 100 pm band. This slot extension was obtained from a lookup 
table for each band which gave the amount of extension as a function of the flux ratio. The edge of the 
fainter-detection slot defining the net cross-scan range is the edge which was extended. For flux ratios 
greater than about 0.6, this extension was negligible. For a flux ratio of 0.4, the extensions in the 12, 25, 
60, and 100 pm bands were 13", 13", 28”, and 34", respectively. For a flux ratio of 0.2, the extensions 
are 20", 20", 43", and 65", respectively. In cases involving only detectors at the cross-scan boundary of 
the entire survey array, a slot extension to allow for the possible passage of the image outside of the 
detector area was added; the values used for this were 45* , 89 , 138 , and 206 in the 12, 25, 60, and 
100 pm bands, respectively. 

The error due to the net remaining slot overlap was modeled as a uniformly distributed random 
variable from this point on. This completed the computation of the position angles and their uncertain- 
ties in the uncorrelated-error frame of reference. The correlated errors (i.e., the relatively slowly varying 
absolute pointing errors) were taken to be those associated with the drop-dead source. These were 
modeled as Gaussian random variables, but were not yet convolved with the other uncertainties because 
they did not enter the band-merging problem (see below). Instead they were merely carried along so that 
they could be included after the band-merging was completed. 


V-22 



The twist angle for the refined position was simply interpolated linearly in time from those of the 
two detections being processed. No uncertainty was carried for the twist angle because the reconstruction 
error could have been an order of magnitude larger than the requirement without significantly affecting 
the position information; the reconstruction appears to have been well within its requirements. 

D.2.i In-Band Seconds-Confirmation Photometric Refinement 

Photometric refinement was performed for pairs of detections which passed several qualification 
rules. When performed, Gaussian estimation was used, where the Gaussian error was assumed to be in 
the logarithm of the flux. The estimates for the a priori errors were derived from analysis of a posteriori 
discrepancy dispersions as functions of signal-to-noise ratio in each band, and the combination of the 
Gaussian assumption with well-modeled variances was found to be quite good in subsequent % 2 . 

To be eligible for flux refinement processing, the ratio of the fainter detection to the brighter one 
had to be greater than 0.67; otherwise the brighter detection alone was used. In cases of confusion 
(except detections discarded in a previously processed case), only candidate fluxes were used, because 
confusion was detected only when it was the drop-dead detection that was composed of more than one 
source image. When the detections were unconfused and met the flux ratio requirement, they were still 
required not to show saturation of the analog-to-digital converter; if both detections were saturated, then 
the brighter flux was used. If one and only one was saturated, then the unsaturated detection alone was 
used. Finally, if the detections were on edge-overlap slots opposite a dead redundant detector (or a 
significantly noisy one which produced no detection confirmable with the pair being processed), then only 
the brighter detection was used for flux information, unless it was saturated, in which case the fainter 
detection alone was used. 

The drop-dead detection always provided the source name (time tag and detector number). The 
numbers of all detectors involved were always recorded in the arrays set up for this purpose (see Table 
X.B.6). When only one detection was used for flux information, it also supplied the signal-to-noise ratio, 
detection correlation coefficient, and baseline value. When refinement was done, the maximum signal- 
to-noise ratio and detection correlation coefficient were retained, and the unweighted average of the two 
baselines was kept. 

The order of the detector numbers in the array was determined as follows. In the double-detection 
mode, the first number is the candidate, and the second one is the drop-dead, unless one and only one 
detection was saturated, in which case the first number is that of the saturated detector. There is no way 
to tell the difference between both detections being saturated and only the candidate being saturated 
without recourse to additional data. In the triple-detection mode, the detector numbers are in decreasing 
detection time order if there was no analog-to-digital converter saturation involved. Otherwise the first 
number is the fainter edge-overlap detector, and the other two follow the rules for double-detection mode 
number order. 

D.2.i In-Band Seconds-Confirmation Statistical Processing 

When detections were found to be confirmable in the absence of confusion, saturation, effects due 
to failed or noisy detectors, fluxes which did not qualify for refinement, or occurrence of the detections 
exclusively at the boundary of the survey detector array area in the focal plane, then the discrepancies in 


V-23 


the in-scan position angle and the flux were computed for statistical analysis. The values and 
significances of these were used to adjust and then verify the a priori error variance models. 

In each band, the relative flux discrepancy standard deviation as a function of signal-to-noise ratio 
was computed. The mean discrepancy was not significantly different from zero, and typical standard 
deviations ran from 20%” for weak sources to about 10% in the 12 pm band, 15% in the 25 pm band, and 
20% in the 60 and 1 0U pin bands for bright sources. These standard deviations were used to examine the 
behavior of the photo me tric error at the sampling intervals involved. After about 30 days of data had 
been accumulated, th e var iances obtained in this way were used to supply new a priori flux error vari- 
ances for the rest of the, m ission. It was not necessary to cycle subsequent results through this process 
again, i.e., the statisti cal pro perties of the photometric discrepancies sampled at intervals of a few seconds 
and a few hours (fro m sim ilar analysis at hours-confirmation) were stable over the lifetime of the space- 
craft. ~~~ a 

D.3 Band-Merging ~ 

D.3.a Overview of Band-Merging 

Band-merging was per formed in a manner very similar to in-band seconds-confirmation, except that 
detection buffers were ma intained for all bands simultaneously, and detections in different bands were 
combined. Again a "drop -dead" source was selected as a nucleus to which detections in other bands 
were attached if they coujcf be associated with it. Up to four attempts were made to obtain some type of 
flux measure to put in_aIL three other bands for each drop-dead. In the first round, a detection from the 
primary source buffer jn each band was sought. These buffers contained sources which were either 
seconds-confirmed (SCJor were on detectors opposite a failed or noisy redundant detector (NSCF). The 
drop-dead source was selected as the SC or NSCF source with the earliest time tag. The drop-dead’s 
band determined the order in which the other bands were searched, according to Table V.D.4. 

Coarse windows liTtime and in-scan position angle were used to limit the search for merging candi- 
dates. When found, such candidates were subjected to a fine position test such as that of in-band 
seconds-confirmation, ajLna test on the in-scan position angle (Table V.D.l). In the cross-scan direction, 
the nominal slot extent™!" the drop-dead and the candidate were required to overlap. 

If more than one can didate was acceptable in a given band, the priority was given to SC candidates 
over NSCF candidate s, _and ~ otherwise the best in-scan match was taken. A bit was set in the confusion 
status word for the source. 5 


Table V.D.4 Order of Band-Merging 

Drop-Dead Band Order in which Other 

Bands Were Searched 


(nm) (pm) 


12 

25, 

60, 

100 

25 

10, 

60, 

100 

60 

25, 

100, 

12 

100 

60, 

25, 

10 


V-24 


When a candidate was merged with a drop-dead, the position parameters were refined immediately 
with the same algorithm as that used for in-band seconds-confirmation. This was done before resuming 
the merging search in the next band, so that subsequent candidates would have to be compatible with the 
position information of all detections merged up to that point. 

D.3.b Band Filling 

After the attempt to merge sources from the SC/NSCF buffers, if any bands remained empty, the 
non-seconds-confirmed (NSC) detections were searched for band-filling candidates. These were the detec- 
tions which failed seconds-confirmation without the alibi of a dead or noisy redundant detector. The 
tests were identical to those for SC/NSCF sources, except that the threshold used was just under 4. 1 a, 
which should reject one true case out of every 20,000. The NSC detections were used for flux estimates 
only; their position information was considered too risky to use for parameter refinement. If more than 
one NSC detection was acceptable for filling a given band, the confusion status bit for that band was set, 
and the best in-scan position match was kept. After the NSC buffers were searched, unused NSC detec- 
tions were passed on to the input file for use at hours-confirmation. 

Any bands which remained empty were filled with low signal-to-noise detections (Section V.C.8) or 
with upper limits based on the noise histories of the detectors which the source image crossed. If any 
detectors were certain to have been crossed, then the upper limit was based on the lowest noise of any 
such detector; otherwise the highest noise on any detector which might have been crossed was used. 

D.3.c Special Considerations Regarding Band-Merging 

Three special considerations entered the band-merging problem. These were deferred in the discus- 
sion above in order to minimize the complexity of the description. The first of these was concerned with 
the selection of the drop-dead detection from one of the SC/NSCF buffers. Rather than selecting the old- 
est detection in any of the four bands, a time offset of seven seconds was applied against detections in the 
100 pm band in order to reduce the amount of cirrus contamination. This compensated for the fact that 
the 100 pm band was the first to register observable point source detections, because of its location at the 
entrance to the focal plane. Without the time offset, cirrus detections had been able to claim nearby 
point source detections before they could be associated with their proper partners. Implementing the 
seven-second delay forced all multi-band sources to begin band-merging with drop-dead detections at 
shorter wavelengths, where cirrus was much less of a problem. 

The second special consideration related to the handling of SC and NSCF detections of the same 
source in the same band. It was not unusual for edge-overlap detectors opposite a noisy detector to parti- 
cipate in a triple-detection with the noisy detector. Usually this triple-detection mode was processed 
without any problem, but occasionally the three detections would fail the tests required for acceptance of 
the triple-detection mode. These tests attempted to prevent confusion of close sources, as described 
above under in-band seconds-confirmation. When these tests failed, it was possible for the source to be 
carried forward in both the SC form and the NSCF form. Experience showed that the SC form was prac- 
tically always a better representation of the source than the NSCF form, but half of the time the NSCF 
had the earlier time tag and hence had first choice of band-merging candidates. In order to correct this, a 
test was added after selection of a drop-dead to see whether it was NSCF and was followed closely by any 


V-25 


SC source in the same band. In such cases, the SC and NSCF sources were swapped in the buffer, so that 
the SC source would be processed for band-merging first. The time windows used were one second in the 
12, 25, and 60 pm bands, and 1.5 seconds in the 100 pm band. 

The thir d special consideration involved the use of position information of NSCF sources in param- 
eter refinement. With three dead detectors and several more with degraded noise properties, NSCF 
sources were too numerous to permit completely ignoring their position information. On the other hand, 
false alarms on detectors opposite dead or noisy redundant detectors had to be accepted as NSCF. The 
only partially distinguishing characteristic of false alarms was their tendency toward low detection correla- 
tion coefficients. A compromise was therefore developed which permitted the use of NSCF position 
information in parameter refinement provided that the detection had a signal-to-noise ratio of at least ten 
with a correlation coefficient of at least 0.98. The only exception to this rule was when no SC or 
qualifying NSCF sources were present, in which case all NSCF position information was used in parame- 
ter refinement. 

D id Focal Plane Geometry Analysis 

The mean in-scan position angle discrepancies for each pair of redundant detectors in a band were 
interpreted as errors in the focal-plane geometry as mapped through the optics onto the slcy. These errors 
could not not be distinguished from timing errors, because all data were taken at the survey scan rate, but 
separation into components was not necessary for survey analysis purposes. Whenever detections in 
different bands were merged and qualified for position refinement, the in-scan position discrepancies were 
computed for all possible band combinations. After sufficient data were obtained, these mean discrepan- 
cies were forced toward zero by modifying the model of the focal plane geometry. 

D.4 Known Source Correlation 

A measure of processing quality was obtained by tracking the progress of certain known sources of 
infrared emission through the various stages of the data reduction. These included most of the IRC and 
AFGL objects, along with about 25,000 K Stars and 2500 numbered asteroids. A dozen comets and the 
major planets within the viewing constraints were also included, as well as about 200 objects which were 
selected on the basis of potential interest to users of the low-resolution spectrometer. 

D.4.a Known Source Prediction 

The a priori fluxes in the IRAS bands for the IRC, AFGL, and K stars were estimated by blackbody 
curves run through measured flux points. The K star estimates were adjusted by global rescaling after a 
month’s accumulation of data. The fluxes for the solar system objects were obtained via standard ther- 
mal models. The low-resolution spectrometer objects had no flux estimates. 

For each survey scan, the known objects which were to be covered were predicted. The predictions 
involved the detectors which would be crossed by each source image, the times of these crossings, the flux 
that should be observed, and a measure of the probability that a detection would result. The scan param 
eters and the known positions of the sources were used with the detailed boresight pointing history to cal- 
culate the geometrical predictions, and once the detectors were identified, their sensitivities and the inter- 
sections of their slots with the image blur were used to obtain the expected flux values. The probability 
of detection was based on the expected signal-to-noise ratio and photometric error, along with the uncer- 


tainty in the reconstructed solar aspect angle, the clearance in the slot of the image center, and the width 
of the image blur. The probability increased with higher signal-to-noise ratio, as photometric error 
became less likely to thwart the detection process, and it decreased with smaller slot clearance, as cross- 
scan limit cycling and image blur width became more likely to prevent the necessary amount of flux from 
arriving at the detector. For the solar system objects, orbital position calculations were necessary. These 
were computed in heliocentric ecliptic coordinates at the time corresponding to the middle of the obser- 
vation period and transformed to spacecraft-centered position angles. The remainder of the task was the 
same as for inertially fixed sources. 

After each anticipated known source image had been mapped through the focal plane, the number 
of detections predicted for it was checked to determine whether the source could reasonably be expected 
to be detected. The flux predicted for each detector was required to be above a certain level 
corresponding to the detection threshold on that detector. At least one SC or NSCF prediction was 
required before any prediction at all was generated. If any prediction was issued, then the estimated 
fluxes in all bands, even those below threshold, were passed on to the correlation analysis, although no 
detectors were associated with any band for which the fluxes were expected to be below the detection 
threshold. 

D.4.b Correlation of Observations With Predictions 

Association of an observed source with a predicted known source was done strictly on the basis of 
position agreement using the non-Gaussian statistics described at the beginning of Section V.D. 1 (Fowler 
and Rolfe 1982). After an identification was made, a subsequent identification of the same prediction 
with another observation could still occur and have a chance to replace the earlier association if appropri- 
ate. In such a case, the two associations were compared, and only the better match was kept. Similarly, 
if one observed source was found to pass the tests for association with more than one prediction, only the 
best match was kept. 

If a predicted source was never matched to an observation, then the predicted probabilities of detec- 
tion were checked to see whether any were above 0.99 and, if so, a warning message issued. A given 
prediction was considered unable to acquire further matches if the latest observation processed was more 
than 41 past it in the scan direction, or if another prediction at least five positions further downstream 
was matched. 

D.4.c Known Source Flux and Position Discrepancy Analysis 

When final match decisions had been made for each predicted known source, those which were 
identified with observations were investigated for discrepancies in the flux and position. The various 
types of known sources mentioned above were all kept separated in the statistical analysis of the 
discrepancies. 

In-scan position discrepancies were used to compute the mean, variance, and statistical significance 
of this error. The information was used in the focal plane geometrical calibration. Only very small mean 
errors were found, and the variances indicated that the a priori errors were slightly conservative. In addi- 
tion, data were grouped in cross-scan cells in order to determine whether the mean in-scan error was a 
function of the cross-scan location of the detectors involved. This would reveal any significant rotational 


V-27 



misalignment of the focal plane about the optical axis; an upper limit of less than half a second of arc 
was found for the impact of this effect on source position reconstruction. 

Cross-scan position discrepancies were studied in three groups; sources observed only in one band, 
sources observed in more than one band, and sources containing any triple detections. These groups 
have distinctive ranges of cross-scan errors, and are listed above in order of decreasing uncertainty. For 
each group, the mean and variance of the cross-scan position discrepancies were computed. General 
agreement with a priori values was found. 

D.5 Overview of Hours-Confirmation 

The next level of source confirmation involved observations with time separations from 100 
minutes up to 36 hours. Most objects were observed on consecutive orbits, but the upper limit was 
chosen to accommodate recovery scans for areas of sky not covered satisfactorily on the first attempt for 
a variety of reasons (see Section III.D). Generally, hours-confirmation was run on goups of three succes- 
sive SOPs, although this rule was violated during the minisurvey (SOPs 29-44) and in two cases of survey 
restarts (SOPs 57-61, 256-257, and 265). 

The "drop-dead" approach was used again to select a source for processing. When the oldest source 
became 36 hours older than the most recently acquired source, then the oldest source became eligible for 
hours-confirmation processing. A maximum of four scans were allowed to participate in one hours- 
confirmation of one source. Match candidates were selected by taking all subsequent unconfirmed obser- 
vations within a coarse window centered on the drop-dead source. This window was 27.5 across in eclip- 
tic longitude and 10.3' high in latitude. When the drop-dead source was within 30' of an ecliptic pole, all 
unconfirmed sources within 30' of that pole were included in the candidate set. 

Both seconds-confirmed/band-merged and non-seconds-confirmed (NSC) sources were included in 
the processing, but drop-dead sources were taken only from the seconds-confirmed set until all eligible 
ones had been processed, after which all eligible NSC sources were used in time order as drop-deads. 

D.5. a Hours-Confirmation Decision 

The candidates were required to come from orbits other than that of the drop-dead source, or else 
they were discarded from the processing of the drop-dead. They were also required to have detected 
fluxes (i.e., not band-filled fluxes) in at least one band in common with the drop-dead. Those remaining 
were then tested pairwise with the drop-dead for position agreement, and any which were not acceptable 
on this basis were discarded. Those that passed were grouped into sets belonging to individual orbits and 
checked for confusion. If no two remaining candidates were from the same orbit, then processing pro- 
ceeded to flux tests and combined flux/position tests. If two or more candidates from the same orbit had 
detected fluxes in at least one common band, then confusion was diagnosed, and status bits were set 
which would identify this condition to all downstream processing (Section V.D.8). If more than one can- 
didate remained from a given orbit, but no common bands were found, then the situation was not 
labeled as confusion, but only one candidate was retained; this was the one which passed the combined 
flux/position test with the highest score. 

If no candidates were found to be acceptable, the drop-dead source was rejected. Otherwise param- 
eter refinement was performed for the flux and position information. Tracking of known sources was 
performed, and various statistical computations were carried out. Candidates which were not confirmed 


with the drop-dead source remained for subsequent use, but confirmed candidates were not available for 
further hours-confirmation. 

D.5.b Position Agreement 

Each pairing of the drop-dead source with a candidate was tested for position agreement with the 
method mentioned at the beginning of this section. The position of the candidate was transformed to the 
local in-scan/cross-scan coordinates of the drop-dead, and the two-dimensional cross-covariance of the 
two position probability density functions was evaluated for the separation observed. If the result was 
below the threshold (Table V.D. 1), the candidate was not considered further for confirmation with the 
drop-dead. Otherwise the result was stored for use later, and the candidate remained viable. 

The threshold was set during simulation tests, and it was verified as acceptable during the processing 
of early survey data. As in the known source correlation processing, it was easier to arrive at a threshold 
experimentally than to attempt to accept a specific fraction of all true cases. This follows from the fact 
that the algorithm’s acceptance fraction of true cases increases as the position uncertainties decrease. 

D.5.c Photometric Agreement 

Each candidate which survived the position test was examined for photometric agreement with the 
drop-dead in all bands in which both sources had detected fluxes. A y} test was used which was based on 
the logarithmic discrepancies of the fluxes and the corresponding variances. The number of degrees of 
freedom was the number of common bands. The parameter tested was the complement of the cumula- 
tive probability; this was required to be above 1 x KT 4 , or else the candidate was discarded. No more 
than one true case out of every ten thousand should be rejected by this test. If the requirement was met, 
then a combined flux/position test was performed which required the product of the flux test score and 
the position test score to be above ten. Again the candidate was discarded or retained based on the out- 
come of this test, which also served as the tie breaker when more than one candidate from the same orbit 
passed all the tests. 

D.5.d Hours-Confirmation Confusion Processing 

When more than one candidate was confirmed with the drop-dead source, a series of pairings of 
each candidate with each other was performed, and the requirements placed on each pairing were to have 
at least one band with detected fluxes in common, to pass the position test, the flux test, and the com- 
bined flux/position test. When one of these requirements was not met, the candidate with the higher 
combined flux/position test score with respect to the drop-dead source was retained, and the other was 
discarded. 

D.5.e Hours-Confirmation Position and Photometric Refinement 

Position refinement was performed for the drop-dead source and all confirmed candidates by apply- 
ing the technique discussed by Fowler and Rolfe (1982) in pairwise fashion to the drop-dead source and 
the first candidate, then to the result of this process and the next candidate, and so on until all source 
observations involved had been processed. The method essentially computes the renormalized product of 
all of the position probability density functions and calculates the parameters which describe this product 
in terms an approximation to the original form of the density function. In this way, the near-optimal 
treatment of the non-Gaussian errors was maintained. 


V-29 


iUilillif 


Pairwise photometric refinement was performed in a manner analogous to the position refinement. 
Bands in which the paired observations had equal flux status were refined, otherwise the flux with the 
higher status was retained alone. Fluxes obtained by detection with the point-source template were taken 
as having the highest status; seconds-confirmed and non-seconds-confirmed both qualified. These were 
refined with Gaussian estimation applied to the logarithms of the fluxes, and so seconds-confirmed fluxes 
were usually weighted more heavily because of their smaller uncertainties. Low signal-to-noise ratio 
fluxes were refined by simple weighted averaging of the fluxes, with the signal-to-noise ratio values used 
as the weights. Upper limits based on noise were refined by retaining the lower value. All detector 
numbers involved were stored in arrays (see Section V.D.l) and kept with the confirmed source with the 
following exceptions. Detectors supplying noise fills which were not used were discarded. When four 
source observations were confirmed, only the first, second, and last set of detector number arrays were 
kept; this was done for purposes of space conservation, and a flag was set to indicate this condition. 

D.S.f Hours-Confirmation Statistical Processing 

Confirmed source positions were written to an output file for downstream tracking of confirmation 
frequency as a function of sky position, and histogram data were maintained to display confirmation fre- 
quency as a function of signal-to-noise ratio. A histogram counter corresponding to the highest signal-to- 
noise ratio in any rejected source was also maintained. 

Flux discrepancies in each band were processed to obtain the mean and the standard deviation as 
functions of signal-to-noise ratio; this was done in a format similar to that of seconds-confirmation, and 
was used in the photometric uncertainty analysis to feed back a posteriori dispersions for use as subse- 
quent a priori errors, y 1 tests on the fidelity of the photometric error modeling showed that this process 
was working as expected. Only confirmed sources devoid of confusion, outer-slot confinement, and other 
questionable symptoms were used to generate the photometric dispersion data. The standard deviations 
were similar to those of seconds-confirmation but slightly larger, indicating more power in the pho- 
tometric error spectrum at orbital frequencies. 

Position discrepancies were examined for in-scan mean error, variance, and statistical significance, 
and the cross-scan discrepancies were studied as functions of latitude. The solar motion caused the scan 
overlap from one orbit to the next to be a function of latitude, and the overlap was designed to minimize 
the total cross-scan error of single-band sources. This anticipated effect was confirmed, and this 
strengthened confidence in the fidelity of the error modeling. Slightly larger discrepancies were found 
very near the ecliptic, and these were probably due to the confirmation of asteroids which moved slightly 
between observations but not enough to preclude confirmation. This effect was not included in the pre- 
launch estimation of position errors at hours-confirmation. The primary effect, larger cross-scan errors 
for single-band sources at latitudes between 40° and 50° from the ecliptic, was clearly present. 

Histograms of the threshold parameters were generated to show their distributions. The main con- 
cern was for the cutoffs to lie in regions where variations in value did not cause significant variations in 
the number of confirmations and rejections. This would be possible if the noise processes were not 
saturating the decision process, and if the cutoffs were placed well above the values produced by most 
correct matches and well below the values produced by most false matches. This appeared to be the case. 

V-30 


I IN 


Confirmation and rejection of numbered asteroids were tracked and printed in the summary. The 
fraction of these which survived hours-confirmation was about half, as predicted before launch. 

D.5.g Special Considerations Regarding Hours-Confirmation 

Some complications to the hours-confirmation processing were omitted from the discussion above 
in order to limit its intricacy. These provisos will be described briefly in this section. 

As mentioned in the discussion of seconds-confirmation, sometimes a source was observed only on 
detectors which lay at the cross-scan boundary of the survey array. It was not possible on a single pass to 
determine whether the source actually missed the detector slots, so that both its flux and position infor- 
mation would be incorrect. The position uncertainties were expanded as described above, as this could 
be done in a way which virtually guaranteed bracketing the true position; there could be no such 
guarantee for the flux. As a result, the flux test was not performed when at least one of the sightings 
involved only outer-slot detectors in all bands containing detected fluxes. Flux refinement was also 
bypassed in these cases unless both sightings being processed were exclusively on outer slots; when this 
happened, the refined flux kept the status of being outer-slot only. If a subsequent pairing involved a 
sighting with detections inside the survey array cross-scan boundaries, its fluxes were retained without any 
averaging, and the outer-slot status was removed. 

When any source was processed which was detected but not seconds-confirmed in the 25 pm band 
next to the gap left by the demise of detectors 17 and 20, special checking was performed. If detections 
in other bands were present, the detector geometry was examined to determine whether actual image 
center passage between the cross-scan limits of the 25 pm band detector could be verified; if so, no 
further special action was taken. Otherwise the flux in the 25 pm band was treated as an outer-slot flux, 
since the image center may actually have missed the slot involved in the detection. This disqualified the 
source’s 25 pm flux from testing, inclusion in the photometric dispersion analysis, and contributing to 
refinement unless no better data were available. 

Particle radiation posed a significant threat to the photometric accuracy, and sources observed in all 
four bands were more likely to be affected by it. The error caused by this source of noise produced a 
broad non-Gaussian component to the photometric dispersion which highlighted the fact that it did not 
conform to the assumptions underlying the x 2 tests used for flux agreement. A rule was implemented, 
therefore, which was applied when the sources tested had at least three bands in common and failed the 
flux test. In such a case, the single most discrepant band was removed from consideration, and the % 2 
test with one fewer degree of freedom was used. This exemption could be invoked only once per source 
pair. 

D.6 Overview of Weeks-Confirmation 

The last stage in the confirmation chain searched for observations of the same sources with time 
separations on the order of a week to six months. All hours-confirmed sources became input to the 
weeks-confirmation processor. Sources which were not confirmed were placed in the WSDB along with 
those which were, and no rejections were enforced until final catalog preparation. As each new regionof 
sky was processed, the new sources were put unconfirmed into the WSDB, where they remained until 
that region was covered again, at which time the sources from the earlier coverage served as candidates to 
confirm the newer ones. 


V-31 


As each hours-confirmed source arrived for processing, a coarse window on the sky was used to 
select candidates from the WSDB. These may or may not have been weeks-confirmed already. The win- 
dow was 10.3' square, and the candidates were required to be separated in time by at least 36 hours from 
the source being processed. No flux tests were performed. Only a position test was used, and if more 
than one candidate passed this test, only the candidate with the highest score was kept as a match. 
Because of the sequential testing, valid multiple sightings should have been confirmed pairwise as each 
new sighting entered the processing, removing the need to identify more than one correct candidate from 
the WSDB. When a choice had to be made, a counter was incremented for possible confusion. Typi- 
cally only about 3% of the confirmed sources were diagnosed as potentially confused at weeks- 
confirmation. 

If there were no acceptable candidates, the new source was placed in the WSDB. Otherwise posi- 
tion refinement and discrepancy statistical computations were performed. 

D.6.a Weeks-Confirmation Decision 

The confirmation decision algorithm was a position test applied to each pairing of the new hours- 
confirmed source with the candidates. The candidate position probability density function was mapped 
into the in-scan/cross-scan coordinate system of the new source. The test was based on the two- 
dimensional cross-covariance of the position probability density functions, evaluated at the observed 
nominal separation. The underlying principle was identical to the position test used at hours 
confirmation, and for relative twist angles of less than 4.6°, the same algorithm was employed. Other- 
wise, a generalized implementation was used which was extensively complicated by the need to handle 
arbitrary relative twist angles between the axes of symmetry of the density functions. The details of this 
generalized version are discussed by Rolfe, Otake, and Fowler (1984). When the value of the cross- 
covariance fell below the threshold, the candidate was released from consideration. 

D.6.b Weeks-Confirmation Position Refinement 

Confirmed source sightings were processed for refinement of the position parameters by renormaliz- 
ing the product of the probability density functions. This procedure was similar in concept to that of 
hours-confirmation, but was complicated by the generalization to arbitrary relative twist angles. When 
the relative twist angle was less than 4.6°, the hours- confirmation algorithm was used. For larger angles, 
the refined density function was sufficiently Gaussian that it was satisfactory merely to assign a value of 
one second of arc to the uniform-error component. 

D.6.c Weeks-Confirmation Statistical Processing 

As at hours-confirmation, histograms of the value of the threshold parameter for confirmed and 
non-confirmed sources were separately accumulated, position data for both types were output for down- 
stream analysis of the sky distribution of these events, and known source tracking was performed. 

Statistical summations for computing the means and variances of position discrepancies were main- 
tained over the SOP period and over periods of approximately 50 days. This was done for the observed 
sightings relative to each other, and separately for known sources relative to refined position parameters, 
where applicable. Good agreement between a priori error modeling and the observed dispersions was 

found. 


V-32 


D.7 Auxiliary Processing for Low Resolution Spectra 

All observations associated with known sources were processed for extraction of low-resolution spec- 
trometer data (Chapter IX). This was also done for any observations above a signal-to-noise ratio of 25 
in the 12 or 25 pm band, even if they were not associated with a known source. The path of the source 
image through the focal plane was checked for passage over the spectrometer aperture, and if this was 
found, then the crossing time and scan rate were inserted into the output source record so that extraction 
requests for sources which survived hours-confirmation could be issued. 

Weeks-confirmed source sightings for which spectrometer extraction requests had been issued were 
processed for index association. This involved generating a record which associated the name of the later 
sighting with that of the earlier one, and which provided some housekeeping information about the 
extraction requests. 

D.8 Flux and Confusion Status Words 

Two status words were used to track the progress of a source through all levels of the confirmation 
processing. Each wavelength band of each hours-confirmed sighting has associated with it a flux status 
(FSTAT) and a confusion status (CSTAT) word which describe the quality of the quoted flux and 
whether the processor detected any possibly confusing neighbor sources nearby. Both of these status 
words were used in deciding which sources were to be included in the final catalog (see Section V.H). 
FSTAT (Table V.D.5) tells whether in a given band the quoted flux is an upper limit, an unconfirmed 
detection(s) or a well-confirmed measurement. CSTAT tells whether a sighting was potentially confused 
by other detections at the times of seconds- or hours-confirmation or during band-merging. Tables 
V.D.3a,b lists values of CSTAT. 


Status 

Table V.D.5 Flux Status (FSTAT) Values 

Meaning 

7 

Hours-confirmed; detector number array 
reveals number of detections. 

6 

One seconds-confirmed detection that was 
never hours-confirmed but was associated 
with at least one low-signal-to-noise detection. 

5 

One seconds-confirmed detection that was 
never hours-confirmed or associated with 
any other detection. 

4 

Two or more NSCF detections hours-confirmed 
and possibly including NSC detections. 

3 

At least one NSC detection associated with 
no more than one NSCF detection. 

2 

One NSC or NSCF detection and possibly any 
number of low-signal-to-noise detections (S/N > 2.9) or 
at least one low-signal-to-noise detection (S/N < 3) 

1 

Lowest three-sigma noise value found 


V-33 



D-9 Conversion of Position Uncertainties to Gaussian Approx imation 

Throughout the processing, the position error model was based on the non-Gaussian formalism 
mentioned in Section V.D.l. Because most of the point source position probability density functions 
evolve toward Gaussian shapes at each stage of refinement, the catalog description was designed to 
approximate the more familiar error ellipses of two-dimensional Gaussians, while retaining most of the 
accuracy of the non-Gaussian model. This involved bracketing the nominal positions on each coordinate 
axis with a single number to be interpreted as a confidence interval on that axis. This is a simplification 
of the optimal cross-scan description, which requires two numbers to define the distribution about the 

mean. 

For sources with strongly non-Gaussian position uncertainties, there is no way to escape the fact 
that some aspects of the error behavior will not generally be Gaussian. An important specific example is 
the relation between error in sigma units and the fraction of all cases exceeding that error. Because of the 
decision to quote the axes of the two-dimensional contour enclosing 95% of the probability mass, the 
error bars given were selected to match the corresponding deviations on the axes of the random variables 
used in the error model. Thus the expected frequency at which the true source position lies outside of the 
quoted contour is 5%. The contour is not generally elliptical, however, as this is a Gaussian feature not 
necessarily shared by the actual distributions. The error incurred by assuming that the contours are all 
elliptical, on the other hand, should not be extreme, and in fact this approximation played a role in map- 
ping the two-dimensional 95% confidence into one-dimensional confidences on each axis. 

The error bars were obtained as follows. For Gaussian errors, the two- dimensional 95% confidence 
would correspond to a * 2 probability of 0.95 with two degrees of freedom. This implies a value of 5.99 
for the x 2 random variable, so that the ellipse crosses each axis at 2.45 sigma, corresponding to a 
confidence of 98.6%. This was used for the confidence level for both axes. On the Gaussian axis of the 
error model, the interval used was 2.45 a y . On the non-Gaussian axis, a lookup table was used to obtain 
the 98.6% confidence deviation as a function of the ratio of the uniform half-width, L z , to the Gaussian 
a z . This table was computed by numerical quadrature for a grid of these ratios ranging from 0. 1 to 30. 
Below 0.1, the probability density was considered Gaussian, and above 30 it was taken as purely uniform. 

E* Overview of Small Extended Source Data Processing 

The small extended source processor used detections at scales larger than that of point sources to 
locate and measure stationary, unconfused sources resolved by the detectors but smaller than 8 . The 
identification of candidates was similar to that of point sources, in that the data stream of each detector 
was examined using zero-sum square-wave filters. The detections were checked for seconds-confirmation, 
then aU the detections in the hours-confirming coverage were assembled to construct a model of the 
source from which estimates were made of the source parameters. Before weeks-confirmation, an effort 
was made to discard sources that were either confused, larger than 8' or fragments of larger structures. 
Sources were then weeks-confirmed and, when possible, band-merged. A set of criteria described in the 
introduction to the Small Extended Source Catalog was used to select sources for inclusion in that cata- 
log. 


V-34 



E. 1 Potential Detections 


To locate potential small extended sources the data stream of each detector was first successively 
compressed by adding consecutive pairs of data samples. Designating the original data stream of each 
detector to be of level 0, new data streams of higher levels were generated according to the rule that the 
jth data sample in the kth level data stream was given by, 

X{j,k) - X(j,k- 1) + X(j+\,k-l) (V.E.l) 

Potential detections in any data stream were found by applying a narrow-bandpass, digital filter at 
each point in the data stream. This filter consisted of an eight-point, zero-sum, square-wave which sub- 
tracted the first and last two data points from the middle four (Eq. (V.C. 1)). When the filter passed over 
a source it gave rise to a characteristic peak, and thus picked out structures whose width spanned approx- 
imately four sample points. If we denote the spatial sampling frequency of the level 0 data stream by S 
(i.e. S is the number of data samples per arc min in the data stream of level 0), then the approximate 
width of an object (in arcmin) detected by the square-wave filtering the kth level data stream is given by, 

W(k) - (4/5 1 ) x (2 k ) (V.E.2) 

The spatial sampling frequency, S , was approximately 4 in the 12 pm and 25 pm wavelength 
bands, 2 in the 60 pm wavelength band, and 1 in the 100 pm wavelength band. Since it was not 
intended to detect structures greater than 8' in extent, the 12 pm and 25 pm detections were retained 
from levels 1, 2 and 3, corresponding to characteristic square-widths of 2', 4', and 8'. At 60 pm the level 
1 and 2 data streams were searched for sources 4' and 8' in size while at 100 pm only level 1 detections 
were kept for small extended sources. 

A peak in the kth level data stream which was picked out by the square-wave filter was accepted as 
a potential detection if its amplitude was greater than 3 times the local noise appropriate for that data 
stream. The amplitude of a square-wave peak was defined as 

A(k) = £7(2*) (V.E.3) 

where E is the maximum value attained above zero in the square-wave filtered data stream. The noise 
was estimated for each data stream as described for point sources. 

Since a source typically produced detections in more than just a single data stream for a given 
detector, each candidate detection from a level k data stream was compared with all candidates from the 
level (k + 1) data stream of the same detector that were within a distance of five level k samples. In such 
a comparison the detection with the largest square-wave peak was chosen as the best representation of the 
source, and other detections were rejected. To reject sources larger than 8’, a comparison was also made 
with the amplitude on a scale of 16'. If, as a result of this comparison, the accepted detection had a 
characteristic size of 16' (corresponding to a level 4 detection at 12 pm and 25 pm, a level 3 detection at 
60 pm, and a level 2 detection at 100 pm), then this source was rejected. Similarly, since point sources 
were also to be excluded, a comparison of each extended source detection was made with any level 0 
detections (point source) in the neighborhood. If the level 1 detection was found to have a square-wave 


V-35 


amplitude less than 0.75 times the level 0 amplitude, then the level 1 detection was discarded. The value 
0.75 proved to be optimal for the discrimination between point-like and extended sources and was 
evaluated after processing data from the Large Magellanic Cloud, where many small extended sources 

could be clearly seen. 

E.2 Seconds-Confirmation 

The focal plane array was designed so that the scan track of each stationary source crossed at least 
two detectors in each wavelength band each time the focal plane array passed over the source. All poten- 
tial detections which passed the comparisons described in Section E.2 above were checked for seconds- 
confirmation, and those which did not seconds-confirm were rejected unless the detector with which they 
would have confirmed had been declared failed. The latter condition arose most frequently at 25 pm. 

E.3 Source Construction and Hours-Confirmation 

In point source processing each pair of seconds-confirming detections was merged into a single 
detection. This was not the case for extended sources, nor was there any explicit check analogous to the 
point source hours-confirmation processing. Each individual seconds-confirmed detection was passed to a 
processor which attempted to piece together detections from the same band and the same hours- 
confirming coverage into a single source. Detections were considered to be from the same hours- 
confirming coverage if they were observed within 36 hours of each other. In most cases this process 
resulted in an hours-con firmed source in the point source processing sense, i.e., the resultant source was 
made up of two or more pairs of seconds-confirming detections whose detection times differed by a few 
hours. However, it was also possible for two or more seconds-confirming detections without a subse- 
quent hours-confirming sighting to result in an acceptable source. 

The piecing together of sources involved several stages. First, seconds-confirmed detections which 
lay sufficiently close to one another on the sky and which had detection times within 36 hours were 
linked together. If a detection failed to link with any other detection, it was rejected and not used in sub- 
sequent processing. Each group of linked detections was mapped onto a rectangular grid, representing 
their distribution on the sky. Each pixel in the grid was a l' square. The size of each detection was just 
the detector size in the cross-scan direction, and W(k) in the in-scan direction (Eq. (V.E.2)). This grid 
consisted of a number of cells which may be referred to by the coordinates (K,L). The number of detec- 
tions, N(K,L), which overlapped each grid cell was found, and each cell was assigned a weight, W(K,L), 
which was proportional to the sum of the intensities in each of the detections overlapping the cell. The 
intensity in a detection was taken to be the amplitude of the square-wave peak divided by the detector 
aperture area. The flux from the source seen by each of the detectors was calculated. This flux was then 
distributed among the grid cells overlapped by the detector according to the weights W(K,L). The total 
intensity associated with each grid cell was then the sum of the contributions from each overlapping 
detector, divided by the number of such detectors. 

Having assigned an intensity to each grid cell, the cells which were regarded as contributing to the 
source were identified. The cell with the maximum intensity was located. If this lay close to the boun- 
dary of the grid then no source was identified. Otherwise all the cells near to the maximum intensity cell 
which had an intensity above 1% of the maximum were linked into a single source; if the intensity distn- 


V-36 



bution had a minimum boundary and then a rise beyond the minimum, the linking process was stopped 
at this minimum. 

After all the cells within the source were identified, a check was made to see if most of the source 
lay on the grid. This was done by requiring the intensity on all four boundaries of the grid to be less 
than 10% of the maximum intensity on the grid. If this was not the case then "source run ofF was said 
to have occurred. If, in addition to source run-off, the maximum intensities on either pair of opposite 
grid boundaries were sufficiently different, indicating the presence of a strong flux gradient, then the 
source was likely to be a relatively small bump on the side of a larger source, and no source was returned. 
If source run-off occurred, but there was no evidence of a strong flux gradient across the grid, then the 
source was accepted, but a flag was set to indicate that it was possibly a fragment of a larger source. 

The flux of the source was taken to be the sum of the intensities of all the grid cells which were 
linked into the source, multiplied by the grid cell area. The intensity-weighted sum of the number of 
detections contributing to each grid cell in the source was found. This was called the confirmation level 
of the source, and gives an indication of how well the various detections contributing to the source were 
put together to form a single source. Only those sources whose confirmation level was above a threshold 
of 2 were passed on for further processing. 

Sources which passed the confirmation level and signal-to-noise thresholds became the input for 
subsequent confirmation processing; they were referred to as intermediate small extended sources, and 
were treated as hours-confirmed sources. The centroid and extent of each of those sources was com- 
puted, the extent being characterized by the matrix of the second moments of the flux distribution. This 
matrix was diagonalized to yield the semi-major and semi-minor axes of the source, together with the 
angle between the semi-major axis and the ecliptic meridian. The number of detections in the source was 
also earned. A detection was counted if it contributed more than half its own flux to the source. 

E.4 Cluster Analysis Processing 

Before weeks-confirmation was attempted, it proved necessary to verify that the candidates were 
well-defined and relatively isolated. 

When the focal plane passed over an object more extended than 1(/, structures within the object 
often gave rise to many detections whose properties depended on the scan direction. The simple detec- 
tion algorithm grossly misrepresented such large sources, since only the flux at the spatial frequencies of 
the square-wave filters was detected. This value could be far less than the true total flux. Large extended 
objects were characterized by strings and clusters of sources. This behavior was found particularly pro- 
nounced at 100 (im. 

A somewhat similar situation arose in regions with a high density of sources, or equivalently, 
because of complex backgrounds, such as the Galactic plane. The results from such regions were particu- 
larly sensitive to scan direction. Fluxes quoted in regions of complex structure should NEVER be taken 
at face value. 

Before weeks-confirmation was attempted, the sources were first analyzed by the cluster analysis 
processor, which looked for clusters of detections at the same wavelength and from the same hours- 


V-37 



confirming coverage. If the cluster was below a certain critical size then its constituent detections were 
merged into a single source which was then ready for weeks-confirmation. If, however, the cluster was 
larger than the threshold size, then it was deemed part of a large extended object, and the sources contri- 
buting to the cluster were rejected from any further processing. All detections which were not found to 
be members of large clusters were kept for subsequent processing. 

Clusters were identified by finding sources observed within 36 hours and which could be linked 
together spatially. A detection was considered to be linked to a cluster if it was a close neighbor to at 
least one other source in the cluster. Two detections were considered to be close neighbors if their link 
parameter, L, was less than a certain threshold. The link parameter, L, is given by, 

L - D/(r { + r 2 ) (V.E.4) 

where D is the distance between the centroids of the two sources and r, and r 2 are the radii of their 
equivalent-area circles, i.e. if a and b are the semi-major and semi-minor axes of the source, then the 
equivalent-area circle has a radius given by 

r - (V.E.5) 

The value of the threshold on the link parameter was taken to be 3.5 for reasons discussed in Sec- 
tion V.E.7.a. 

When the detections in a cluster were merged, the flux of the merged object was taken to be the 
sum of the fluxes of the constituent sources. The position was the flux-weighted centroid from all the 
constituent detections in the cluster. The extent of the source was obtained by combining the covariance 
matrices about the new centroid; the new matrix was then diagonalized to yield new estimates of the 
semi-major and semi-minor axes and twist angle. Relatively bright sources were not greatly altered when 
merged with fainter detections. 

The maximum semi-major axis that a merged cluster could have before being rejected was set at 5 , 
since extended source processing was not intended for larger sources and could not accurately estimate 
fluxes on size scales greater than about 8'. 

E.5 Weeks-Confirmation 

After cluster analysis processing, the remaining small extended sources were tested for repeatability 
on the weeks/months timescales of weeks to months. Only those detections found on at least two hours- 
confirming coverages were accepted as reliable sources. Two sources were regarded as being weeks- 
confirmed if, 

(a) they were in the same wavelength band; 

(b) they had detection times which differed by at least 36 hours; 

(c) they satisfied a position confirmation criterion detailed below; and 

(d) the ratio of the flux of the brighter detection to that of the fainter was less than 3. 

The link parameter for weeks-confirmation was taken to be the % 2 statistic of the two-dimensional 
position-match confidence of the two detections; in computing x 2 , each detection was modelled as a two- 


V-38 


III 



dimensional Gaussian distribution characterized by the extent matrix. Two detections positionally 
confirmed if their link parameter was less than 1.4 (see Section V.E.7.b). 

When a group of two or more weeks-confirming sources was identified, it was merged into a single 
weeks-confirmed source. The flux was taken to be the average of the fluxes of the two sources. The cen- 
troid and extent matrix of the weeks-confirmed source were calculated, and the extent matrix was diago- 
nalized to yield semi-major and semi-minor axes and the twist angle between the semi-major axis and the 
ecliptic meridian. 

If all the hours-confirmed sources contributing to a weeks-confirmed source positionally confirmed 
each other, then the source was said to be mutually confirmed. If an hours-confirmed source (denoted by 
A) weeks-confirmed with two or more other sources (denoted by Bl, B2, etc.) which did not mutually 
confirm, the following was done to extract the best weeks-confirmed source. First, a decision function, 

D, was evaluated for each of sources Bl, B2, etc., where 

D = (x 2 ) x (1 + X 2 ) + (R ~ 1) x R (V.E.6) 

and x 2 is the link parameter between sources A and B and R is the ratio of the brighter flux to that of 
the fainter. 

The source for which D was the smallest was found, and the ratio D/D min was found for each of 
the sources Bl, B2, etc. All sources for which D/D min was above a threshold of 2 were discarded. If 
only one source remained, then this was chosen as the weeks-confirming partner of source A. Otherwise, 
the source was considered confused and the candidates were rejected. 

E. 6 Band-Merging 

The weeks-confirmed, single-band sources were analyzed to see if they could merge with weeks- 
confirmed sources in other bands to become multi-band sources. Those sources in the four wavelength 
bands which positionally confirmed were linked together. The sources were tested for positional 
confirmation in the same way as in the cluster analysis processor, i.e. two sources positionally confirmed 
if their link parameter, x 2 as in Section V.E.5, was less than a certain threshold. If all the single-band 
sources contributing to a band-merged source confirmed with one another, then the source was said to be 
a mutually confirmed merger. 

The information describing each band-merged source became the basis for the small extended 
source catalog. Where two weeks-confirmed sources from one band competed to merge with a single 
source in another band, the band-merging attempt was abandoned and the various single band com- 
ponents were listed separately in the catalog. A flag was set to indicate that the sources could not be suc- 
cessfully band-merged. 

As an additional check on previous processing steps, the band-merging processor was allowed to 
look for possible band-mergers between weeks-confirmed sources from the same band. A total of 166 
cases were found. Upon inspection, they were all found to occur in areas with at least four hours- 
confirming coverages, each of which had yielded an hours-confirmed detection. These four detections 
formed a pattern that generated two weeks-confirmed sources; the failure of all four to link together could 


V-39 


be due to scanning structures in different directions, flux estimation errors, pointing errors, or to artifacts 
of the processing. The two weeks-confirmed sources obtained were naturally larger than the individual 
hours-confirmed sources by a sufficient amount that the band-merging processor managed to link them. 

All these cases of weeks-confirming candidated discovered too late were treated just like cases where 
several sources from the same band competed to band-merge with a single candidate from another band, 
the "band-merging" attempt was abandoned, and the various single band components were kept and sub- 
jected to clean-up processing individually (see catalog introduction). 


E.7 Optimizing the Processor 

The intermediate file of hours-confirmed sources (in the restricted sense of Section V.E.3 above) 
accumulated as the satellite data were processed. Cluster analysis, weeks-confirmation, and band-merging 
were run repeatedly on this intermediate file to optimize the the thresholds in these processors. This sec- 
tion describes how the thresholds were arrived at and discusses the implications of the choices. 

It was determined from preliminary analysis that the threshold on the link parameter used in the 
cluster analysis processor would have to be larger than 2, and that the weeks-confirmation threshold 
would have to be in the vicinity of 1 . It also became clear that in the range of interest, the clustering 
threshold had the greatest influence on the output. This threshold was therefore optimized with the 
weeks-confirmation threshold held at 0.8; then the weeks-confirmation threshold was chosen. The final 
changes in confirmation did not affect the clustering enough to require more tuning. 

The goal in optimizing the processor was enhanced reliability; completeness was only a secondary 
concern. Reliability included the requirement that the source be free of potentially confusing neighbors. 
Regions of low source density (such as high Galactic latitudes) were the prime areas where the processor 
was expected to perform well. 

E.7.a Choosing the Clustering Threshold 

As stated earlier (Section V.E.4), cluster analysis was meant to filter out fragments of sources that 
were larger than 8 f and sources that were confused. 

Neither problem was common at high Galactic latitudes, and while it was necessary to apply cluster 
processing in these areas, it was not possible to select an optimal threshold by studying these areas alone. 
Figure V.E.l illustrates this point clearly by showing that the number of weeks-confirmed sources at high 
Galactic latitudes was essentially independent of the clustering threshold at 12 pm and 25 pm; at 60 pm 
and 100 pm the number of sources dropped as the clustering threshold increased, indicating the presence 
of complex structure at these wavelengths. Figure V.E.l displays the results of processing in a region 
(henceforth Region A) at high Galactic latitudes, defined in ecliptic coordinates by 0° < P < 90° and 
135° < A. < 205°; its total area was about 4010 sq. deg, about 10% of the sky. The source density was on 
the order of 0.02 to 0.05 per sq. deg, too low for confusion to be a problem. 

Figure V.E.2 displays the number of weeks-confirmed sources as a function of clustering threshold 
for Region B, which includes a crowded portion of the Galactic plane. It was defined in ecliptic coor- 
dinates by 25° < (5 < 45°, and 280° < X < 300°; its total area was about 326 sq. deg, which implies a 


V-40 


400 



CLUSTER PROCESSING THRESHOLD 


Figure V.E.l This shows that cluster analysis processing does not greatly affect the number of 

small extended sources that are weeks-confirmed at high galactic latitudes. Only at 
100 jim is there a substantial dependence, because of the cirrus. 

source density of about 0.3 to 6.0 per sq. deg. Both figures V.E.l and V.E.2 were obtained with a weeks- 
confirmation threshold of 0.8. 

The shape of the curves in Figure V.E.2 suggested that a natural choice of clustering threshold could 
be based on keeping the final source density near the confusion limit. The confusion limit was obtained 
by requiring a minimum number of 25 beams per source (as in Section V.H.6), which corresponds to a 
probability less than 0.1% that two sources will be found in the same beam, and a probability of about 
1% that two adjacent beams will both have sources in them, assuming Poisson statistics. To use this cri- 
terion, an estimate was needed for beam size. A close upper limit was simply the in-scan width of the 
largest detector template (about 10') times twice the cross-scan width of a detector (about 10 f ), leading to 
an effective "beam size" of 1/36 sq. deg. 

In practice, however, and especially at 12 and 25 qm, detections on smaller templates were com- 
mon, and the effective beam size was found to be smaller than the upper limit by a factor two or more. 
To estimate the effective beam size, the average density of small extended source detections per survey 
coverage per band was found in the five most crowded bins on the sky. Each bin was approximately a 
sq. deg in size. This density was the same at 12 and 25 jim, namely 77 sources per sq. deg, with a popu- 
lation dispersion of 4; at 100 |im, the density was 40 ± 5 per sq. deg At 60 ^m, the average density in 
the five most crowded bins was 78 ± 20; if the highest-density bin was thrown out, the average in the 
next five was 68 ± 10. The result at 100 qm was quite close to the upper limit estimate above, as 
expected, since only the largest template was available at 100 \im. 


V-41 




Figure V.E.2 In contrast with Figure V.E.l, regions of high source density are heavily affected by 

cluster analysis. 


Adopting as effective beam sizes 1/40, 1/68, 1/77 and 1/77 sq. deg, respectively in the 100 pm, 60 
pm, 25 pm and 12 pm, the critical densities are 1.6, 2.7, 3.1 and 3.1 sources per sq. deg. To find the 
corresponding critical clustering threshold, three small heavily populated windows within Region B, with 
a total area of 9.8 sq. deg were used. The average density of weeks-confirmed sources dropped quickly as 
the clustering threshold increased from 2 to 3, and then leveled off, in a way similar to, but steeper than, 
what was seen in Fig. V.E.2. The critical source density was reached in all bands for thresholds between 
3 and 4. A value of 3.5 was chosen for all four bands. 

E.7.b Choosing the Weeks-Confirmation Threshold 

Figure V.E.3 shows the number of weeks-confirmed sources as a function of the weeks-confirmation 
threshold for Region A. Clearly, almost all confirmations were acquired by a threshold of 2; the slow rise 
beyond that point was roughly linear, as expected for false confirmations. The linear rise with threshold 
was expected because the search area (rather than the search radius) scales linearly with the confirmation 
threshold. Reliability dictated minimizing these false confirmations, while completeness demanded keep- 
ing as many of the better positional matches as possible. A value of 1.4 for weeks-confirmation threshold 


V-42 





Figure V.E.3 


Effect of weeks-confirmation threshold on the number of sources. The chosen thres- 
hold is indicated by the vertical broken line. 


was selected because it marked the boundary between the steep climb due to true confirmations and the 
gradual climb due to the false confirmations. 

E,7.c Choosing the Band-Merging Threshold 

Because both spatial resolution and source properties changed with wavelength, an astronomical 
object could appear extended in one band and point-like in others. In view of that, band-merging was 
carried out after confirmation, in contrast to point source processing. 

Figure V.E.4 shows the output of the band-merge processor as a function of the band-merging link 
parameter threshold in Region B (low Galactic latitudes). As anticipated, most sources turned out to be 
single-band sources. Past a threshold of 1.4 very little new band-merging took place. A threshold of 1.4, 
the same as for weeks-confirmation, was adopted. 

E.7.d Summary and Discussion 

It was evident that the performance of the small extended source processor at high Galactic lati- 
tudes varied slowly as a function of the cluster processing threshold. In contrast, crowded regions pro- 


V-43 



LOW GALACTIC LATITUDES 
AREA ■ 326 sq. deq. 
CAP THRESHOLD ■ 3.5 



BAND MERGING THRESHOLD 


Figure V.E.4 


The optimal threshold for band-merging is indicated by the vertical broken line. It is 
the same as for weeks-confirmation. 


vided the testing ground for selecting an optimal clustering threshold. With a first determination of 3.5 
as the clustering threshold, high Galactic latitudes provided the optimal choice of 1.4 as the weeks- 
confirmation threshold (Fig. V.E.2). A value of 1.4 was also chosen as the optimal threshold for band- 

merging. 

The final iteration was to repeat clustering optimization using the final choice for confirmation 
threshold. This was done using Fig. V.E.5, where the density of weeks-confirmed sources is shown as a 
function of the clustering threshold in the three crowded regions mentioned in Section V.E.7.a. The con- 
fusion limits were 3.1, 3.1, 2.7, and 1.6 sources per sq. deg in the 12, 25, 60, and 100 pm bands, respec- 
tively. This critical density was reached for all bands between clustering thresholds of 3 and 4; as 
expected, a value of 3.5 was still the optimal common choice for all bands. 

To assess the significance of this choice, one can estimate the size of the area searched for close 
neighbors by the cluster analysis processor both in relative and absolute terms. 

If an extended source is thought of as a square-wave in one dimension with total width W, then its 
corresponding rms size is W/{2 x 3° 5 ). The template used for detecting this source would have been 
itself square-wave shaped with a width W, and baseline segments W/2 on each side. A clustering thres- 
hold of 3.5 implies that 2 sources are considered close neighbors as soon as the baseline segments of their 


V-44 








CLUSTER PROCESSING THRESHOLD 

Figure V.E.5 A final check on the optimal thresholds: the weeks-confirmation threshold used here 

is the final one ( 1 .4); the effects of cluster analysis are quite drastic, as expected. The 
tick mark on each curve indicates the critical source density; in all cases this density 
is obtained at thresholds greater than the optimal choice of 3.5. 


respective detection templates start to overlap. This was clearly a reasonable, though somewhat conserva- 
tive, way of guarding against confusion. 

To estimate the angular distances involved in cluster analysis, the mean size of a sample of 1 1 1 
sources in each band was calculated after clustering and weeks-confirmation; these mean sizes (always 
close to the medians as well) were 1.5', 1.5', 1.8', and 2.2' at 12, 25, 60, and 100 pm. The largest size in 
any band was 3'. On average, therefore, cluster analysis treated as "close neighbors" two sources within 
10' of each other at 12 and 25 pm, 12' at 60 pm, and 15' at 100 pm. 

Cluster analysis fulfilled its objective in recognizing and setting aside large structures that were frag- 
mented into small extended sources; this was the reason for the decrease in 100 pm and 60 pm source 
counts with increasing clustering threshold in Region A (Fig. V.E. 1): cirrus was integrated into larger 
structures and dropped from further processing. It should be stressed, however, that cirrus is not absent 
from the small extended source catalog. 


V-45 



Cluster processing also fulfilled its objective as a confusion processor, as shown by the reduction by 
an order of magnitude of the source density in crowded areas (Fig. V.E.4) as the clustering threshold was 
varied from 1 to greater than 3.5. The sources that survived in densely populated areas were either very 
isolated or locally dominant. Isolated sources had no neighbors within the search window. Dominant 
sources were so much brighter than their neighbors that when the latter were combined with them the 
source parameters were barely altered, so that the size in particular did not grow beyond the maximum 
cutoff value. In confused areas most sources were dropped because they combined with a neighbor 
within the search radius, but far enough away that the combined structure exceeded the size limit. Such 
occurrences were recognized by the rejected source having an axial ratio much larger than unity. 

It should, therefore, be stressed that the absence of a small extended source where one was expected, 
in crowded or uncrowded regions, may be due to the presence of a neighbor; the two sources may have 
combined into too large a source. 

Table V.E. 1 traces the number of small extended sources that were processed through clustering 
analysis and weeks-confirmation with the final choice for the thresholds in Regions A and B. The frac- 
tion of sources that survived cluster analysis and went on to weeks-confirm was much higher in Region A 
(high latitude) than in Region B (low latitude). At 12 and 25 pm, about 90% of the sources surviving 
cluster analysis did not pass weeks- confirmation and were therefore discarded; this percentage decreased 
at longer wavelengths but remained substantial. The main reason for this high failure rate was the lack of 
a rigorous requirement for hours-confirmation, such as was required for point sources. Detector noise or 
other transients could trigger detections which seconds-confirmed, and then were used to construct a 
source that was discarded only at weeks-confirmation. 

The excess of 25 pm detections in Region A was a direct result of the lack of hours-confirmation: 
the dead detectors in this band relaxed the seconds-confirmation filter, and therefore allowed many more, 
stray detections than in other bands. The problem was hardly noticeable in Region B because most 
detections there were triggered by real but complex structure on the sky. That difference just reflects the 
contrasting definitions of Regions A and B: A has a low surface density of sources at the survey sensi- 
tivity, and the noise was dominated by the detector noise; B was dominated by confusion noise, in the 
sense that it was densely packed with detectable sources. The result was that most detections were dis- 
carded by cluster processing in Region B, and by weeks-confirmation in Region A. When all bands were 
combined it turned out that in both regions about one out of every seven detections ended up contribut- 
ing to a weeks-confirmed source; this fraction was remarkable similar for Regions A and B. 


E.8 The Small Extended Source Catalog 

Sources included in the small extended source catalog were obtained from the pool of weeks- 
confirmed sources and had to survive additional filtering to reject spurious or unreliable sources due to 
crosstalk, or to weak point sources combining with radiation hits and noise (see the catalog introducton). 

Associations were made between small extended sources and objects in other astronomical catalogs, 
following the procedure used for the point source catalog (Section V.H.9). The only difference was that 
the minimum search radius for all associations was 120". 


V-46 


Table V.E.la Processing Results in Region A (High Galactic Latitude) 


Band 

Number of 
Seconds/ 
Hours 
Confirmed 

Number 

Surviving 

Cluster 

Processing 

Number of 
Weeks- 

Confirmations 

NCOVR(2) 

NCOVR(3) 

NCOVR(>3)* 

12 jim 

2593 

2507 

88 

74 

13 

1 

25 jim 

3339 

3249 

101 

70 

31 

0 

60 jim 

2170 

1989 

196 

147 

45 

4 


1753 

1249 

209 

174 

30 

5 


Table V.E.lb Processing Results in Region B (Galactic Plane) 


Band 

Number of 
Seconds/ 
Hours 
Confirmed 

Number 

Surviving 

Cluster 

Processing 

Number of 
Weeks- 

Confirmations 

NCOVR(2) 

NCOVR(3) 

NCOVR( > 3)* 

12 jim 

5095 

2212 

350 

274 

70 

6 

25 jim 

4988 

2206 

364 

261 

95 

8 

60 jim 

8706 

2104 

364 

296 

65 

3 

100 jim 

4036 

1687 

369 

273 

86 

10 


* NCOVR(n) is the number of weeks-confirmed sources that were detected on n hours-confirming sur- 
vey coverages. 


F. Asteroids and Comets 

To achieve the objective of high reliability in the measurement of inertially fixed sources, the survey 
required that a source be repeatedly detected on several time scales (see discussion in Section V.D.l). 
This multiple confirmation process provided the means of detecting and rejecting objects moving at a 
variety of angular rates with respect to the inertial sky and the orbiting satellite. Detections which failed 
seconds-confirmation were due to radiation hits and infrared sources near the spacecraft, such as material 
emitted by IRAS itself, space debris, and Earth orbiting satellites. Failures at hours- and weeks- 
confirmation were used to detect comets and asteroids, effectively rejecting them from the catalog. 


V-47 





Solar system objects moving across the sky more rapidly than about l' per hour failed the hours- 
confirmation test. This test was used during the mission to search for fast-moving objects and resulted in 
the discovery of six comets, an extensive cometary debris trail, and two Apollo asteroids, one of which 
may be an extinct cometary nucleus. The details of the search for fast moving objects are given in Sec- 
tion ni.D.1. 

Asteroids and comets moving more slowly than 1 per hour would hours-confirm, and thus reside in 
the WSDB. To assess the efficiency of the weeks-confirmation filter, the coordinates of all hours- 
confirmed sources were positionally associated with the coordinates predicted from the orbital elements of 
asteroids numbered 1 through 2736, 12 periodic comets, and the 6 outer planets. The tagged sources 
enabled one to trace known solar system objects through the confirmation process. An analysis of these 
sources is given in Section VII.F. 

G. Extended Source Products 
G.l Processing Overview 

The basic approach to the compilation of the two-dimensional images and related time-ordered files 
was to select high-quality data meeting criteria based on observing conditions and the performance of 
individual detectors. Since extensive automatic confirmation tests were not applied as in the case of the 
discrete sources, internal consistency between different detectors within each survey observation was 
imposed by adjusting individual baselines and responsivities to produce the same mean brightness as the 
band average for that observation (the "destriping" operation). Weighted averages of the data were used 
to mosaic data from multiple survey scans of the same region into single digital images. Consistency 
between the scans forming an hours-confirming coverage was checked by human scanning of difference 
images j significantly discrepant data were deleted from the images. Finally, since the zodiacal emission 
toward any given direction on the celestial sphere depended upon the Earth’s orbital position and Sun- 
referenced observation angles at the time of the observation, data obtained in the three hours-confirming 
coverages of the sky were processed and presented separately. A file needed to reconstruct the observing 
geometry for each observation of a given region of the sky, the Zodiacal Observation History File, was 
produced. 

G.2 Quality Checking. Selection, and Weights 

Calibrated time-ordered detector data were used for the extended emission files if they had been 
obtained under acceptable observing conditions with respect to nuclear radiation event rate, adequate 
off-axis distance from strong potential sources of stray light, and satellite scanning rate. Data were 
deleted near the beginning and end of scans when the scan rate deviated from the nominal survey value, 
or if the detector signal saturated the analog-to-digital converter. 

The radiation blanking time introduced by triggering of the spike "deglitcher" in two detectors of 
each band was averaged over 2.5 sec intervals. Data were deleted from any band during any interval in 
which the dead time introduced by the deglitcher exceeded 10% in that band. 

The processing was designed to allow definition of a cone of avoidance for closest approach to the 
Earth, Moon, Mars, Jupiter, and Saturn. Stray light performance of the instrument was examined early 


V-48 


in the mission, and it was concluded that the nominal scan constraints on closest approach to the Earth 
and Jupiter produced data of acceptable quality for extended source processing, and that no constraint on 
data near Saturn or Mars would be imposed. No data within 30° of the Moon were included in the 
extended source products. This resulted in a series of lune shaped gaps spread across the ecliptic plane in 
each of the three sky coverages (Fig. III.C.6). 

A weight factor was assigned to each sample from each detector as a quality indicator for that sam- 
ple. The weights followed the samples through all subsequent averaging, image creation, and image 
mosaic processes, yielding a weight map corresponding to each intensity map produced. A general weight 
was assigned to data from each detector based on the nominal sensitivity and noise of that detector 
observed early in the mission. Inoperative detectors were assigned zero weight. The weight was propor- 
tional to the variance of the data sample, so that it could be used as a statistical weight in an averaging 
process. 

The extended source images were not intended to give high-quality photometric information for 
point or small extended sources. However, to maintain consistency for small sources, data from the two 
edge detectors in each band which were substantially smaller than nominal detectors in that band (detec- 
tors 26, 47, 39, 42, 11,31,4, and 55) were deleted. 

G.3 Phasing, Sorting, and Gaps 

In preparation for smoothing and binning the data for the various extended source products, the 
time-ordered calibrated data were adjusted on a detector-by-detector basis to align samples obtained at 
the same sky position of the telescope boresight. This phasing took account of both the scan rate of the 
telescope and of electronic phase shifts in the sampling of individual detectors. Detector data were also 
sorted into cross-scan order in each band according to their centroid positions. Because of the phasing 
operation, some data were lost at the beginning of each scan and whenever a gap in the telemetry stream 
occurred. The start-up time for the phasing was 1 1 seconds, or about 3/4° on the sky. 

G.4 Conversion to Surface Brightness 

Detector data were converted from calibrated in-band fluxes (W m -2 ) to surface brightness, (W 
m -2 sr -1 ), using an effective instantaneous field of view for each detector obtained from slow scans of 
point sources across the array (Table IV.A.l). 

G.5 Compression and the Time-Ordered Files 

The detector data were compressed in the in-scan direction prior to projection into the image 
domain. Compression was accomplished by filtering and then decimating in order to prevent aliasing. A 
symmetric filter was desired to provide zero phase shifts, and a Lanczos single smoothed filter ((sinjc/x) 2 
weighting in the window) was selected. The filter width was chosen for each band to provide two output 
samples per second, corresponding to a nominal 2 ' sample spacing. No compression or interpolation was 
applied in the cross-scan direction because the construction of the telescope’s focal plane array produced 
samples with 2' spacing. These compressed data formed the time-ordered input for the sky plates and 
Galactic plane maps. 

A lower resolution time-ordered file of intensity data was also created. The Zodiacal Observation 
History File (ZOHF) contains the data samples obtained by averaging data from all detectors in each 


V-49 


lit is Hi IIIU II I ;«tl I i m II JiUHIiltli.l -Mb 


band (according to their weights) over an interval of 8 seconds of time. This provides data samples at 
approximately 30' intervals in both the in-scan and cross-scan directions. The file also retains, in addi- 
tion to celestial coordinates, the Sun-referenced observation angles and time, and information necessary 
to evaluate the zodiacal emission contribution to the intensity in each sample (see discussion below). 
Data from this file were used to create the low-resolution all-sky maps. 

G.6 Destriping 

In the course of reducing the scan data into images, it was found that, in spite of the efforts to cali- 
brate the gain and offset for each detector individually, small inconsistencies remained between detectors 
in a given band within a single scan. Such systematic inconsistencies showed up clearly as stripes , i.e., 
persistent streaks in the in-scan direction in the images on the cross-scan scale of a detector length. A 
"destriping" procedure was developed to remove these high spatial-frequency inconsistencies while 
preserving the best available photometric information for the band as a whole. The procedure consisted 
of determining a responsivity and baseline correction for the data from each detector for each survey 
scan, and then applying that correction to the time-ordered data file prior to projection into images. The 
corrections were determined by requiring that the intensity from each detector follow the weighted band- 
average intensity during that observation. This procedure depended on the scan passing over areas of 
both low and high average background so that responsivity effects could be separated from baseline 
effects. Prominent small sources (high curvature regions) in the data were excluded in computing the 
averages. If insufficient data were available in a given observation, or if the correlation of the output of a 
detector with the band average was poor, the corrections determined in the previous observation were 
used. 

This process reduced the striping on a scale of a few arc minutes introduced by variations among 
the detectors in the focal plane. It did nothing to remove the 1/2° wide stripes caused by variations in 
calibration or sky brightness between different scans of the same piece of sky. The treatment of these 
wider stripes is discussed below under the topic of consistency checking in Section G.8. 

G.7 Projection into Sky Maps 

After phasing, smoothing, and resampling to produce the compressed data, the time-ordered detec- 
tor data formed the equivalent of a two-dimensional array 15 or 16 samples wide (depending on the 
band) and several thousand samples long. Exact positional information was carried for every other sam- 
ple of the two detectors at the edges of the array in each band, the "tie points". The map projection was 
accomplished by projecting the exact positions of the tie points into the line and sample space of the map 
images with the mapping transformation, and doing a bi-linear interpolation in line and sample space to 
get the projections of the other samples. All projections were done so that the sky coordinate associated 
with a pixel refers to the position at the center of the pixel. The map projections used are discussed in 
Section X.D. Once assigned a position in image space, a particular data sample was binned by dividing 
the weight of the sample among the nearest four bins with an inverse bi-linear interpolation, multiplying 
the sample value by the divided weights and accumulating the weighted samples into the four bins of the 
intensity image and the weights into the four bins of the weight image. Maintaining these separate 
weighted intensity and weight images facilitated combining data from multiple scans into a single image. 
The final average intensity image was produced by dividing the weighted intensity image, pixel-by-pixel, 


I II 


V-50 


by the weight image. These procedures applied to the processing of the low-resolution all-sky map data 
as well as the sky plate data. 

G.8 Consistency Checking and Removal of Bad Data 

After the data had been assembled into maps, the images were examined to check for anomalies. 
One coverage of the sky consisted of 1/2 ’-wide survey scans spaced every 1/4°. By selecting alternate 
scans, a coverage could be split into two sets of non-overlapping scans, each of which essentially covered 
the whole sky. The check consisted of examining the difference image between these two halves of a sky 
coverage. Anomalies discovered in this way were then classified and, if necessary, examined further in 
the two direct intensity images. If the anomaly was severe, the offending data were removed from the 
map. Three classes of anomalies were removed: 

1) Obvious flooding of the focal plane by near-field objects. 

2) Scans or sections of scans which were substantially noisier than their neighbors. The random 
noise level needed to be about two times larger than normal before this effect was noticeable, so 
relatively few data were removed for this reason. More frequently, the destriping strategy 
would fail and noise in the form of objectionable 2 -wide stripes would appear. Scans causing 
such stripes were rejected. 

3) The brightness of the scan differed substantially from the generally smooth background level of 
its neighbors. The criterion used for removal was a difference of more than 10% of the baseline 
value from neighboring scans on adjacent orbits. A substantial amount of data was removed 
for this reason. Particularly offensive in this respect were 25 pm scans made near edges of the 
South Atlantic Anomaly (SAA) where bias-boost (Section m.C.4) procedures were not used. 

Data removal was accomplished by adding the offending scan or section of a scan back into the 
map with the negative of its original weight. 

Several known residual anomalies remain from this cleaning process. No attempt was made to 
remove or even find small scale discrepancies between different scans of the same piece of sky, and many 
small, generally point-like, sources remain in the maps which DO NOT appear in the same place in two 
different scans. This point cannot be overemphasized; when using the sky maps to look at small 
sources, even bright ones, at least two separate coverages MUST BE COMPARED or THE USER WILL 
BE FOOLED. Many, maybe most, of the non-confirming sources are due to asteroids; however, they do 
occur in all parts of the sky, and caution must be exercised in all cases. 

Noticeable l/2°-wide stripes run across many of the sky plates. Regularly spaced patterns of these 
stnpes were due to the effects of non-boosted SAA crossings, as mentioned above. Isolated stripes were 
due to some undetermined calibration anomaly. Almost none of the stripes exceed the 10% difference 

criterion; however, an occasional stripe which differs by more that this has been left in to prevent creation 
of a hole in the map. 

It is not known at this time whether the scans deleted due to non-boosted SAA crossings, or their 
neighboring scans, were correct. The maps have been made flatter by the deletions of data, but the possi- 
bility of small, systematic DC calibration errors in the middle of scans, deemed unlikely, cannot be 
excluded. 


V-51 


No attempt has been made to exclude the effects of photon-induced responsivity enhancement in 
the data (see Section IV.A.8). For instance, passage across bright regions in the Galactic plane increased 
the responsivities of the 60 and 100 pm detectors for some time after passage. Similar effects near bright 
discrete sources are evident as well. Again, comparison among the three sky coverages should be made to 
avoid unrecognized problems due to these effects. All three coverages should be compared for this pur- 
pose, since the first two coverages generally passed over a given piece of sky in nearly the same direction, 
while the third coverage, where present, generally passed in a different direction. 

G.9 Final Man Image Generation 

The final steps in production of sky plate maps were the reduction to average intensity, conversion 
from in-band average intensity to brightness density (Jy sr -1 ) and application of final calibration data. 
Reduction to average in-band intensity was accomplished by division of the weighted intensity image by 
the corresponding weight image. The conversion to brightness density was based on the knowledge of 
the shapes of the filter passbands of the instrument and the assumption of a radiation source with an 
energy distribution flat in flux per octave. Small corrections modified the baseline and rescaled the result- 
ing brightness on a pixel-by-pixel basis to reflect the final calibration. 

Maps of the sky within 10° of the Galactic plane were made by remapping the pixels from the sky 
plates into Galactic coordinates using the cylindrical projection described in Section X.D. The remap- 
ping was done by the tie point technique described above, where the full map transformation was applied 
only to a subset of the pixels and linear interpolation is applied to the rest. Binning of a data sample into 
the nearest four pixels was also done as described in Section V.G.7. 

The low-resolution all-sky maps and the Zodiacal Observation History File were subject to the same 
modifications as the sky plates for final generation processing. 


H. The Point Source Catalog 

H. 1 Processing Overview 

After the Working Survey Data Base (WSDB) was completed by the sequential processing of every 
SOP, several programs processed that data base to create the catalog. 

The major steps were as follows: 

1. Final calibration corrections, 

2. Clean-up processing, which forced sources closer than 30" (in-scan) and 90" (cross-scan) to 
weeks-confirm, 

3. Point source and small extended source neighbor tagging, 

4. Cirrus (highly structured 100 pm emission) flagging, 

5. High source density processing, 

6. Variability tagging and average flux computation, 

7. Association and classification of low-resolution spectra, 

8. Associations with previously known astronomical sources. 


V-52 



9. Catalog source selections, 

10. Transformation of coordinates, assignment of position uncertainties and assignment of source 
names. 

Each of these steps is discussed separately below except for step 1. Step 1, incorporating the final 
calibration, is discussed in Section VLB. 

H.2 Clean-Up Processing 

An analysis of sources in the WSDB that were within several arc minutes of each other showed that 
roughly 5% of all sources did not have every possible hours-confirmed detection combined into one 
source. There were three reasons for this. First, sources scanned more than five times within one hours- 
confirming coverage were separated into two distinct hours-confirming detections due to an upper limit 
of four scans per hours-confirming detection imposed by the processing. Only one of these two detec- 
tions was allowed to weeks-confirm with other hours-confirmed detections already present in the WSDB. 
Second, sources that suffered band merging difficulties occasionally produced two hours-confirmed detec- 
tions out of one hours-confirming coverage. Only one piece was allowed to weeks-confirm. Finally, due 
to the way the minisurvey scans were scheduled, the first survey coverage was not allowed to weeks- 
confirm with some of the minisurvey sources. Each of these cases was sufficiently well defined to allow a 
clean-up processor to give these pieces a chance to weeks-confirm with each other. Additional problems 
were caused by radiation hits, source confusion and by multiple scans inadvertently spaced too close 
together in time to allow weeks-confirmation. These sources were cleaned up in a two-step 
process. First, all weeks-confirmed sources were given the opportunity to confirm using the standard 
threshold with any hours- or weeks-confirmed neighbors located within 30" in-scan and 90” cross-scan. 
Most sources that confirmed at this stage were ones that had earlier been denied the opportunity for 
purely technical reasons. A single source with a refined position resulted. 

The second stage of clean-up was to force all weeks-confirmed sources to merge with any neighbors 
located within 30" in-scan and 90" cross-scan, whether or not the sources were confirmable according to 
the weeks-confirmation decision process. No position refinement attempt was made. This step proved 
necessary because analysis of near neighbors showed that most were caused by the incorrect splitting of 
single sources in confused regions by the various confirmation processors. 

A separate processing problem was partially solved by the clean-up processor as well. Two normal 
survey scans were placed within one layer of the minisurvey (Section III.C.ll, Table III. 1). These scans 
were inadvertently processed twice; once as part of the minisurvey and once as part of the regular 
survey. Thus some sources (those missed by that layer of the minisurvey) had identical hours- 
confirming detections present as two separate entries within the source. One of these was deleted by the 
clean-up processor, although nothing could be done to remove the double weight given to that sighting in 
the position refinement process. 

H.3 Neighbor Tagging 

The relative isolation of a source provides the user with an indication of its quality. Associated 
with each source are the numbers of hours- and weeks-confirmed point sources and small extended 
sources. The search area used was a box 6 ' (half- width) in-scan, which was the largest detection shadow- 


V-53 


ing distance (Section V.C.7), and 4.5' (half-width) cross-scan, which was a detector length. Neighbors 
were tagged before the clean-up processor was applied. However, neighbors that were confirmed with the 
source during clean-up were not counted, nor, due to an error, were any sources within a box that had 
been successfully cleaned up. 

If no neighbors were found, it is unlikely that processing or confusion problems exist. The presence 
of neighbors should make the user check those neighbors to investigate the possibility that the source was 
extended, was in a confused patch of sky, or suffered some other problem. 

H.4 Cirrus Flagging 

At 100 pm the infrared sky is characterized by emission from interstellar dust on all spatial scales, 
known affectionately as “infrared cirrus". A significant chance exists that ANY catalog source has been 
affected by components of this long-wavelength emission on the point source scale. Four separate quanti- 
ties were derived for each source to assess the importance of contamination by cirrus. While these are 
imperfect flags (a source may be affected by cirrus without any of these flags being set, or may not be 
affected at all even at high values of the flags) they work for the majority of sources and provide a simple 
description of the sky at 100 pm in the vicinity of each catalog object. 

Flag 1 (CIRRI) is the density of WSDB sources (hours and/or weeks-confirmed) detected only at 
100 pm within a ±1/2° ecliptic coordinate box of the source position. Cirrus typically produces strings 
of such sources, giving rise to a high density. Weak cirrus, producing only a few 100 pm sources, may 
still affect the quality of catalog sources. 

Flags 2 (CIRR2) and 3 (CIRR3) come from the 100 pm extended source data at a resolution of 
1/2° (see Section V.G.5). Although the angular resolution is a serious limitation for these flags, since the 
angular scale of point sources is 10-30 times smaller, cirrus on the point source scale usually shows struc- 
ture on the 1/2° scale as well. 

CIRR2 was derived from a spatially filtered version of the 1/2° averaged 100 pm emission. The 
value of the filtered sky brightness at a point, J,, is given by 

7, - -0.5 x /,_! + /, -0.5 x I i+l (V.H.l) 

where adjacent measurements, /„ of the total intensity were separated by 1/2° in the in-scan direction. 
This filtered surface brightness was converted into a “cirrus flux", F c , by multiplying J, by the solid angle 
of a typical 3'x5' 100 pm detector. The cirrus flux was then compared with the 100 pm flux (or upper 
limit) of each catalog source, F s . The value of CIRR2 was scaled logarithmically to fit in the range 1 to 9 
according to 

CIRRI - (8/3) x log (F c /F s ) + (19/3) (V.H.2) 

A value of CIRR2-1 corresponds to the cirrus flux being less than or equal to 0.01 of the source 
flux while CIRR2-9 corresponds to the cirrus flux being equal to the source flux. As described in Sec- 
tion VII.H, CIRR2 less than about 4 is indicative of little or no cirrus contamination, while larger values 
probably indicate significant contamination. A value of CIRR2-0 indicates that no 100 pm extended 
emission data were available. 



Flag 3 (CIRR3) is the total intensity of the 100 pm extended emission in a 1/2° beam (MJy sr -1 ). 
High values are indicative of a large column density of interstellar dust. 

A fourth cirrus indicator is the presence of nearby small extended sources at 100 pm. Since the 
small extended sources have sizes much closer to point sources than the 1/2° data, this flag (SES1), which 
gives the number of hours-confirmed small extended sources, will normally be an accurate cirrus indica- 
tor. However, because of the higher thresholds used for this processing, weak cirrus may not be detected. 

As discussed in more detail in Section VTI.H, if any of these flags takes on a large value, one should 
be cautious. Cirrus can affect the point source flux in any band by causing band-merging difficulties and 
can certainly affect the quality of a quoted 100 pm flux. 

H.5 Ave rage Flux Computation and Variability Analysis 

In the computation of an average flux for each source, three levels of flux quality were recognized: 

1) High quality fluxes had no confirmation problems other than missing detections due to failed 
detectors, with at least one hours-confirming detection with no problems. For aficionados, at 
least two hours-confirming sightings with flux status, FSTAT -4 or 7, with at least one 7, were 
required (see Section V.D.8). 

2) Moderate quality fluxes were believed to be reliable, but were missing some detections, usually 
because the source was at the detection threshold. Two hours-confirming sightings with 
FSTAT-3-7 were required. 

3) Upper limits were given for bands lacking high or moderate quality fluxes. Measurements in 
these bands had no more than one measurement with FSTAT-3-7. The upper limits were 
derived from all of the measurements. An analysis showed that measurements in a given band 
with only a single detection (FSTAT— 2) within an hours-confirming coverage were severely 
contaminated by radiation hits and were unreliable. These single "detections" were used only 
as upper limits. Quoted upper limits are nominally 3 ct values. 

High and moderate quality fluxes were obtained by log-flux averaging of individual acceptable 
hours-confirmed fluxes weighted by the inverse of the log-flux variances. If the flux discrepancy flag was 
set (see below) and the source had more than two hours-confirmed fluxes, then one flux was removed 
from the averaging process. The rejected flux was the one farthest from the median flux in units of the 
error. The median flux was arbitrarily picked to be the larger of the middle two fluxes when there was an 
even number of fluxes present. The rationale behind this was that discrepant fluxes occasionally occurred 
because of radiation hits, a passing asteroid, a failed detector, or processing problems. With more than 
two measurements it was possible to weed out the discrepant flux. For quoted fluxes, a x 2 test was 

applied to each band. If the value of the reduced % 2 was greater than 9, the flux discrepancy flag was set 
in that band. 

Uncertainties quoted for high and moderate quality fluxes were the maximum of: a) the uncer- 
tainty of the mean derived from the accepted fluxes; or b) the uncertainty of the mean derived by pro- 
pagating the quoted uncertainties of the individual, accepted flux measurements. 

Fluxes quoted as upper limits were obtained from the the largest flux (or limit) for sources with 
only two hours-confirmed sightings and from the median flux for sources with more than two 


V-55 


detections. The median was taken as the largest of the two middle fluxes for sources with an even 
number of observations. The flux discrepancy flag was set if the largest and smallest flux from all the 
limits for low quality detections differed by more than a factor of 3. 

Sources with high or moderate quality fluxes at both 12 and 25 pm were examined for 
variability. A percentage probability of being variable was quoted for sources whose fluxes at BOTH 
12 and 25 pm either increased or decreased significantly from one hours-confirmed sighting to another. 
In other words, the changes in flux at 12 and 25 pm had to be correlated. Although variability could 
have been detected in other ways, this method led to a simple and reliable determination of variability of 
a large number of sources. For sources with more than one pair of sightings, the value determined for 
the pair of sightings with the greatest likelihood of variability was given. 

The probability that a source brightness varied was calculated by comparing the number of sources 
with correlated flux excursions exceeding m a at 12 and 25 pm, N corr (>m), with those sources showing 
anti-correlated flux excursions exceeding mo in both bands, N anti {>m). The probability that a source is 

variable is given by 


NcorA>rn) _ ^(> m ) 
Ncorr(^°) Ngnti(^o) 

p{m) “ NcorA>rn) + A U(>™) 

Ncorr(^°) Nantii^O) 


(V.H.3) 


This approach was adopted because there are many ways in addition to true variability that a 
discrepant flux could be obtained, including radiation hits, processing problems caused by failed detectors 
and confusion with transient spurious sources which would increase the flux in one band at one sighting. 
Thus, while Eq. (V.H.3) is not rigorous, it gives a good measure of the significance of any flux excursion. 
As discussed in Section VI1.D.3, the sources deemed likely to be variable (p>0.9) represent approxi- 
mately 20% of the 12 and 25 pm objects in the catalog. 

H.6 High Source Density Regions 

The density of sources in the WSDB exceeds the resolving capability of the instrument over some 
parts of the sky The nature of both the sky and the instrument make this a wavelength-dependent effect; 
at 12 pin the instrumental resolution was less than 1' and the sky was dominated by point sources, while 
at 100 pm the resolution was about 4' and much of the sky was dominated by highly structured diffuse 

emission. 

The overriding concern in developing a strategy for regions of high source density was to insure the 
reliability of the information presented in the point source catalog. Reliability, in this context, has three 
meanings: 1) a point source must be an inertially fixed, celestial source (the basic meaning of reliability 
throughout the catalog); 2) a source must represent a very compact structure that, to the extent possible, 
is not merely a fragment of a complex background; and 3) the intensity measurements reported in the 
catalog must repeatably represent the brightness of a source above the local background. 


V-56 


The first step in processing high source density regions was to determine whether the number of 
sources in a given wavelength band in a 1 sq. deg bin exceeded a "confusion limit". Sources in bins with 
more than the threshold number of sources were subject to additional criteria for inclusion in the catalog. 
The decision to apply the criteria was made independently on a band-by-band basis. Thus, a multi-band 
source might have high source density rules applied to its 100 pm measurements, but not to its 12 pm 
measurements. Once the high source density rules were invoked in a band, each source in that band in 
that bin was examined for the quality and repeatability of its individual measurements and for its isola- 
tion from other potentially confusing sources. The high source density processor determined whether the 
measurement of the source in that band was of high, medium or low (upper limit) quality, where these 
definitions differ from those in lower density areas, and calculated an average flux for the source in that 
band. Many sources were found wanting and were excluded from the catalog. 

The high source density selection criteria used only information contained within the WSDB to 
judge the relative merits of a source. The confusion processing algorithms were developed and tuned in 
two confused regions, one in the Large Magellanic Cloud (LMC) and one near the Galactic Center. 
Detector strip charts were examined for many of the sources to verify the validity of the selection criteria. 

H.6,a Location of High Source Density Regions 

The confusion-limited source density is determined by the instrumental angular resolution and how 
many detector beam areas are, on average, required per source for a reliable measurement. A conserva- 
tive limit of 25 beams per source leads to maximum source densities of 50, 50, 25 and 12 sources per sq. 
deg at 12, 25, 60 and 100 pm (assuming a detector area of the nominal in-scan detector size times the 
detector cross-scan width). Aitoff projection maps in Galactic coordinates showing the density of sources 
with moderate or high quality fluxes (according to the basic rules described in V.H.5) in a given 
wavelength band are shown in Figs, V.H.la-d. At 12, 25 and 60 pm the region within 3-5° of the Galac- 
tic plane and within ±100° of longitude of the Galactic center plus small regions in Orion, Ophiuchus 
and the LMC present the major problem areas. The 100 pm sky presents a more complicated picture 
because the regions of high source density cover large parts of the sky due to the infrared cirrus. 

To set the thresholds for high source density processing, the numbers of sources in 1 sq. deg and 
0.25 sq. deg bins were examined in the WSDB. It was determined that processing all 1 sq. deg bins con- 
taining at least 45, 45, 16, and 6 sources (at 12, 25, 60 and 100 pm) would clean up 90% of all 0.25 sq. 
deg areas with confusion-limited source densities. Thus, any 1 sq. deg bin (defined in ecliptic coordi- 
nates, cf. Appendix X.l) containing at least the threshold number of sources in a particular band was 
processed according to high source density rules in that band. No attempt was made to join together 
high source density regions into a few simply connected areas. A list of the high source density bins in 
each band is available with the machine readable version of the catalog. 

High source density rules were occasionally invoked for wavelength bands with fewer than the thres- 
hold number of sources in a bin if failure to include that band would have resulted in a non-adjacent set 
of bands being processed, i.e., a "spectral hole". Thus, for example, if the number of sources at 25 and 
100 pm, but not at 60 pm, exceeded the thresholds, the decision was made to process 60 pm sources as 
well. 


V-57 






25 tim HIGH SOURCE DENSITY BINS 





60pm HIGH SOURCE DENSITY BINS 



ORIGINAL PAGE IS 
OF POOR QUALITY 




Figure V.H.1.4 The parts of the 100 pm sky processed according to high source density rules are shown in an Aitoff 
projection in Galactic coordinates. The black regions contain more than 12 sources per sq. deg, the con- 
fusion limit; the grey areas contain more than 6 sources per sq. deg, the threshold used for high source 
density processing. 




* * 


rl! - 

i i .. /- 

H.6.b Catalog Selection Criteria in High Source Density Regions 

In unconfused regions where the reliability of a single hours-confirraed source is high, the criteria 
for including sources in the catalog emphasized completeness. In high source density regions, however, 
the criteria emphasized reliability. Thus, the first step toward weeding out unreliable point sources was to 
stiffen the requirements for accepted detections and valid weeks-confirmed sources in a given band. 

The concept of high, medium and low quality fluxes is crucial to understanding how sources were 
included in the catalog. High quality measurements were those which passed through the entire data pro- 
cessing chain without blemishes such as more than one missing detection (per hours-confirmed sighting), 
even those missing due to failed detectors. Medium quality fluxes could suffer from a variety of problems 
due either to the faintness of the source, the presence of failed detectors, the complexity of the back- 
ground or peculiarities of the data processing. Low quality fluxes were upper limits based on an estimate 
of the local noise. The definition of the relative quality of fluxes and the way that sources were selected 
for inclusion in the catalog on the basis of their quality differ depending on whether a source was found 
in a region of high or low source density. 

In regions of high source density, tests were applied successively to the measurements of a source to 
determine its quality. If all of the criteria described below were satisfied, the source was deemed to have 
a high quality measurement in that band. If only a subset of the criteria were satisfied, then the source 
was considered to have a medium quality measurement in that band. If none, or only a few, of the cri- 
teria were satisfied, then an upper limit was given for the flux of that source. To be included in the cata- 
log a source had to have a high quality measurement in at least one band. 

For high source density regions the requirements for a high quality flux were stiffened by demand- 
ing that at least two hours-confirmed sightings in one band each meet all of the following criteria: 1) 
seconds-confirmed detections on at least two orbits; 2) a minimum correlation coefficient of 0.97, 
corresponding to a local signal-to-noise ratio of about 6; and 3) fewer than four detections within the 
seconds-confirmation window on a given scan. Simple edge detections (3 detections) were permitted. 

A high quality measurement had to be repeatable. The ratio of brightest to faintest high quality 
flux measurements in a band had to differ by less than a factor of 3.0. If this repeatability requirement 
was not satisfied, then the brightest measurement was taken as a medium quality 
measurement. Although this requirement in principle biases the catalog against variable sources in high 
source density regions, in practice, relatively few sources were rejected for this reason. 

H.6.c Weaker Neighbors 

The above rules all pertain to a single source, independent of its environment. The weaker neigh- 
bor algorithm was designed to remove sources adversely affected by brighter, neighboring sources. The 
"neighbor" of a source with a high quality flux in some band must itself have at least two hours-confirmed 
measurements in that same band and lie within a wavelength dependent position window. If the bright- 
est flux of one source was a factor of 1 .2 weaker than the faintest flux of the other object in the same 
band, then the fainter object was marked as a "weaker neighbor" in that band. The flux status of the 
weaker neighbor was demoted to low quality and the upper limit assigned was that of the flux of the 
brighter object. The half-widths of the weaker neighbor window at (12, 25, 60 and 100 pm) were (90”, 
90",180" and 360") and the half-lengths were (270”, 272”, 286”, 300"--a detector length). 


V-62 


H.6,d Confused Neighbors 

A good indicator that a source was part of a more extended structure was to have two or more 
sources located very close together. A source with a high quality measurement in one band was exam- 
ined for neighbors in the same window as described for weaker neighbors. If another source were located 
sufficiently close to the first and if any confusion status flags other than edge detection flags (Section 
V.D.2) were set in either source, then the fluxes of both were marked low quality and the brightest of the 
two fluxes was assigned as an upper limit to each source. If no confusion status flags were set, neither 
source was changed. The definition of a neighbor and the size of the confused neighbor window were as 
described above for the weaker neighbor test. 

H.6,e Very Near Neighbors 

One of the original selection rules was to exclude from the catalog all sources, in all bands, with a 
neighbor closer than 30” in-scan and 90” cross-scan. Because the clean-up processor described in Section 
V.H.2 forced the weeks-confirmation of all sources within such a box, there should have been no such 
sources remaining to delete. Due, however, to slight differences in the way in which the neighbor boxes 
were calculated by the two programs, about twenty sources at the periphery of the neighbor box were 
deleted by the high source density processor. 

H.6.f Moderate Quality Fluxes 

Measurements that failed one or another of the above rules could still carry useful information 
about the strength of a source in a given band. Moderate quality fluxes had to have at least two hours- 
confirmed measurements with flux status, FSTAT=3,4,5 or 7, and correlation coefficient >0.95. Fluxes 
of otherwise moderate quality could be demoted to low quality by either the weaker or confused neighbor 
rules. 

H.6.g Low Quality Fluxes (Upper limits) 

Measurements in a given band that failed to meet either the high or moderate quality criteria 
become upper limits for the source. If multiple detections were available, the brightest one was 
given. As mentioned above, confused or weaker neighbors were given the brightest flux of either of 
the two sources as an upper limit. Sources of low quality were examined for brighter neighbors and were 
assigned the brightest neighboring flux (see above for neighbor rules). Sources with no detections on any 
sighting (i.e., with only noise fills present) and no neighbors were given an upper limit based on 12 times 
the 1 a value reported by the noise estimator. This value was chosen since the correlation coefficient 
requirement of >0.97 corresponds to roughly a local signal-to-noise ratio of 12. 

H.6,h Flux Averaging and Uncertainties 

In all cases, averages and uncertainties were computed using the logarithms of the fluxes, not the 
fluxes themselves. However, the technique used to obtain the average values for both high and moderate 
quality fluxes in high source density regions was different from the one used outside of dense 
regions. Rather than computing inverse-variance weighted averages in regions where the noise estima- 
tor was suspect, simple averages of the logarithms of the valid high (or moderate) quality fluxes were 
computed. The uncertainty associated with the measurement was the greater of: 1) the average of the N 
uncertainties quoted for the N valid hours-confirmed measurements, divided by N l/2 \ or 2) the standard 
deviation of the mean of the N values averaged together to obtain the flux average. 


V-63 


H.7 Catalog Source Selection 


Outside of regions of high source density, two rules were used to select weeks-confirmed WSDB 
sources for inclusion in the catalog: sources had to have either a high quality flux in at least one band or 
two or more moderate quality fluxes in adjacent bands (Section V.H.5). The minimum time separation of 
123,000 sec. required for two hours-confirmed sightings to "weeks-confirm" was found to be long enough 
to reject even slow moving asteroids. 

Within high source density regions, as described in V.H.6, sources had to have at least one high 
quality flux in one band and be relatively isolated from other sources to be included in the catalog. 

H.8 Low-Resolution Spectral Associations 

For each source in the WSDB, a check was made to see whether a low-resolution spectrum was 
available. See Chapter IX for more details. 

H.9 Associations 

Positional associations of IRAS sources are made with objects in other astronomical catalogs. The 
associations in the IRAS catalog were based purely on positional agreement, with no attempt made to 
distinguish between multiple sources associated with a particular IRAS source. Any number of sources 
from a variety of catalogs could be associated with a specific IRAS object as long as the position test was 
met. No attempt was made to "identify" an IRAS source with a source from another catalog, by requir- 
ing the IRAS source to have a "reasonable” energy distribution for the identification. The only attempt to 
preclude spurious associations was to forbid a source with only a 100 pm flux (i.e., a source having a 
significant likelihood of being Galactic "cirrus") from being associated with stars. 

Because the catalogs used for associations had a wide range of positional accuracies, and the IRAS 
positions were intermediate between the very high positional accuracy catalogs and the lower accuracy 
catalogs, the following procedures were adopted to make the associations. For the high-positional-accuracy 
catalogs, a window of half width 8” in-scan and 45" cross-scan was used. If the position of the IRAS 
source lay within this distance of the cataloged source, then an association was made. For those catalogs 
where the positional uncertainty was greater than that of the IRAS source, an association was made if the 
source position and the IRAS position agreed to within a fixed radius or to within the radius (see below) 
of the cataloged source, if available. 

In addition to the other catalogs, associations with IRAS small extended sources have been 
included, so that the user will know if a small extended source is close to the point source. 

Table V.H. 1 identifies the catalogs used and gives the search radius for the association and an indi- 
cation of whether a source size was used for the association radius. The information earned with the 
association is described in detail in Section X.B and includes such information as visible magnitude, size, 
spectral type, and morphological type. In addition, the distance and position angle from the IRAS posi- 
tion to the location of the cataloged source are given. 


V-64 



Table V.H.l Catalogs Used for Associations with IRAS Sources 


Catalog 

No. 

Catalog 

Type* 

Catalog Name 

Search 

Radius 

Source Size 
Used in 
Association 

01 

2 

General Catalogue of Variable Stars 
Kukarkin, et al. 

90” 


02 

2 

Dearborn Observatory Catalogue of 
Faint Red Stars, Lee, et al. 

90” 


03 

3 

Air Force Geophysical Laboratory 
Four-Color Survey, Price & Walker 

90" 


04 

2 

Two Micron Sky Survey 
Neugebauer and Leighton 

90" 


05 

3 

Globule List 
Wesselius 

90” 

X 

06 

1 

Second Reference Catalogue of Bright 
Galaxies, de Vaucouleurs, et al 

90" 


07 

2 

Early Type Stars with Emission Lines 
Wackerling 

90” 


08 

3 

Equatorial Infrared Catalogue 

90" 


09 

1 

Uppsala General Catalogue of Galaxies 
Nilson 

90" 


10 

1 

Morphological Catalog of Galaxies 
Vorontsov-Velyaminov, et al 

180" 



* Catalog types include (1) extragalactic, (2) stellar or (3) other, e.g. dark clouds, HIT regions, etc. 


V-65 




II Pi" I Ilfll'W! 


Catalog 

No. 

Table V.H. 

Catalog 

Type* 

1 Catalogs Used for Associations with IRAS Sources (Cont.) 

Source Size 

Catalog Name Search Used in 

Radius Association 

11 

3 

Strasbourg Planetary Nebulae 

90 " 


12 

1 

Catalogue of Galaxies and Clusters 
of Galaxies, Zwicky, et al. 

90" 


13 

1 

Smithsonian Astrophysical Observatory 
Star Catalog 

45"(x-scan) 

x8"(in-scan) 


14 

3 

ESO/Uppsala Survey of the ESO (B) 
Atlas, Lauberts 

90" 


15 

2 

Bright Star Catalogue - 4th Edition 
Hoffleit 

45"x8" 


16 

2 

New Catalog of Suspected Variable Stars 
Kukarkin, et al. 

90" 


17 

2 

General Catalogue of Cool Carbon Stars 

90" 


18 

2 

Catalog of Nearby Stars 
Gliese 

45"x30" 


19 

2 

General Catalog of S Stars 
Stephenson 

90" 


20 

3 

Parkes HII Region Survey 
Haynes, et al. 

120" 

X 


* Catalog types include (1) extragalactic, (2) stellar or (3) other, e.g. dark clouds, HII regions, etc. 


V-66 


:| Hi 



Table V.H.l Catalogs Used for Associations with IRAS Sources (Cont.) 

Source Size 

Catalog Catalog Catalog Name Search Used in 

No. Type* Radius Association 

21 3 Bonn HII Region Survey 80" 

Altenhoff, et al. 

22 3 Catalog of CO Radial Velocities 80" X 

Toward Galactic HII Regions, Blitz 
et al. 

23 3 Catalogue of Dark Nebulae X 

Lynds 

Comparison Catalog of HII Regions 
Marsalkova 

Catalog of Star Clusters and Associations 
Alter, et al 

Catalogue of Bright Diffuse Galactic 
Nebulae, Cederblad 
Untersuchungen Uber Reflexionsnebel 
am Palomar Sky Survey, Dorschner and 
Gurtler 

A Study of Reflection Nebulae 
van den Bergh 

Catalog of Southern Stars Embedded in 
Nebulosity, van den Bergh and Herbst 

24 2 Two Micron Sky Survey with Improved 45"x8" 

Positions, Kleinmann and Joyce 

* Catalog types include ( 1 ) extragalactic, (2) stellar or (3) other, e.g. dark clouds, HII regions, etc. 


V-67 




Table V.H.l 

Catalogs Used for Associations with IRAS Sources (Cont.) 

Catalog 

No. 

Catalog 

Type* 

Catalog Name 

Search 

Radius 

Source Size 
Used in 
Association 

25 

1 

Catalog of Dwarf Galaxies 
van den Bergh 

90" 


26 

1 

Atlas of Peculiar Galaxies 
Arp 

120" 


27 

1 

Galaxies with an Ultraviolet 
Continuum, Markarian, et al. 

90" 


28 

1 

Catalog of Extragalactic Radio 
Sources Having Flux Densities 
Greater than 1 Jy at 5 GHz, 
Kuhr, et al. 

60" 


29 

1 

Catalogue of Quasars and Active 
Nuclei, Veron-Cetty and Veron 

90" 


30 

1 

Lists of Galaxies 
Zwicky 

90" 


31 

1 

Atlas and Catalog of Interacting 
Galaxies, Vorontsov-Velyaminov 

120" 


32 

3 

IRAS Small Scale 
Structure Catalog 

min(120", 

diameter/2) 

diameter/2 

39 f 

3 

Ohio State University 
Radio Catalog 

120" 


40 

2 

University of Michigan 
Spectral Atlas 

60" 
45" x 8" 

(Vol. 1) 
(Vols. 2,3) 

41 

3 

IRAS Serendipitous 
Survey Catalog 

60" 



"“Catalog types include (1) extragalactic, (2) stellar or (3) other, e.g. dark clouds, HII re- 
gions, etc. 

f Catalog numbers 33-38 reserved for internal use. 



Authors: 


T. Chester, B. T. Soifer, C. Beichman, J. Fowler, T.N. Gautier, G. Helou, H. McCallon, M. 
Rowan-Robinson, T. Conrow, M. Hauser, D. Walker, R. Walker. 

References: 

Acker, A., Marcout, J., and Ochsenbein, F., 1981, Astron. Astrophys. [Supp.l 43, 265. 

Altenhoff, W.J., Downes, D., Pauls, T., and Schraml, J., 1979, Astron. Astrophys. [Suppl.l 35, 23. 

Alter, G., Balasz, B., and Ruprecht, J., 1970, Catalogue of Star Clusters and Associations. Budapest: 
Akademiai Kiado. 

Arp, H., 1966, Ap.J. [Suppl.l 14, 1. 

Blitz, L., Fich, M. and Stark, A. A., 1982, Ap.J. [Supp.l 49, 183. 

Cederblad, S., 1946, Lund Annals, Ser. 2, No. 119. 

de Vaucouleurs, G., de Vaucouleurs, A., and Corwin, Jr., H.G., 1976, Second Reference Catalogue of 
Bright Galaxies. Austin: University of Texas Press. 

Dorschner, V.J., and Gurtler, J., 1964, Ast. Nachr. 287, 257. 

Fisher, J.R., and Tully, R.B., 1975, Astron. Astrophys. 44, 151. 

Foltz, C.B., Peterson, B.M., and Boronson, T.A., 1980, A. J. 85, 1328. 

Fowler, J.W., and Rolfe, E.G., 1982, Journal of the Astronautical Sciences, Vol. XXX, No. 4, 385 

Gliese, W., 1969, Veroffentl. Astron. Rechen-Instituts Heidelberg, No. 22. 

Haynes, R.F., Caswell, J.L., and Simons, L.W.J., 1979, Aust. J. Phys. Astrophys. [Supp.l, No. 48. 

Hoffleit, D„ and Jascheck, C., 1982, The Bright Star Catalogue, Fourth Revised Edition. New Haven: 
Yale University Observatory. 

Kleinmann, S.G., and Joyce, R.R., 1984, private communication. 

Kojoian, G., Chute, P.A., and Aumann, C.E., 1984, A. J. 89, 332. 

Kojoian, G., Elliott, R., and Bicay, M.D., 1981, A. J. 86, 816. 

Kojoian, G„ Elliott, R„ and Bicay, M.D., 1982, A. J. 87, 1364. 

Kojoian, G., Elliott, R., and Tovmassian, H.M., 1978, A. J. 83, 1545. 

Kojoian, G., Elliott, R., and Tovmassian, H.M., 1981, A. J. 86, 81 1. 

Kuhr, H„ Witzel, A., Pauliny-Toth, I.I.K., and Nauber, U., 1981, Astron. Astrophys. [Supp.l 45, 367. 


V-69 


Kukarkin, B.V., Kholopov, P.N., Artirukhina, N.M., Fedorich, V.P., Frolov, M.S., Goranskij, V.P., 
Gorynya, N.A., Karitskaya, E.A., Kireeva, N.N., Kukarkina, N.P., Kurochkin, N.E., Medvedevna, G.I., 
Perova, N.B., Ponomareva, J.A., Samus’, N.N., and Shugarov, S. Yu., 1981, New Catalog of Suspected 
Variable Stars (on magnetic tape), Moscow. 

Kukarkin, B.V., Kholopov, P.N., Efremov, Yu. N., Kukarkina, N.P., Kurochkin, N.E., Medvedeva, G.I., 
Perova, N.B., Fedorovich, V.P., and Frolov, M.S., 1970, General Catalog of Variable Stars, Vol II. 
Sternberg Inst., Moscow State University. 

Kukarkin, B.V., Kholopov, P.N., Pskovsky, Yu.P., Efremov, Yu.N., Kukarkina, N.P., Kurochkin, N.E., 
Medvedeva, G.I., Perova, N.B., Fedorovich, V.P., and Frolov, M.S., 1971, General Catalogue of Variable 
Stars, Vol. III. Sternberg Inst., Moscow State University. 

Lauberts, A., 1982, The ESO /Uppsala Survey of the ESO(B) Atlas. Mucnich: European Southern Obser- 
vatory. 

Lee, O.J., Baldwin, R.J., and Hamlin, D.W., 1943, Annals of the Dearborn Observatory V, Part l A. 

Lee, O.J., and Bartlett, T.J., 1944, Annals of the Dearborn Observatory V, Part IB. 

Lee, O.J., Gore, G.D., and Bartlett, T.J., 1947, Annals of the Dearborn Observatory V, Part 1C. 

Lynds, B.T., 1962, Ap. J. [Suppl.l 7, 1. 

Markarian, B.E., 1967, Astrofizika 3, 55. 

Markarian, B.E., 1969a, Astrofizika 5, 443. 

Markarian, B.E., 1969b, Astrofizika 5, 581. 

Markarian, B.E., and Lipovetskii, V.A., 1971, Astrofizika 7, 511. 

Markarian, B.E., and Lipovetskii, V.A., 1972, Astrofizika 8, 155. 

Markarian, B.E., and Lipovetskii, V.A., 1973, Astrofizika 9, 487. 

Markarian, B.E., and Lipovetskii, V.A., 1974, Astrofizika 10, 307. 

Markarian, B.E., and Lipovetskii, V.A., 1976a, Astrofizika 12, 389. 

Markarian, B.E., and Lipovetskii, V.A., 1976b, Astrofizika 12, 657. 

Markarian, B.E., Lipovetskii, V.A., and Stepanyan, D.A., 1977a, Astrofizika 13, 225. 

Markarian, B.E., Lipovetskii, V.A., and Stepanyan, D.A., 1977b, Astrofizika 13, 397. 

Markarian, B.E., Lipovetskii, V.A., and Stepanyan, D.A., 1979a, Astofizika 15, 201. 

Markarian, B.E., Lipovetskii, V.A., and Stepanyan, D.A., 1979b, Astrofizika 15, 363. 

Markarian, B.E., Lipovetskii, V.A., and Stepanyan, D.A., 1979c, Astrofizika 15, 549. 


V-70 



Markarian, B.E., Lipovetskii, V.A. and Stepanyan, D.A., 1981, Astrofizika 17 , 619. 

Marsalkova, P. 1974, Astrophys. Sp. Sci. 27, 3. 

McCallon, H.L. and Kopan, E.L., IRAS Telescope Pointing Reconstruction, JPL Publication No. 85-1, to 
be published. 

Neugebauer, G., and Leighton, R.B., 1969, Two-Micron Sky Survey, NASA SP-3047. Washington, D.C.: 
National Aeronautics and Space Administration. 

Nilson, P., 1973, Uppsala General Catalogue of Galaxies. Act Universitatis Upsaliensis, Nova Regiae 
Societatis Upsaliensis, Series V:A, Vol. 1 . 

Peterson, D.S., 1973, A. J. 78, 811. 

Pnce, S.D., and Murdock, T.L., 1983, The Revised Air Force Geophysics Laboratory Infrared Sky Survey, 
AFGL-TR-83-0161. Hanscom Air Force Base, Massachusetts: Air Force Geophysics Laboratory. 

Rolfe, E.G., Otake, H., and Fowler, J.W., 1984 (in preparation for submission to the Journal of the 
Astronautical Sciences). 

Smithsonian Astrophysical Observatory Star Catalog (4 vols.) 1966. Washington, D.C.: Smithsonian 
Institution. 

Stephenson, C.B., 1973, Publications of the Warner and Swasey Observatory, Case Western Reserve 
University, 1, No. 4. 

Stephenson, C.B., 1976, Publications of the Warner and Swasey Observatory, Case Western Reserve 
University, 2, No. 2. 

Sweeney, L.H., Heinsheimer, T.F., Yates, F.F., Maran, S.P., Lesh, J.R., and Nagy, T.A., 1978, Interim 
Equatorial Infrared Catalogue, TR-0078(3409-20)-l. Los Angeles: The Aerospace Corporation. 

van den Bergh, S., 1966a, A. J. 71, 922. 

van den Bergh, S., 1966b, A. J. 71, 990. 

van den Bergh, S., and Herbst, W., 1975, A. J. 80, 208. 

Veron-Cetty, M.P., and Veron, P., 1984, A Catalogue of Quasars and Active Nuclei. Munich: European 
Southern Observatory. 

Vorontsov- Velyaminov, B.A., 1959, Atlas and Catalog of Interacting Galaxies. Sternberg Inst., Moscow 
State University. 

Vorontsov- Velyaminov, B.A., and Arhipova, V.P., 1963, Morphological Catalog of Galaxies, Part III. 
Moscow State University. 

Vorontsov- Velyaminov, B.A., and Arhipova, V.P., 1964, Morphological Catalog of Galaxies, Part II. 
Moscow State University. 


V-71 


Vorontsov-Velyaminov, B.A., and Arhipova, V.P., 1968, Morphological Catalog of Galaxies, Pan IV. 
Moscow State University. 

Vorontsov-Velyaminov, B.A., and Arhipova, V.P., 1974, Morphological Catalog of Galaxies, Pan V. 
Moscow State University. 

Vorontsov-Velyaminov, B.A., and Krasnogorskaja, A.A., 1962, Morphological Catalog of Galaxies, Part 
I. Moscow State University. 

Wackerling, L.R., 1970, Mem. R.A.S. 73, 153. 

Wesselius, P.R., 1979, unpublished. 

Zwicky, F., and Herzog, E„ 1963, Catalogue of Galaxies and of Clusters of Galaxies, Vol II. Pasadena: 
California Institute of Technology. 

Zwicky, F„ and Herzog, E., 1966, Catalogue of Galaxies and of Clusters of Galaxies, Vol III. Pasadena: 
California Institute of Technology. 

Zwicky, F„ and Herzog, E., 1968, Catalogue of Galaxies and of Clusters of Galaxies, Vol IV. Pasadena: 
California Institute of Technology. 

Zwicky, F., Herzog, E„ and Wild, P., 1961, Catalogue of Galaxies and of Clusters of Galaxies, Vol 
I. Pasadena: California Institute of Technology. 

Zwicky, F„ Karpowicz, M„ and Kowal, C.T., 1965, Catalogue of Galaxies and of Clusters of Galaxies, 
Vol. V. Pasadena: California Institute of Technology. 

Zwicky, F. and Kowal, C.T., 1968, Catalogue of Galaxies and of Clusters of Galaxies, Vol. VI. Pasadena: 
California Institute of Technology. 

Zwicky, F., Sargent, W.L.W., and Kowal, C.T., 1975, A. J. 80, 545. 

Zwicky, F„ and Zwicky, M.A., 1971, Catalogue of Selected Compact Galaxies and of Post-eruptive 
Galaxies. Zurich: Offsetdruk L. Speich. 


V-72 


VI. FLUX RECONSTRUCTION AND CALIBRATION 


Three quite separate functions must be performed in order to reconstruct the flux of a source from 
the telescope data stream. First, the effects of the telescope transfer function must be removed from the 
data. Second, the results must be transferred to relative photometric units. In the case of IRAS, this 
amounted to establishing the relationship between source amplitudes observed during the survey scans 
and outputs of flashes of the internal reference source. Third, the relative photometry must be put on an 
absolute scale, i.e., the flashes of the internal reference source must be calibrated in an absolute sense. 

In the following discussion, the processing by which the signal received at the ground station was 
converted to the effective detector current is described, followed by the description of the process by 
which the relative and, finally, the absolute photometry was achieved. 

A. Processing - Removal of Telescope Transfer Function 

The electronic chain by which the current caused by photons incident onto the detector is 
transferred to the signal received at the ground station is described in Section II.C.5. Briefly, the detector 
acted as a current source as a result of the incident photons. This current was input to a trans- 
impedance amplifier which transformed the signal into a voltage equaling the total current times the feed- 
back resistor plus a small offset voltage due primarily to the imbalance in the cold JFETs at the inputs to 
the operational amplifiers. The output voltage from the trans-impedance amplifier was amplified and 
shaped by analog electronics, fast rise time nuclear particle spikes were removed and the analog voltage 
was converted into a digital signal for transmission in a compressed format to the Earth. In the process- 
ing, the route was reversed. As described below, the received signal, in data numbers, was retraced to the 
current at the detector. 

A. 1 Digital Electronics 

Pre-launch ground tests determined that there were no significant errors or non-linearities intro- 
duced by the digital electronics or by the analog to digital conversion. Any errors introduced by the 
compression of digital data on board the telescope and reconstruction at the processing center were 
assumed to be random and to be in the least significant bit and were ignored. The raw input was there- 
fore assumed to correspond exactly to the output of the analog to digital converter. Since a linear analog 
to digital converter was assumed, the voltage into the analog to digital converter, i.e., the voltage at the 
output of the analog chain was given by: 

Vanig - Co + Cl X DN (VI. A. 1) 

where 


V an [ g — voltage at output of the analog amplifier - voltage at the analog digital converter 
Co — inherent voltage offset of the analog digital converter 
DN - observed signal in data numbers. 

C i - data number voltage conversion factor (volts/DN) 


VI- 1 



Preflight measurements showed that any adverse effect on the data due to the multiplexer should be negli- 
gible and it was therefore assumed to be linear. 

A.2 Analog Electronics Amplifiers 

The analog electronics contained the commandable gains and offsets. The input voltage to this part 
of the electronics can be recovered by 



where, 

V, ia — voltage at trans-impedance amplifier output 
Gj — commandable gain, i — 1,2,3. 

Sj — commandable offset, j — 1 , 2 ,..., 8 . 

The effects on the flux calibration due to the pole zero amplifier and the pulse circumvention circuit 
(deglitcher), which also lie in the analog electronics board between the output of the preamplifier elec- 
tronics and the commandable offset electronics, were ignored. The analog electronics also contained a 
filter to remove high frequencies in the signal. Adverse effects due to the filter were considered to be 
small at the frequencies of interest and were therefore ignored. 

A.3 Trans-impedance Amplifier 

In order to recover the extended component of the incident flux, it was necessary to determine, and 
remove, any electronic baseline offset. The voltage at the output of the trans-impedance amplifier was 
taken to be of the form: 


V lia [observed] - Ij x Rf + V tia [off] (VI.A.3) 

where Id is the current through the detector and Rf is the feedback resistor. 

The offset voltage, V tia [ off] represents a sum of terms due to (a) offsets in the trans-impedance 
amplifier, including the cold JFETs, and (b) offsets in the analog electronic boards. Since the effects are 
indistinguishable at the output, both have been lumped into one equivalent offset voltage at the output of 
the trans-impedance amplifier. Although these offsets were in principle temperature dependent, it was 
concluded from pre-flight tests of the relevant temperature coefficients, and the relative stability of the 
electronics boxes in flight, that this effect was small enough that it was not necessary to model the tem- 
perature dependent aspects of V tia [off], A single fixed value for the electronic baseline was therefore used 
for each SOP. The details of the determination of V tia [off] are given in Section VI.B.3. 

After launch, the electronic baselines at 60 and 100 pm were found to be affected by the bias boosts 
induced during crossings of the South Atlantic Anomaly (SAA). The effect is thought to be due to the 
differential heating of the JFETs because of the large current passing through the JFET connected to the 



(VLA.4) 


detector. This effect was modeled by adding a time dependent exponential function to V tia [off] 

Corr - (a + b x T) x exp(— A/x) 

where: 


T — duration of the most recent bias boost 
a, b - constants 

A - time interval between time of interest and end of most recent bias boost 
x - characteristic decay time of the baseline shift 

This correction was added to the electronic baseline value before removing the electronic baseline from 
the data. For a typical passage through the SAA, the magnitude of this baseline shift was about 20% and 
80% of the signal from the TFPR at 60 pm and 100 pm. The time constant x was typically 52 
minutes at 60 pm and 42 minutes at 100 pm. 

A. 4 Removal of Coherent Detector Noise 

Detector 5 (100 pm) exhibited a stable, coherent 0.25 Hz waveform. By adding successive 4-second 
intervals of data coherently, it was possible to determine the shape of the waveform to better than 1% of 
the remaining rms noise. This wave form was then subtracted from the detector 5 data stream. This 
resulted in an order of magnitude improvement in the rms noise and greatly reduced the number of 
spurious point source detections for detector 5. After removal of the wave form, the character of the 
resulting detector 5 data stream was typical of 100 pm detectors. 

Detector 19 (25 pm) exhibited a spike, of randomly varying amplitude, every second at sample #9 
of the 16 Hz sampled data stream. The typical amplitude was approximately 6-8 times the cleaned rms 
noise for detector 19. The amplitude was measured by subtracting the average of the neighboring sam- 
ples, which showed no evidence of electronic cross-talk, from the value of the spike sample. If the spike 
amplitude was less than a maximum allowable amplitude (approximately 10 times the rms noise of the 
cleaned detector 19 data), then the spiked sample was replaced by the average of its neighboring samples. 
If the spike amplitude was greater than the maximum allowable amplitude, then the maximum allowable 
amplitude was subtracted from the spike sample. The latter case can occur in the 6% of detector 19 
detections for which the spike appears on top of a point source. If the signal to noise ratio of such a 
point source after the spike was removed is less than 20, the result will be a "topped" source. This pro- 
cess can lower the correlation coefficient below threshold for some of the weaker sources. Sources 
brighter than signal to noise ratio — 20 will have progressively higher correlation coefficients. After 
removing the spikes, the overall rms noise for detector 19 was reduced by a factor ~ 3. 

Detector 43 exhibited a 1 Hz waveform that consisted of a positive spike with signal to noise ratio 
~ 2 at sample #9 of the 16 Hz sampled output, followed by a small negative spike, with signal to noise 
ratio — 0.5, at sample #10. The amplitude varied randomly, but remained small, with a signal to noise 
ra t*° ^ 3. In this case a constant waveform of the above shape was subtracted, resulting in a reduction 
of the detector 43 rms noise by a factor of ~ 1.3. All subtractions were effectively done at the output of 
the trans-impedance amplifier. 


VI-3 



Fieure VI. A. 1 The fit used to represent the feedback resistor is shown in the four wavelength bands. 

The flux densities corresponding to the appropriate voltages across the resistor are 
shown. The same fit, that shown in Fig. II.C.12, was used for all the detectors. 


A. 5 Feedback Resistor 

The feedback resistors had a nominal resistance of 2 x 10 10 £1, but in fact the resistance of the feed- 
back resistor was dependent on the voltage across the resistor. The fit to the feedback resistor used in the 
data reduction is a three-piece curve shown in Fig. VI.A.l.a-d. It is the same fit as shown in Fig. II.C.12. 
This was the only nonlinearity assumed in the processing. 


A.6 Summary 

The processing to remove the instrumental transfer functions can be summarized by the following 
equation for the current in the detector: 


h- 


Co + Cix DN _ _ 

G, 


/ Rf 


(VI.A.5) 


VI-4 



B. Determination of Relative Flux 

The processing outlined above resulted in a current which was assumed to be proportional to the 
flux incident on the detector except for small effects due to the dwell time on a source. The conversion 
of the current to relative flux was effected by a comparison of source amplitudes with amplitudes result- 
ing from flashes from the internal reference source (see Section II.C.3) which provided a secondary rela- 
tive calibration source. The internal reference source was in turn absolutely calibrated with respect to 
celestial sources using NGC 6543 as a secondary transfer standard; see Section VI.C below. The process- 
ing to obtain the relative photometry of the IRAS sources is discussed in this section. 

B. 1 Overall Procedure to Determine Relative Photometry 

Each survey scan was initiated and terminated with a pair of flashes of the internal reference source 
to monitor the responsivity of the system. Because it was assumed that responsivity changes were slow, 
responsivity values between internal reference source flashes were determined by linear interpolation 
between these flashes. In principle, therefore, the responsivity was determined separately for each scan. 
If, as discussed below, both of the flashes at one end of a scan were unusable, the extrapolation was con- 
tinued to the closest accepted flashes on adjoining scans. There was, however, never an interpolation 
across a bias boost (see Sections II.C.5, III.B.6, III.C.4 and IV.A.7) since the bias boost largely erased the 
detector’s memory. In this case, the detector’s responsivity was assumed to equal that calculated from 
the most recent flash accepted. 

The responsivity relative to the flashes of the internal reference source, as determined from the 
flashes at the ends of the scans, was used to scale both point sources and extended emission. As seen in 
Fig. IV.A.2, the actual variations in the responsivity as measured by the flashes of the internal reference 
source are typically < 10% in the 12, 25 and 60 pm bands and < 15% in the 100 pm band. 

As discussed in Section IV.A.4, the responsivity is a function of the spatial frequencies present in 
observing a source. Multiplicative factors appropriate for either the point source responsivity or the 
extended emission ("DC") were applied in the source detection processing (Sections V.C, V.G). No 
attempt has been made to adjust the calibration of the small extended sources to account for the 
appropriate spatial frequencies encountered (Section V.E). Additionally, the observations were reduced 
assuming a responsivity independent of the total flux falling on the detector; as indicated by the observa- 
tions of Fig. IV.A.4, this assumption was clearly wrong in the 60 and 100 pm bands. 

The design of the entire detector and electronics system used in IRAS was DC coupled, thus allow- 
ing, in principle, a measurement of the absolute brightness of the sky from a measurement of the total 
photo-current. There was, however, no absolute radiometric reference included in the experiment in 
order to evaluate V tia [off] directly, an essential quantity in determining the true sky brightness. The 
reconstruction of the offset V tia \ off] for each of the detector channels was the most difficult part of the 
effort to produce properly calibrated absolute sky brightness data for all parts of the sky. For example, in 
order to determine the total incident flux on each detector to an accuracy of one percent, it would be 
necessary to determine the absolute value of the DC voltage level at the input of each amplifier to about 
1 microvolt. Despite this handicap, the extraordinary DC stability of the system, when coupled with the 
calibration procedures described below, provided an accurate and sensitive measurement of the sky 
brightness in all four wavelength bands. 


Vl-5 



B.2 Photometry of Point Sources and Small Extended Sources 

The point source responsivity obtained from the measurements of the internal reference source as 
described above were applied directly to the output of the point source recognizer. No attempt was 
made in the reconstruction to correct for the cross-scan response profile other than in a statistical sense. 
However, as part of the calibration procedure using the flashes of the internal reference source and the 
celestial source NGC 6543, the shape of the cross-scan response of each individual detector (Section 
IV.A.3) was combined with the knowledge of the cross-scan position of the standard calibration track to 
produce a single correction factor for the standard calibration response. This correction factor would 
result in zero average error for point sources passing over uniformly distributed cross-scan positions. In 
the 25 pm band the difference between the measured amplitude and the average amplitude never 
exceeded 10%, while in the other bands the difference did not exceed 5%. The variations in the ampli- 
tudes caused by the cross scan response implies that a minimum uncertainty of 5% must be assigned to 
all fluxes determined in the survey. 

The small extended sources were treated like the point sources except that a 10% and 8% increase 
in responsivity was assumed at 12 and 25 pm to approximate sources typically 4 times the nominal point 
source template in spatial extent. Except for this small correction, no attempt was made to take the 
variation of the responsivity with size of the source into account; see Section IV. A. 4. Although this pro- 
cedure incurred large uncertainties, the uncertainties introduced because of a lack of knowledge of the 
true spatial distribution of the sources dominated the basic photometric uncertainties. 

B.3 Photometry of Extended Emission 

Since there were neither on-board calibration techniques to establish the electronic offset voltage 
directly for each detector under zero photon flux conditions nor celestial sources of known total sky 
brightness, it was necessary to follow a special procedure to determine V tia [off\. The procedures can be 
summarized as follows: 

a) An area of the sky with a smoothly varying sky brightness, free from point sources and near 
the north ecliptic pole and NGC6543, was selected to serve as a secondary calibration source 
for the diffuse emission. The area was called the total flux photometric reference or 
TFPR. Since the viewing geometry of the TFPR through the zodiacal dust cloud varied 
throughout the duration of the mission, the signal from this area of the sky varied with time. 

b) On several occasions during the mission direct measurements of V, ia [off] were obtained while 
pointing at the TFPR. These yielded direct comparisons of the flux from the TFPR, 
FtTFPR], with flashes from the internal reference source (see Section B.3.a below). Hence 
values for the total sky brightness, B[TFPR], were obtained using measured values for the 
solid angles in each band. More accurate and more frequent measures of the annual variation 
of BITFPR] were obtained by differencing the fluxes from the north and south ecliptic poles 
obtained on individual scans which passed over both poles (see Section B.3.b below). These 
data are shown in Fig. VI.B.l along with the best fit models for the annual variation. Table 
VLB. 1 gives the parameters for the model in each band and summarizes the errors of meas- 
urement. 


VI-6 


1 II 




Figure VLB. 1.1 The variation of total sky brightness at 12 and 25 pm at the north ecliptic pole is 
shown as a function of day number of the year (1983, January 1 (UT) is day 1). See 
Section VI.B.3.b for explanation of symbols. 

c) In order to update the values of V, /a [off] as needed for regular processing of the survey data, 
observations of the TFPR were made twice daily and the signals for each detector derived 
from the models of B [TFPR] versus day number were subtracted. 

B.3.a Determination of F[TFPR] 

The determination of F[TFPR] was based on the possibility of changing the detector responsivity, 
and hence the total photocurrent, while leaving the offset voltage V tia [off! unchanged. In the case of the 
12 pm detectors, the responsivity change was easily achieved since two bias voltage levels were available 
within the operating range of the detectors; see Section II.C.5. In the other wavelength bands, where the 
second bias levels resulted in channel saturation, the required change in responsivity was obtained as a 
result of passage through a dense portion of the SAA uncorrected by bias boosting (see Section IV.A.7), 


VI-7 




[TFPR] (MJysr 



Figure VI B 1 2 The variation of total sky brightness at 60 and 100 pm at the north ecliptic pole is 
shown as a function of day number of the year (1983, January 1 (UT) is day 1). See 
Section VI.B.3.b for explanation of symbols. 


and the change in responsivity was achieved by applying bias boost while the TFPR was still observable. 
The special sequences to do this were started and ended by 13/16 second and 15 second multiple flashes 
of the internal reference source. 

The reduction of the special calibration observations is illustrated with the equations representing 
the voltage at the trans-impedance amplifier as: 

V lia [TFPR, 1] -Rl 1] x FlTFPR] + V, ia [off ] 

and (VI.B.l) 

V tia [TFPR , 21 - R[2] x FlTFPR] + V tia l off] 

VI-8 


"1 Tii 



where the symbols [1] and [2] indicate high responsivity and low responsivity conditions of the detector. 
F[TFPR ] and V t ialoff] are assumed constant during the observation. 

/?[1] and R [2], the responsivities in the two states, can be determined from the differential 
responses to the internal reference source: 


UYoallRStf) - jR[l] x FURS] (VI.B.2) 

and 


A{V tia [IRS,2]) -R[ 2] x FURS] 


where ^ r [//?5'l is the source strength of the internal reference source. The effects of the backgrounds and 
offsets cancel out since the measurements are differential. In practice, the flashes used for these special 
measurements were 15 seconds long; they therefore had to be calibrated relative to the 13/16 seconds 
flashes normally used. 

The relative source strength of the TFPR can now be found from: 


FlTFPR] - 


V lia [TFPR, 1] - V,JTFPR,2] 
A(V lia [IRS, 1]) - A(K„J/R.S,2]) 


x FURS] 


(VLB.3) 


The quality of these special calibration observations varied with the wavelength band; see Fig. 
VI.B.l. At 12 pm the observations yielded a smooth sine wave which agreed in phase and amplitude 
with measurements of the ecliptic polar differences seen in pole-to-pole scans; see below. At 25pm it 
was necessary to use only one of the two detector modules (module 2A) since bias boosting altered the 
electronic offset appreciably compared to the relatively small change in responsivity produced by a 2 to 3 
minute bias boost. Module 2A had been modified to overcome pre-flight wiring failures inside the dewar 
and, fortunately, these modifications greatly reduced the effects of bias boosting on V„ a [off]. 

Some independent information concerning the zero point for each detector was available during the 
first week of the mission when the cryogenically cooled cover was still in place. This afforded essentially 
zero background conditions for the 12 and 25 pm detectors and allowed their electronic offsets to be 
measured directly. These measurements agreed with the results of the special calibration measurements 
described above within 6% and 10% at 12 and 25 pm respectively. Unfortunately the internal dewar 
background was not low enough with the cover in place at 60 and 100 pm to give useful results. 

At 60 and 100 pm the change in responsivity of the detectors as a result of the SAA dose and a 
short bias boost application was substantial and it was possible to obtain useful data on a number of 
occasions by delaying the bias boost until the telescope was pointed at the TFPR. All detectors responded 
well to this procedure and repeatable results were obtained even though the data are not as well behaved 
as those obtained in the 12 pm band (see Fig. VI.B.l); in part this is due to the fact that the sky bright- 
ness at 60 and 1 00 pm is considerably lower than at shorter wavelengths. 

Studies of the variation of V„ a [off] throughout the mission in each of the bands show similar 
behavior; the relatively smooth long term drift appears to be well correlated with measured changes in the 
temperatures of the warm electronics boxes. 


VI-9 



Table VI.B.l TFPR Model Parameters' 


Effective wavelength (pm) 12 

Parameter: 2 

Bo (MJy/sr) 3 13.5 

statistical uncertainty 5 0.1 

total uncertainty 6 1 6 

Bi (MJy/sr) 3 1.40 

statistical uncertainty 5 0.03 

total uncertainty 6 .17 

<p (day) 4 -23.8 

statistical uncertainty 5 1.2 

total uncertainty 6 3 


25 

60 

100 

27.6 

7.7 

8.3 

0.8 

0.2 

0.2 

3.6 

1.0 

1.6 

2.3 

0.58 

0 

0.1 

0.07 

0.7 7 

0.3 

0.10 


-22.3 

-26 

- 

1.3 

9 

- 

3 

10 

- 


' The parameters have been converted to sky brightness (MJy sr -1 ) in order to illustrate the relative mag- 
nitudes of the parameters. The parameters were originally derived relative to the flashes of the internal 
reference source. 

2 At a time t in days the model assumes BlTFPR] to be: 

B[TFPR] - B 0 + x sin[(27t/365.25) x (/ - q>)] 

3 The usual convention of using a flat spectral distribution for the sources was followed in deriving the 
flux densities. 

4 1983 January 1 (UT) is day 1. 

5 The statistical uncertainty corresponds to a 1 standard deviation in the fit to the observations. 

6 The total uncertainty incorporates the estimated systematic uncertainties: 5, 6, 7 and 10% in the abso- 
lute calibration at 12, 25, 60, and 100 pm respectively; 5% for the frequency response and 10% for the 
solid angle. Note that these uncertainties are estimated to apply at the TFPR; they are clearly exceeded 
in other places in the sky. 

7 3 a upper limit. 


It must be kept in mind that no tests were available to validate the design goal that the instrumental 
background would be negligible at 60 and 100 pm. The tests carried out to measure the out of field radi- 
ation, however, indicate that stray radiation from the Sun and Earth were negligible (see Section IV.C.3). 
It is assumed that no significant radiation from within the instrument was able to reach the detectors. 




B.3,b Determination of the TFPR Annual Variation 


As discussed by Hauser et al. (1984) the component of the sky brightness produced by the zodiacal 
dust emission varies as the Earth makes its way through the interplanetary dust cloud. It was possible to 
measure this variation without knowledge of the electronic offset voltage by assuming that v,j off] 
changes negligibly on the time scale of half an IRAS orbit. The observed difference between the bright- 
ness of the north and south ecliptic poles on a single scan is therefore a measure of the true difference in 
the sky brightness. This difference was found to vary seasonally and to be well represented in all 
wavelength bands except at 100 pm by a sinusoidal modulation of half amplitude B, [TFPR], where the 
phase angle does not vary with wavelength. 

The pole-to-pole differences determined the time dependence of the TFPR brightness much more 
accurately than the special calibration measurements. The polar difference data were used to determine 
the phase angle and modulation amplitude of the TFPR variation, while the special calibrations provided 
the average DC brightness B 0 [TFPR], The best fit sinusoidal models in each band are presented in Fig. 
VI.B.l, along with the half values of the ecliptic polar differences. B 0 was determined by a least squares 
fit of the special calibration data to the independently determined sinusoidal model for the annual varia- 
tion. 

In Fig. VLB. 1 the solid dots represent half of the observed difference between the north and south 
ecliptic poles measured on single pole-to-pole scans. The total sky brightness measured with the special 
calibration observations is shown by solid squares for the north calibration polar region (TFPR) and by 
solid triangles for the south polar region. The polar difference data were positioned vertically by fitting 
them to the north polar special calibration observations. The model of sky brightness variation derived 
from these data and given in Table VLB. 1 is shown as the solid curves. The special calibration observa- 
tion data for the south polar region was inverted in phase before plotting in the figure. The dotted curve 
at 100 pm shows the best fit sine function. The open diamonds show the TFPR brightness after the 
cover was ejected, assuming the telescope background with the cover on was zero at 12 and 25 pm. 

At 100 pm the annual variation was so small that the polar differences were dominated by small 
variations resulting from small differences in scan tracks passing over "cirrus" clouds near the poles (see 
Low et al. 1984). As a result, the value of B, [TFPR] at 100 pm was taken to be zero. An upper limit to 
the value of Bj [TFPR] at 100 pm was found by fitting a sinusoid of the same phase found at shorter 
wavelengths to the special calibration data. As can be seen in Fig. VI.B.l, the result is a poor fit to the 
data and the 3 a upper limit is listed in Table VI.B.l. 

Over half of the observations of the survey went through two passes of processing. The first pass 
of the data reduction was used to determine parameters of the model for the time dependent behavior of 
the TFPR. Once the values for B! and B 0 listed in Table VI.B.l were derived as described above, they 
were used during the second pass to recompute the twice daily determinations of V ua [TFPR] which were 
then used to correct all of the survey observations for small variations in the electronic offsets. No 
attempt was made to include higher order terms in the model for the annual variation. 

Systematic uncertainties in the value of B(TFPR) remain due to the limitations of the total flux cali- 
bration procedure in fixing the average value of the TFPR brightness, the assumption of sinusoidal time 
variation of the TFPR, the methods used to determine the amplitude of the variation, the ratio of point 


VI-11 




source to DC responsivity (Section IV.A.4) and the effective detector solid angles (Section IV.A.3). The 
statistical uncertainties listed in Table VI. B. 1 are derived from the fits to the observations. The total 
uncertainties listed in Table VI. B. 1 include the uncertainties in the absolute calibration (Section 
VI.C.2.c), an estimated 5% uncertainty in the frequency response for sources as bright as the TFPR, and 
a 10% uncertainty in the determination of the solid angles. The uncertainty in determining the average 
value of the TFPR brightness using the responsivity-switching procedures is difficult to estimate, but the 
comparison with the cover on/cover off test suggests this may be on the order of 10% at 12 and 25 pm. 
No independent corroboration is available at 60 and 100 pm. Refer to Section V.G for a discussion of 
additional effects in directions other than the ecliptic poles, such as the fixed pattern noise and correc- 
tions for the variations caused by motion through the zodiacal cloud. 

B.4 Problems 


B.4.a Rejected Internal Reference Source Flashes 

The ability to extract an accurate measure of the amplitude of the internal reference source flash 
was complicated because of the structure of the sky over which the internal reference source was flashed. 
In particular, a good measure of the baseline accompanying the flash was required. The difficulty in 
extracting this information increased as the complexity of the sky emission increased. Several checks 
were implemented in the software to assure that an unreliable internal reference source measurement was 
rejected. 

Errors could arise if extrapolations had to be made over a rejected flash to the next available one 
and also if an internal reference source flash was accepted but should have been rejected. This usually 
occurred when the majority of internal reference source measurements in a band were rejected due to 
complex sky background, but one or two slipped past the threshold checks in the software. Reference 
flashes were rejected in fewer than 7% of survey scans and this usually occurred at 60 or 100 pm. The 
photometric error associated with the extrapolation to the next available reference outside the current 
scan is estimated to be less than 5% in the 12, 25, and 60 pm bands. In the 100 pm band possible pho- 
ton induced responsivity enhancement (Section IV.A.8 and VI.B.4.c) may result in photometric errors up 
to 15%. The error associated with the use of reference flashes which should have been rejected is 
estimated to be less than 10% at 12 and 25 pm and less than 20% at 60 and 100 pm. 

B.4.b Radiation 

The flux reconstruction took into account the effects of the bias boost by making sure that no 
detector’s responsivity was interpolated across a bias boost event and by implementing a model to 
account for the time dependent exponential baseline decay after the bias boost event. Any linear 
change in the responsivity after a bias boost was corrected by the interpolation between reference flashes, 
but residual nonlinear responsivity variations and changes associated with low level radiation exposure, 
such as at the polar horns, were ignored in the flux reconstruction process. It is estimated that not 
accounting for radiation exposure should introduce an uncertainty of less than a few percent in the data. 


VI- 12 


in: 


B.4.c Photon Induced Responsivitv Enhancement 
Point Sources 


As discussed in Section IV. A. 8, analysis after the mission showed that passage over a bright infrared 
source such as Saturn or, more typically, the Galactic plane increased the responsivity of the 100 pm 
detectors by as much as 15%. The enhancement decayed with a long but unknown time constant, thus 
invalidating the basic assumption of a linear change in the responsivity in the time between flashes of the 
internal reference source. 

The sign and magnitude of the effects of photon induced responsivity enhancement on the pho- 
tometry depends crucially on the relative location of the detected source, on the location of the bright 
source of infrared radiation, i.e., of the Galactic plane, and on the time of the flashes of the internal refer- 
ence source relative to the bright source exposure. This is illustrated in Fig. VI.B.2. As an example, if 
the flashes occurred far from the Galactic plane, (flashes (3) and (1) in Fig. VI.B.2. a), the flux of sources 
located on the trailing side of the plane would be overestimated and sources on the leading side of the 
plane would be underestimated. On the other hand, if the flash occurred near the Galactic plane, (e.g., as 
shown in Fig. VI.B.2.a with flashes (2) and (1)) the linear interpolation of the responsivity would lead to 
an underestimate of the flux of the same source. 

The enhancement had not been measured or modeled before flight. In principle, this effect could 
be accounted for on a scan by scan basis by calculating a nonlinear time history for the responsivity 
between flashes of the internal reference source. In practice, however, because the effect was discovered 
only after a significant amount of processing was completed, it was not possible to make a proper correc- 
tion for this effect, and a statistical approach was taken. The basis of the approach is to assume that a 
series of scans were laid down with sufficient regularity that on the average a given area of the sky relative 
to the Galactic plane was affected uniformly. For example, if the flashes were all taken sufficiently far 
from the plane that the effects of the enhancement had decayed, the observed fluxes would resemble 
those in Fig. VLB.2.C. The ratio of fluxes observed on descending scans to those seen on ascending scans 
would then be as indicated in Fig. VLB,2.d. Such a signature was in fact observed at a number of eclip- 
tic longitudes; see Fig. VI.B.3.1. 

Unfortunately, a significant area of the sky was not covered with overlapping descending and 
ascending scans. The areas which were covered with scans in both directions and for which > 50 sources 
were found in bins 20 x 5° in ecliptic longitude and latitude are indicated in Fig. VI.B.4. On the basis 
of this figure, a correction factor was assigned at the locations where at least 50 overlapping values of F D 
and F a , the fluxes measured on descending and ascending scans were measured. The asymmetry about 
the Galactic plane seen in the figure is a result of the manner in which the survey scans were made. For 
example, at ecliptic longitudes 60 to 160°, some of the descending scans ended prior to crossing the 
Galactic plane giving virtually no responsivity enhancement while other descending scans crossed the 
plane resulting in significant responsivity enhancement. Likewise the same affect occurs for the ascending 
scans between ecliptic longitudes of 240° to 340°, giving lower corrections for the negative Galactic lati- 
tudes. The correction was quantized into four non-zero values, ± 2.5% and ± 7.5% on the basis of the 
value of log (Fd/F a \ shown in Fig. VI.B.4. An arbitrary extrapolation of these correction factors was 


VI-13 



IRS 

FLASH 

AMPLITUDE 


(3) 


(b) 


IRS 

FLASH 

AMPLITUDE 


ASSUMED 

RESPONSIVITY 





0 


1 


GALACTIC LATITUDE 


IRS 

FLASH 

AMPLITUDE 


ASCENDING SCAN 


(c) 


QUOTED 

FLUX 


1 I 1 

0 

GALACTIC LATITUDE 


(d) 

ft) 


__j i — i — 

0 

GALACTIC LATITUDE 



DESCENDING 


ASCENDING 



Figure VI B 2 A schematic representation of the effects of photon induced responsivity enhancement. 

8 Sketches (a) and (b) show the true and assumed responsivity for ascending and des- 

cending scans. It is seen that if the first flash of the internal reference source occurs 
before the plane crossing, the sign of the effect is reversed from the case when it occurs 
after the plane crossing. Sketches (c) and (d) show the effect on the quoted fluxes and 
on the ratio of fluxes obtained on descending and ascending scans. 


1 TFI 


VI- 14 


LOG (F D /F A ) LOG (F_ / F 


0.20 


60 < ECLIPTIC LONGITUDE < 80 


BEFORE 



ECLIPTIC LATITUDE 

Figure VI.B.3 Results of comparisons of fluxes at 100 pm of sources observed on descending and 
ascending scans for sources within the ecliptic longitudes indicated. The ordinate is 
^og(F D /F A ) where F A is the flux found on ascending scans and F D is the flux found 
on descending scans. 

, . , . . , * * > * 

T\ ■ A HtXJT y 


T IT i C 







LEGEND: CORRECTION 


ORIGINAL PAGE IS 
QE POOR quality: 

saniuvi )iidii d 3 



sarunvi )iidii 33 


Figure VI.B.5 


The correction factors, assigned on the basis of Fig.VLB.4, used to correct for the 
photon-induced responsivity enhancement at 100 um, See text. 






assumed for areas where sufficient coverage by both ascending and descending scans was missing; the 
resulting corrections are shown in Fig. VI.B.5. Subsequently, the correction was applied to each source 
found at 100 pm on a descending scan and its negative was applied to the sources found at 100 pm on 
ascending scans whether or not it was in a region of overlapping ascending and descending scans. Clearly 
this procedure is very approximate. Figure V1.B.6 shows the results of applying these corrections to the 
same observations presented in Fig. VI.B.3. 

Extended Emission 

The effects of photon induced responsivity enhancement on the extended emission data were 
observed by comparing, pixel by pixel, low resolution (0.5° square pixels) maps of the first sky coverage 
to the last sky coverage. Photon induced responsivity enhancement from bright portions of the Galactic 
plane appears to affect the extended emission in both the 60 and 100 pm bands to a larger extent than it 
does the point sources. 

This problem has not been well enough understood at the release of this catalog (November 1984) 
to properly remove the effect. No correction has been applied to the quoted surface brightnesses, but 
separate data sets for the three sky coverages have been produced. When all of these data sets are 
released, the user will be able to determine if photon induced responsivity enhancement effects are impor- 
tant in a particular area of interest. The zodiacal history file (Section X.D.6) should be consulted to 
determine if the different sky coverages did indeed cover the area in different directions. 

B.4.d Variation of Frequency Response with Total Flux 

The observations shown in Fig. IV.A.4 indicate that in the 60 and 100 pm bands, at least, the fre- 
quency response depends on the brightness of the source. The observations have not been analyzed in 
detail, but sources up to ten times brighter than a Lyr apparently follow the pattern shown by a Lyr. It 
is not clear whether the change for brighter sources comes in the point source response at survey rate or 
in the low frequency responsivity. Thus the photometry of all sources brighter than about 100 Jy at 60 
or 100 pm has uncertainties ranging to 30% at 60 pm and 70% at 100 pm that are not understood. This 
caution must apply equally to the extended emission atlases and to the small extended sources where the 
variations in the responsivity with frequency apparently change sign in going from weak to bright sources 
and from point sources to regions of large scale emission. 

C. Absolute Calibration 

C.l General Philosophy 

The absolute calibration of the IRAS point source observations is tied directly to the absolute cali- 
bration by Rieke et al. (1984) of the ground based photometric system at 10 pm. Specifically, the 12 pm 
IRAS band is calibrated via measurements of a Tau, with the assumption that its absolute flux density is 
as given in Rieke et al. Extrapolation to the 25 and 60 pm bands is achieved using models of stars, nor- 
malized to observations of the Sun. In this latter respect, the IRAS calibration differs in principle from 
the ground based calibrations of Rieke et al. and of Tokunaga (1984) who assumed that the flux density 
of a Lyr between 10 and 20 pm varied as that of a 10,000 K blackbody. The extrapolation of the abso- 
lute calibration from 60 to 100 pm is based on observations and model calculations of asteroids whose 
absolute flux at 60 pm was obtained using the stellar calibration. 


VI-19 


The spectral response of the bands is sufficiently broad (Section ILC.2) that it is necessary to specify 
the continuum energy distribution of the source being observed when defining flux densities at a given 
wavelength. The approach used for IRAS was to assign effective wavelengths of 12, 25, 60 and 100 tun 
for the four bands. The effective bandwidth of each band was then calculated such that the quoted flux 
densities are correct if the source has an energy distribution with a flux per logarithmic frequency interval 
v x / v - X x f x which is constant with frequency v. Any other continuum distribution, and in particular 
that of hot stars, requires a color correction. This color correction, which ranges up to 50% for astro- 
nomically interesting continua, is discussed in detail in Section VI.C.3 below. 

C.2 Point Source Calibration 

The point source amplitudes used to determine the absolute calibration of the IRAS bands were all 
measured using the pointed observation mode, rather than using the results of the survey per se. These 
pointed mode measurements have less intrinsic scatter than measurements obtained during the survey 
since the cross-scan position of a source with respect to individual detectors is well defined and flashes of 
the internal reference source were taken at the same time as the calibration source measurement. Each 
source was crossed at least five times on at least two separate occasions. In addition the reductions of the 
pointed observations were available before the survey processing was completed. Table VLG1 shows, 
however, that statistically the results of the survey reductions and pointed observations agree. The three 
piece fit to the feedback resistor curve shown in Figure VI. A. 1 was assumed throughout. Except for this 

the system was assumed to be linear. 

C.2.a Stellar Calibration 

The absolute calibration at 12 pm was set so that the color corrected flux density of a Tau at that 
wavelength was 448 Jy, in agreement with ground based measurements by Rieke et al. (1984) at 10 pm. 
Alpha Tau was chosen on the basis that it was the primary stellar source in the absolute calibration of 
Rieke et al. and that it was well measured, with high signal to noise ratio, by IRAS. Although the abso- 
lute scale was set using observations of this one star, the IRAS measurements of a significant subset of the 
stars used by Rieke et al. are in excellent agreement with the ground based observations. In Table 
VI.C.2 the flux densities at 12 pm of stars measured using the pointed mode of IRAS are compared with 
the results of Rieke et al. at 10.6 pm. In this table, and in the following discussion, the flux densities 
obtained from IRAS have been color corrected assuming the energy distribution follows that of a hot 
blackbody. The flux densities obtained from the ground based observations have been extrapolated from 
10.6 pm to 12 pm again assuming the energy distribution of the star follows that of a hot 
blackbody. The average ratio of the ground based flux densities to those obtained by IRAS is 1.01 ± 

0 . 01 . 

In order to extrapolate the absolute calibration which was established for the 12 pm band to the 
longer IRAS wavelengths, it is convenient to define a magnitude system for inter-comparisons of the pho- 
tometry in different wavelength bands which normalizes out the energy continuum of hot stars. Because 
the energy distributions of the stars are very nearly black bodies, the magnitude system has been defined 
such that zero magnitude corresponds to a flux density in Janskys: 

UO.O mag) - 1.57 x 1(T 16 x 5^10,000 K) (VI.C.l) 


VI-20 



Table VI.C.l Difference Between Survey and Pointed Observations 


Band 


12[im 


25nm 


60pm 


lOO^im 

Star 

mag 1 

dif 2 

mag 

dif 

mag 

dif 

mag 

dif 

a Tau 

-3.08 

-0.08 

-3.02 

-0.05 

-3.03 

i 

o 

o 

oo 

-2.81 

-0.02 

a Aur 

-1.90 

+0.01 

-1.95 

-0.03 

-1.85 

+0.01 

-1.85 

-0.06 

a CMa 

-1.36 

-0.00 

-1.38 

-0.06 

-1.23 

+0.04 

- 

- 

a CMi 

-0.71 

+0.03 

- 0.72 

+0.00 

-0.72 

-0.02 

- 

- 

P Gem 

-1.21 

-0.01 

-1.19 

-0.02 

-1.22 

-0.05 

-1.19 

+0.19 

a Leo 

+ 1.58 

-0.08 

+ 1.63 

-0.09 

- 

- 

- 

- 

y UMa 

+2.39 

+0.01 

+2.33 

+0.03 

- 

- 

- 

- 

£ UMa 

+ 1.99 

-0.00 

+ 1.93 

-0.09 

- 

- 

- 

- 

a Boo 

-3.22 

-0.07 

-3.09 

+0.01 

-3.02 

+0.09 

-3.02 

+0.21 

y Dra 

-1.44 

-0.01 

-1.50 

-0.03 

-1.42 

+0.06 

-1.48 

+0.07 

a Lyr 

-0.02 

-0.03 

-0.17 

+0.01 

-1.94 

+0.01 

-3.00 

+0.12 

P Gru 

-3.40 

+0.02 

-3.50 

-0.05 

-3.52 

-0.02 

-3.48 

+0.19 

Avg 


-0.019 


-0.030 


+0.005 


+0.10 

Sigma 3 


0.036 


0.040 


0.053 


0.11 


1 mag - the magnitude derived from survey observations. 

2 dif * the difference between the magnitude derived from the 
survey and pointed observations. 

3 sigma - the population standard deviation. 


where B v is the Planck function in Jy sr“ l . For the IRAS effective wavelengths / v [0.00 mag] is 28.3, 
6.73, 1.19 and 0.43 Jy at 12, 25, 60 and 100 pm. On this system a Lyr has a 12 pm magnitude of [12 
pm](a Lyr) - 0.02 mag if it behaves as a 10,000 K blackbody beyond 2.2 pm and [12 pm](a Tau) — 
-3.00 mag. 

The extrapolation of the stellar model to 25 pm was based on the compilation of solar photometry 
presented by Vemazza, Avrett, and Loeser (1976). A comparison of several stellar model calculations 
(Gustafsson, et al 1975; Kurucz, 1979; Bell, 1984) with the data in Vemazza et al predicted a smaller 
color difference, by about 2%, in the models than observed for the Sun. The models did show, however, 
that stars with a wide range of effective temperatures and surface gravities have the same [12 pm] - [25 
pm] and [25 pm] - [60 pm] colors. The following colors, obtained from the observations of the solar 
fluxes, were thus adopted for the average of a set of calibration stars: 


VI-21 




Table VI.C.2 Comparison with Ground-Based Observations 


/ v (Rieke et al.)* 
/ v (IRAS, 1 2 pm) 


a Tau 

1.01 

a Aur 

1.01 

aCMi 

0.98 

P Gem 

1.01 

a Boo 

1.00 

y Dra 

1.05 

a Lyr 

1.00 


"■extrapolated to 1 2pm using hot black body energy distribution 


[12pm] - [25pm] - - 0.03 mag (VI.C.2) 

[25pm ] - [60pm ] — - 0.03 mag 

In Table VT.C.3 the stars used to develop the calibration are listed along with the resultant magnitudes 
from IRAS pointed observations. The stars were selected because reliable models were available for these 
stars and because there was no obvious excess at 60 pm relative to a Tau. 

C.2.b Asteroid Calibration 

The fluxes from the asteroids Hygiea, Europa and Bamberga were observed on eight occasions using 
the pointed mode of IRAS. Ten individual crossing of the focal plane were measured for Hygiea and 
Europa while 20 crossings were recorded for Bamberga. Color temperatures were derived for each obser- 
vation using the stellar calibration at 25 and 60 pm, and the assumption was then made that these color 
temperatures would remain unchanged between 60 and 100 pm. The flux at 100 pm was then calculated 
for each observation from the appropriate blackbody function and the stellar calibration at 60 pm. The 
average of the predicted 100 pm fluxes for the three asteroids was then adopted as the basis of the 100 
pm absolute calibration. The average color temperatures and the ratio of the predicted 100 pm flux to 
that adopted are given in Table VI.C.4. 




Table VI. C. 3 IRAS Magnitudes of Calibration Stars in Pointed Mode 1 


Star 

[12 pm] 

[25 pml 

[60 pm] 

(mag) 

(mag) 

(mag) 

a Tau 

-3.00 

- 2.97 

-2.95 

a Aur 

-1.91 

-1.91 

-1.86 

a Car 

-1.40 

-1.38 

-1.33 

a CMa 

-1.36 

-1.32 

- 1.27 

a CMi 

-0.74 

- 0.72 

-0.70 

P Gem 

-1.20 

-1.18 

-1.17 

a Leo 

1.66 

1.72 

— 

a Boo 

-3.15 

-3.10 

-3.11 

‘See Section VI.C.2.a 

for definition of magnitude scale. 



The asteroids Hygiea, Europa and Bamberga were selected for the basis of the 100 pm calibration 
because the analysis of ground based observations (Lebofsky, 1984) indicated that the standard asteroid 
model was a good fit out to 20 pm and could thus reasonably be expected to fit the longer IRAS 
wavelengths as well. This expectation was in fact borne out by the IRAS observations. 

Several other asteroids were observed with IRAS. These are included in Table VXC.4 in order to 
show the dispersion in the method. It should be emphasized that those selected as the basis of the 100 
pm calibration were those which a priori fit the "standard asteroid model" described below between 12 
and 60 \im. 

As a check on the above procedure and the simple assumptions concerning asteroid colors, the 
observed flux densities for all the asteroids measured already, including the three used in the 60 to 100 
pm extrapolation, were compared to calculations based on the "standard asteroid model" of Morrison 
(1973) and Jones and Morrison (1974). The infrared emissivity of the surface was taken to be 0.9 
independent of wavelength and the thermal modeling constant p was 0.9; the albedo was taken from the 
TRIAD file (Zellner, 1979). The temperature distribution on the surface was assumed to follow: 

T = T ss cos 0 25 z (V7.C.3) 


VI-23 



Table VLC.4 

Color Temperatures 

of Asteroids between 25 and 60 pm 

Asteroid 

Color Temp. 1,3 

Obs./Pred. Flux 2 3 


K 


Europa 

228 ±15 

0.97 ±0.12 

Bamberga 

243 ±10 

1.01 ±0. 1 5 

Hygeia 

232± 16 

1.01 ±0.13 

Eukrate 

292 ±29 

1.19±0.12 

Egeria 

259± 19 

0.95 ±0.1 3 

Ceres 

234 ±12 

1.13±0.12 

Flora 

297 ±40 

1.43 ±0.25 

Berberic 

301 ±32 

1.52±0.26 

Pallas 

232 ±7 

1.05 ±0.08 


1 The color temperature is based on the 25 and 60 pm fluxes 

2 The predicted 100 pm flux is extrapolated from the 60 pm flux density obtained from the stellar 
calibration and the color temperature. 

3 The uncertainties are the population standard deviations 


where T ss is the subsolar point temperature and z is the zenith angle of the Sun. The temperature of the 
dark side was taken to be 0 K; this assumption does not lead to a significant error since mainly the sunlit 
sides of the asteroids were observed by IRAS. 

The asteroid diameters were adjusted for each observation to match the 60 pm stellar calibration 
exactly, i.e. all ratios of observed/model fluxes at 60 pm are identical to unity. The 100 pm calibration 
was adjusted such that for the asteroids the mean of the ratio of the observed flux to the model flux was 
equal to unity. The resultant ratios of observed fluxes to those derived from the model are given in Table 

VI.C.5. 

C 2 c Estimated Accuracy 

The estimated absolute accuracy of the stellar plus asteroid calibration relative to the 10 pm ground 
based calibration is 2%, 5% and 5% at 12, 25 and 60 pm, based on the uncertainty in the stellar models 
and the scatter in the standard stars, and 10% at 100 pm based on the uncertainty of the asteroid model 
extrapolation. The stated accuracy of the 10 pm absolute calibration is 3% (Rieke et al. 1984). The 


VI-24 




Table VI.C.5 Ratio of Observed to Model Fluxes of Asteroids 


Asteroid Observed/Model Fluxes* 



25 pm 

60 pm <2) 

100 pm 

Europa 

1.026 

1.000 

1.000 


1.026 

1.000 

0.965 

Hygiea 

1.052 

1.000 

1.005 


1.062 

1.000 

1.015 

Bamberga 

1.000 

1.000 

0.980 


1.000 

1.000 

1.035 


1.082 

1.000 

1.015 


1.066 

1.000 

1.025 

Average Asteroid 

1.039 

1.000 

1.000 (3) 

Population Dispersion 

±0.031 


±0.029 


**' Observed 25 and 60 (im fluxes are based on stellar calibration. The model fluxes are based on 
the standard asteroid model (Morrison, 1973; Jones and Morrison, 1974) with the beaming fac- 
tor p - 0.9, emissivity e - 0.9, and albedo from the TRIAD file (Zellner, 1979). 

* 2) The asteroid diameter is used to normalize the model flux to the observed flux at 60 pm, i.e. 
observed/model — 1 .000 for each asteroid. 

* 3) By definition of the calibration procedure, the mean ratio of observed/model — 1 .000. 


error introduced into the absolute calibration due to uncertainties in the spectral pass-bands is estimated 
to be less than 4% for objects as warm, or hotter than asteroids. 

The ratios listed in Table VI.C.5 indicate that, for the three asteroids used in the calibration, the 
flux at 25 pm is higher than expected on the basis of the stellar calibration and the asteroid models by 4 
± 1% once the fluxes have been normalized at 60 pm. This is well within the expected accuracy of the 
stellar and asteroid models. The dispersion of the ratio of observed to model fluxes at 100 pm is only 
3%, the mean being identical to unity by definition of the asteroid calibration procedure. The good fit at 
25 pm and the small dispersion at 100 pm gives confidence that the extrapolation of these asteroids is a 
valid procedure. 


VI-25 


c 3 



TABLE Suppl.VI.C.6 Color Correction Factors, K 1 



INTRINSIC 

RATIO OF FLUX DENSITIES 


CORRECTION FACTOR 


5 


POWER LAW 2 

BEFORE COLOR-CORRECTION 







a 

fX 12 pm) 

fX 25 pm) 

/„( 60 pm) 

K(12 pm) 

K(25 pm) 

K(60 pm) 

K{100pm) 



fX 25 pm) 


fX lOOpm) 



-3.0 

0.113 

0.063 

0.21 

0.91 

0.89 

1.02 

1.02 



-2.5 

0.162 


0.275 

0.92 

0.91 

1.00 

1.01 



-2.0 

0.232 

0.164 

0.355 

0.94 

0.93 

0.99 

1.00 



-L5 

0.333 



0.97 

0.96 

0.99 

1.00 

= 


-1.0 

0.480 



1.00 

1.00 

1.00 

1.00 


! 

-0.5 

0.694 


0.786 

1.04 

1.04 

1.02 

1.00 


1 


1.005 

1.045 

1.037 

1.10 

1.10 

1.05 

1.01 


! 

0.5 

1.459 

1.642 

1.378 

1.17 

1.16 

1.09 

1.02 

■ 


1.0 

2.123 

2.567 

1.843 

1.25 

1.23 

1.15 

1.04 



1.5 

3.094 

3.992 

2.484 

1.35 

1.32 

1.23 

1.06 



2.0 

4.519 


3.373 

1.47 

1.41 

1,32 

1.09 



2.5 

6.610 

iHfe 

4,617 

1.61 

1.53 


1.12 

- 


3.0 

9.681 

14.475 

6.370 

1.78 

1.67 

■Hi 

1.16 

_ 

s 


RATIO OF FLUX DENSITIES 


CORRECTION FACTOR 


i 


INTRINSIC 

BEFORE COLOR-CORRECTION 





- 

. 









| 

TEMP (K) 

A3 12 nm) 
fX 25 pm) 

/ v (25 pm) 

/ v (60 pm) 

/ v (60 pm) 
/ V (100 pm) 

K(12 pm) 

K(25 pm) 

K(60 pm) 

K(100 pm) 


i 

10000 

4.345 

6.050 

3.350 

1.45 

1.41 

1.32 

1.09 

- 

: 

5000 

4.172 

5.931 

3.327 

1.43 

1.40 

1.32 

1.09 


; 

4000 

4.086 

5.872 

3.316 

1.42 

1.40 

1.31 

1.09 


i 

3000 

3.944 

5.773 

3.297 

1.41 

1.39 

1.31 

1.09 


1 

2000 

3.666 

5.578 

3.259 

1.38 

1.38 

1.31 

1.09 


1 

1000 

2.891 

5.005 

3.145 

1.27 

1.34 

1.29 

1 .08 


* 

800 

2.545 

4.730 

3.088 

1.22 

1.32 

1.28 

1.08 


; 

600 

2.036 

4.287 

2.995 

1.15 

1,29 

1.27 

1.08 


z 

500 

1.692 

3.950 

2.920 

1.09 

1.26 

1.26 

1.08 


= 

400 

1.272 

3.478 

2.810 

1.01 

1.22 

1.24 

1.08 


] 

i 

300 

0.785 

2.780 

2.630 

0.92 

1.15 

1,21 

1.07 



290 

0.734 

2.693 

2.606 

0.91 

1.15 

1.21 

1.07 


= 

280 

0.684 

2.602 

2.580 

0.90 

1.14 

1.20 

1.07 


* 

270 

0.633 

2.506 

2.553 

0.89 

1.13 

1.20 

1.07 


“ 

260 

0.583 

2.407 

2.523 

0.88 

1.12 

1.19 

1.07 


- 

250 

0.534 

2.304 

2,491 

0.87 

1.11 

1.19 

1.07 



240 

0.486 

2.196 

2.457 

0.86 

1.09 

1.18 

1.07 


- 

230 

0.438 

2.084 

2.420 

0.85 

1.08 

1.18 

1.07 



220 

0.392 

1.967 

2.381 

0.85 

1.07 

1.17 

1.07 


* 

210 

0.347 

1.845 

2.338 

0.84 

1.06 

1.16 

1.06 


- 

200 

0.304 

1.719 

2.291 

0.83 

1.04 

1.16 

1.06 


= 

190 

0.263 

1.589 

2.240 

0.83 

1.02 

1.15 

1.06 


■ 

180 

0.224 

1.455 

2.184 

0.83 

1.0! 

1.14 

1.06 



170 

0.188 

1.317 

2.124 

0.83 

0.99 

1.13 

1.06 



160 

0.154 

1.176 

2.057 

0.84 

0.97 

1.12 

1.06 


= 

150 

0.124 

1.034 

1.983 

0.85 

0.95 

1.11 

1.05 



140 

0.097 

0.892 

1.901 

0.87 

0.93 

1.09 

1.05 


- 

130 

0.073 

0.751 

1.810 

0.90 

0.91 

1.08 

1 .05 


~ 

120 

0.053 

0.614 

1.709 

0.94 

0.89 

1.06 

1.04 


2 

110 

0.036 

0.484 

1.595 

1.01 

0.86 

1.04 

1.04 


| 

100 

0.023 

0.363 

1.468 

1.12 

0.84 

1.02 

1.04 


1 

95 

0.018 

0.307 

1.400 

1.19 

0.83 

1.01 

1.03 


m 

90 

0.014 

0.256 

1.326 

1.28 

0.83 

1. 00 

1.03 


d 

85 

0.010 

0.208 

1.249 

1.39 

0.82 

0.99 

1.03 


- 1 

80 

0.007 

0.165 

1.168 

1.54 

0.81 

0.97 

1.02 



75 

0.005 

0.127 

1.082 

1.74 

0.81 

0.96 

1.02 



70 

0.003 

0.095 

0.993 

2.01 

0.81 

0.95 

1.01 


- 

65 

0.002 

0.067 

0.898 

2.40 

0.82 

0.94 

1.01 


i 

60 

0.001 

0.045 

0.801 

2.97 

0.83 

0.93 

1.00 


i 

55 

__ 

0.028 

0.700 

3.86 

0.86 

0,92 

1.00 


= 

50 


0.0 1 6 

0.597 

5.35 

0.90 

0.91 

0.99 



45 


0.008 

0.493 

8.09 

0.97 

0.92 

0.98 



40 

— 

0.003 

0.391 

13.79 

1.08 

0.93 

0.98 



1 f v [actual] - / v [quoted]/# 

0 $ 

2 /.-V 


spectral response given in Table Suppl.II.C.5. 


VI-26 


ORIGINAL PAGE IS 
PC POOR QUALITY 


I II: 




An alternative explanation of the results in Table VLC.5 is that the difference between the observed 
and predicted 60 to 25 pm flux ratio could be a manifestation of a short wavelength leak in the 60 pm 
spectral response such as reported from the preflight component tests (Section II.C.4). An examination of 
the 60 pm to 12 and 25 pm colors of the calibration stars (Table VI.C.3), however, shows no change in 
color larger than 0.02 mag for different effective stellar temperatures; whereas the colors should change by 
0.04 mag for effective temperatures varying from 4000 to 10,000 K, if the short wavelength leak were 
real. 

C.3 Color Corrections 

In the catalog, the flux densities have been quoted for the effective wavelengths and an input energy 
distribution which is constant in the flux per logarithmic frequency interval (essentially the flux per 
octave); i.e., the flux density with frequency v goes as / v <r V _1 while the flux density with wavelength X 
goes as/*, a A. . No loss in generality in incurred by such a procedure, nor does it prejudice subsequent 
correction for the correct source distribution. 

If the input energy distribution is not constant in flux per octave a correction, the "color correction", 
must be applied to the quoted flux densities. This correction depends on the shape of the intrinsic energy 
distribution and on the details of the wavelength response of the system. The color corrections for a 
number of input energy distributions are given in Table VI.C.6. 

The flux F measured by a detector is given by: 


-fv, [actual] f (fv/fv) [actual] RJv 

(VI.C.4) 

-fv, [quoted] f (fv/fv) [quoted ]Rvdv 

(VI.C.5) 

fv, [quoted] r 

— Jr J (fv/fv) [actual] R v dv 

(VI.C.6) 


In these equations, Vo is the effective frequency of 25, 12. 5, and 3 x 10 12 Hz corresponding to the 
effective wavelengths of 12, 25, 60 and 100 pm, f v is either the actual or quoted flux density of the 
source, (f v /f v ) is the flux density normalized to the effective frequency of the band and R v is the relative 
system response listed in the last column of Table II.C.5. Equations VI.C.4 and VI.C.6 show that the 
true flux density at v 0 is given by: 

fv, [actual] =/ Vo [quoted ]/AT (VT.C.7) 


where, from Eq. VI.C.4 and VI.C.5, 

K ~\S (fv/fv) [actual] /Mv)/| / (fv/fv) [quoted] /MvJ (V7.C.8) 

It should be emphasized that there is no prejudice in the procedure or assumptions. The uncer- 
tainty in the quoted flux densities results directly from the lack of knowledge of the spectral response of 
the system. Deriving the true flux densities requires a knowledge of the intrinsic energy distribution of 
the astronomical sources. 


VI-27 


The sensitivity of the color corrections to uncertainties in the spectral band passes was checked by 
numerical simulations. For many sources of interest in the catalog, the 60 pm band flux densities are 
especially sensitive to errors caused by lack of knowledge of the spectral pass bands. Specifically, in this 
band, uncertainties in the short wavelength parameters of the transmission of the optics, and in the detec- 
tor efficiency, affect the response for extremely cold objects (T — 60 K) in a significantly different way 
than for the stellar energy distribution of the calibrating sources. The numerical tests show that the sensi- 
tivity of the color corrections to wavelength shifts of the entire response is approximately 7%/pm; pre- 
launch measurements of the spectral shape should be accurate to ~ 0.3 pm. On the other hand, if the 
effective responsivity of the system increases by an additive 1% over the entire bandwidth, starting at 24 
pm, the calibration for 60 K sources changes by 16% relative to the stellar calibration. Uncertainties of 
this magnitude in the effective response are the maximum expected from pre-flight measurements. Errors 
in the long wave cutoff of the 100 pm band could also result in significant errors in the flux densities of 

objects colder than 30 K. 

C. 4 Absolute Calibration of Extended Emission 

The absolute calibration of the extended emission observed by IRAS was based directly on the abso- 
lute calibration of the point sources. Thus the value of the equivalent absolute flux density assigned to 
the flashes of the internal reference source was carried directly to the calibration of the extended emission 
with the additional corrections for the frequency response and effective detector areas; see Sections IV.A.3 

and IV.A.4. 

D. Comparison of IRAS Observations with Ground Based Observations 

In Table VI.D.l a comparison is given between the IRAS measurements of selected stars made in 
the pointed mode and the results of ground based measurements by Tokunaga (1984) and Rieke et al. 
(1984). At 12 pm, the IRAS measurements agree with the ground based magnitudes within 1%. The 

difference is largely due to y Dra. 

If a Lyr is omitted, the average difference (and the uncertainty in the mean) between the IRAS 
magnitudes and those of Tokunaga (1984) at 25 pm is -0.01 ± 0.02 while the difference between the 
IRAS magnitudes and those of Rieke et al. (1984) is 0.04 ± 0.01. The 0.02 mag difference in zero point 
between the IRAS magnitudes and those of Tokunaga has been incorporated in these differences as well 
as the difference of 0.03 mag assumed in the 12 to 25 pm color. A corollary of this good agreement is 
that the excess in a Lyr does not start until longward of the long wavelength cutoff of the filters used in 
the ground based observations. 

The most accurately calibrated measurements to compare with the IRAS 60 and 100 pm observa- 
tions are those of the planets Uranus and Neptune by Hildebrand et al. (1984) using NASA’s Kuiper Air- 
borne Observatory (KAO). Hildebrand et al ' s measurements of Uranus are compared to observations 
taken early in the IRAS mission in Figure VI.D.l. It is seen that the flux densities obtained by IRAS 
are ~ 20% lower than those obtained by Hildebrand et al. The cause of the discrepancy is not under- 
stood at this time. 

It should be noted that the temperature of Uranus is ~ 60 K. The planet thus has an energy distri- 
bution quite different from that used in the calibration procedure. Although a short wavelength leak in 


VI-28 


the 60 pm filters as large or larger than 20% would account for this discrepancy, the good fit of the 
asteroid and stellar models at 25 and 60 pm and the near constant stellar colors with effective tempera- 
ture appear to rule out a short wavelength leak of this magnitude. It should be noted, however, that all 
the measurements from the KAO derive their absolute calibration from observations of Mars and a single 
model thereof. Thus the internal agreement of the KAO observations at different wavelengths does not 
reflect independent calibrations. An analysis of the final IRAS survey flux densities, as contrasted to 
those obtained from the pointed observations, is given in Chapter VTI. 


Table VI.D.l Comparison with Ground-Based Observations 




[12 pm] 



[25pm] 


Star 

IRAS 

Ay 

&R 

IRAS 

Ay 

Ar 


(mag) 

(mag) 

(mag) 

(mag) 

(mag) 

(mag) 

a Tau 

-3.00 

-0.01 

-0.01 

- 2.97 

+0.07 

+0.07 

a Aur 

-1.91 

-0.01 

-0.01 

-1.91 

-0.03 

-0.02 

a CMa 

-1.36 

-0.04 

-- 

-1.32 

-0.01 

-- 

a CMi 

-0.74 

-0.00 

-0.02 

-0.72 

-0.04 

-0.01 

P Gem 

-1.20 

-0.02 

-0.01 

-1.18 

-0.02 

-0.04 

a Boo 

-3.15 

-0.00 

-0.00 

-3.10 

-0.02 

+0.06 

y Dra 

-1.43 

-- 

-0.05 

-1.47 

-- 

+0.04 

a Lyr 

0.02 

0.00 

+0.00 

-0.18 

+0.23 

+0.23 


Ay- and A R are the differences between the IRAS [12 pm] and [25 pm] magnitudes and the correspond- 
ing values from Tokunaga (Ay) and Rieke et ah (A/*). Differences are in the sense A — [IRAS] - 
[ground-based]. In displaying these differences, the 10.1 pm magnitudes of Tokunaga have been adjusted 
by 0.02 mag to account for a different zero point convention, and the 20.1 and 21 pm magnitudes of 
Tokunaga and Rieke et al. have been similarly adjusted by 0.05 and 0.03 mag to agree with Eq. (VI.C.2). 


VI-29 





Authors: 


G. Neugebauer, S. Wheelock, F. Gillett, H. H. Aumann, T. N. Gautier, F. J. Low, P. Hacking, M. 
Hauser, S. Harris, P. Clegg. 


References 

Allen, D.A. 1970, Nature, 227, 158. 

Allen, D.A. 1971, in Physical Studies of Minor Planets, editor T. Gehrels, Univ. Ariz. Press, p. 41. 

Bell, R. 1984 private communication. 

Gustafsson, Bell, Ericksson, Nordlund, 1975, Astr. Ap., 42, 407. 

Hauser, M. et al. 1984, Ap. J. (Lett.), 278, LI 5. 

Hildebrand, R.H., Lowenstein, R.F., Harper, D.F., Orton, G.S., Keene, J., and Whitcomb, S. 1984, 
pre-print 

Jones, T.J., and Morrison, D. 1974, A. J., 79, 892. 

Kurucz 1979, Ap. J. (Suppl.), 40, 1. 

Lebofsky 1984, personal communication. 

Low, F.J., et al. 1984, Ap. J. (Lett.), 278, LI 9 
Morrison, D. 1973, A. J., 194, 203. 

Rieke, G.H., Lebofsky, M. and Low, F.J. 1984, preprint. 

Tokunaga, A. 1984, A. J., 89, 172. 

Vemazza, J.E., Avrett, E.H., and Loeser, R. 1976, Ap. J. (Suppl.), 30, 1. 

Zellner, B. 1979 in Asteroids, editor T. Gehrels, Univ. Ariz. Press, p. 1011. 


VI-31 



VII. ANALYSIS OF PROCESSING 


A. Overview 

The purpose of this chapter is to describe the general characteristics of the catalogs of point and 
small extended sources and of the sky brightness images. The distribution on the sky of different types of 
infrared sources is presented not for scientific purposes, but to provide a context within which to interpret 
the characteristics of specific objects (Section VTI.B). Detailed discussions are presented of the positional 
and photometric accuracy of the point sources (Sections VII.C,D). Although the important questions of 
the completeness and reliability of the catalogs are deferred until the next chapter, some of the results of 
the processing relevant to these topics are discussed in Sections VII.E and VII.F. An overview of the 
results of associating IRAS sources with objects appearing in other catalogs is given in Section VII.G. 
Because it is crucial that the user gain an understanding of the various flags that accompany the observa- 
tions of a source, the properties of these flags, in particular the threshold values that indicate that the 
measurement of a source may be suspect, are described in Section VII.H. The properties of the catalog 
of small extended sources are presented in Section VII.I. The extended emission images are discussed in 
Section VII.J. 

B. General Statistics of the Point Source Processing and Catalog 
B.l The Generation of Reliable Point Sources 

The complementary aims of completeness and reliability resulted in the design of detection and 
confirmation programs that acted as a series of filters to weed out spurious sources at each level of the 
data processing. Because the detection thresholds were set low to improve completeness, subsequent 
stages of confirmation were used to establish the reliability of the IRAS sources. 

Approximately 500,000 events per day triggered the square-wave filter detection algorithm (Section 
V.C.2). Approximately 20 to 40% of these detections passed the signal-to-noise and correlation 
coefficient thresholds. The largest fraction of the detections that were accepted, about one-third of the 
total, came at 25 pm. Detections at 12 pm and 100 pm, each contributed one-quarter of the total 
number, while the remainder, about one-sixth of the total number, came at 60 pm. About two-thirds of 
the accepted detections were sightings at different wavelengths of approximately 25,000 seconds- 
confirmed band-merged sources produced daily. About one-third of the seconds-confirmed sources were 
used in the roughly 4,000 hours-confirmed sources (hereafter denoted as HCONs) generated per day of 
surveying. 

Hours-confirmed sources were placed into a Working Survey Data Base (WSDB) to await possible 
weeks-confirmation. At the completion of the initial data processing the WSDB contained about 1.2 mil- 
lion hours-confirmed sources spread over the entire sky. Of these, some 317,000 failed to be matched 
with a weeks-confirming partner during the initial processing. The. remaining 885,000 HCONs were 
combined by the weeks-confirmation processor during the normal processing into 319,000 distinct 
sources each consisting of two or more HCONs. The clean-up processing discussed in Section V.H.2 
forced close neighboring sources into single weeks-confirmed sources, decreasing the number of distinct 
objects while increasing the average number of HCONs per source. After clean-up there were about 


VII- 1 


304,000 single HCONs and 314,000 weeks-confirmed sources, each consisting of two or more HCONs. 
The breakdown of the catalog by numbers of HCONs is given in Table VII.B.l. 

There were three reasons why even a weeks-confirmed source could be excluded from the catalog. 
First, the source could fail to satisfy the criteria of having at least one high quality flux in one band or 
two or more moderate quality fluxes in adjacent bands (Section V.H.5); some 12,000 sources were 
rejected for this reason. Second, the high source density processing could eliminate a source due to the 
strict selection rules imposed in such areas (Section V.H.6 and VII.E. l.b); about 57,000 sources which 
had all survived the first test were rejected by these criteria. Third, a source could be deleted because it 
was a spurious object produced near a bright source (Section VII.E.4); there were 22 such sources outside 
of the Galactic plane. Of the 314,000 weeks-confirmed sources in the WSDB, 245,839 survived all these 
criteria to appear in the catalog. 

B. 2 Distribution of Sources in the Catalog 

As mentioned in Section I.C, there is a meaningful astrophysical classification for almost all of the 
IRAS sources. Objects detected at 12 pm and having / v ( 12 pm) > / v (25 pm) are typically stars (Fig. 
I.C.2, type S in Table VII.B.2). Objects detected at 60 pm and having f v (6 0 pm) > / v (25 pm) are either 
galaxies if they are located away from the Galactic plane or warm Galactic objects if located near the 
plane (type G, Fig. I.C.3). Objects detected only at 100 pm are predominantly cirrus sources, although 
some at very high Galactic latitudes may be external galaxies (type C, Fig. I.C.4). Sources meeting none 
of these criteria are classified "other" and lie mostly in the Galactic plane. Figure VII.B. 1 shows histo- 
grams of the distribution of the four types of sources as a function of Galactic latitude. 

All cataloged sources can also be classified into 15 different spectral groups depending on which 
bands have measurements of either moderate or high quality. The numbers in each category are given in 
Table VII.B.2, where a code has been used to specify which bands were reliably detected. If the presence 
or absence of a measured flux density at 12, 25, 60 and 100 pm is considered in that order to correspond 
to a "1" or "0" in a four digit number, then a source measured only at 12 pm is denoted by "1000", while 
one measured at 12 pm and 60 pm is denoted by "1010", etc. The distribution on the sky of each of 
these color combinations is shown in a series of equal area projections in Galactic coordinates. These 
figures are shown in an appendix at the end of the chapter (Figs. VTI.Ap.1-15). 

C. Positional Accuracy 

C. 1 Positional Accuracy of Catalog Sources 

Stars from the Star Catalog of the Smithsonian Astrophysical Observatory (1966, henceforth the 
SAO Catalog) and UGC galaxies with accurately determined positions (Dressel and Condon 1976, hen- 
ceforth the Dressel and Condon Catalog) were used to investigate the accuracy of the positions quoted in 
the IRAS point source catalog. 

Small areas of half-width 20" in-scan by 75 " cross-scan around sources brighter at 12 pm than at 
25 pm and located more than 20° from the Galactic plane, were searched for SAO stars. Stars with spec- 
tral class O, B, and A, and those with no spectral classification were specifically excluded. Inadvertently, 
stars of spectral types MA, MB, NA and NB were also excluded. This reduced the size of the sample of 
stars used, but, as a subsequent examination revealed, had no effect on the statistics reported below. 


VII-2 



Table VII.B.l Number of HCONs in WSDB and Final Catalog 


Number of 
HCONs 

Number 
of WSDB 
Sources 

Number of 
Sources after 
Clean-up 

Number of 
Catalog sources 

1 

317,764 

304,334 

0 

2 

129,336 

122,705 

76,724 

3 

150,545 

= 151,039 

132,974 

4 

27,233 

28,028 

24,902 

5 

6,906 

7,289 

6,479 

6 

2,731 

2,958 

2,617 

7 

1,565 

1,730 

1,586 

8 

406 

456 

433 

9-23 

100 

124 

124 

Total weeks- 




confirmed 

318,822 

314,329 

245,839 

(> 2 HCONs) 





Combination 

Table VII.B.2 

Number 

Spectral Classification of Catalog Sources 

Type S Type G Type C 

Type O 

1000 

67,332 

57,315 



10,017 

1100 

67,015 

61,776 



5,239 

1110 

13,233 

9,387* 

4,114* 


1,454 

mi 

6,343 

2,287* 

4,583* 


290 

0100 

3,936 




3,936 

0110 

3,642 


3,470 


172 

0111 

3,873 


3,860 


13 

0010 

19,264 


18,494 


770 

0011 

22,702 


22,197 


505 

0001 

33,146 



33,146 


1101 

2,025 

1,478 



547 

1010 

1,100 

647* 

1,062* 


33 

1011 

520 

297* 

515* 


3 

1001 

1,207 

905 



302 

0101 

501 




501 

Total 

245,839 

134,092* 

58,295* 

33,146 

23,782 

*3,476 sources 

are both Type S and G categories. 





VII-3 







Figure VII.B. 1 The number density (per sq. deg) of four kinds of IRAS source (see text) is shown as 
a function of Galactic latitude, averaged over all Galactic longitudes. 


Areas of the same size around sources detected at 60 pm, which were brighter at 60 pm than at 25 pm, 
and located more than 30° from the Galactic plane, were searched for a galaxy in the Dressel and Con- 
don catalog. If there was more than one cataloged star or galaxy within the search box of the IRAS 
source, the source was not used. Matches with position differences greater than seven times the a priori 
uncertainty on either axis were also discarded. No galaxies and fewer than 0.1% of the stars were rejected 
for this reason. 

The analysis was done separately for bright sources (flux densities greater than 1.2 and 1.9 Jy at 12 
and 60 pm respectively) and faint sources (flux densities lower than the above limits). The number of 
such sources in each sample are given in Table VII.C.l. 


VIM 







Table VII.C. 1 

Absolute Position Difference Statistics 





In-Scan 



Cross-Scan 



Number 

Mean 

Population 


Mean 

Population 



N 

difference a 

i 2 /n 

difference 

a 

x 2 /n 



(") 

(") 


(") 

o 


SAO Stars 








Bright 

4757 

0.0 

1.9 

0.42 

0.1 

8.4 

0.55 

Faint 

10558 

0.0 

2.8 

0.65 

0.2 

15.6 

0.87 

Galaxies 








Bright 

340 

-1.2 

5.3 

1.04 

1.0 

12.8 

1.30 




(3.5*) 



(12.2*) 


Faint 

814 

-1.2 

5.7 

0.91 

0.4 

15.2 

1.00 




(4.1*) 



(14.7*) 



*4" uncertainty removed. 


C.l.a Accuracy of the Absolute Positions 

The absolute positional differences between the 1 2 pm sources and the associated stars are shown in 
Figs. VII.C. 1 and VII.C.2 for the in-scan and cross-scan directions. The same information is given for 60 
pm sources and the associated galaxies in Figs. VII.C. 3 and VII.C.4. In the top panel of each figure a 
histogram of the number of bright sources is plotted as a function of the absolute position difference. 
The bottom panel repeats this plot for the fainter sources. In each case an equal-area Gaussian distribu- 
tion with the mean and standard deviation of the sample is plotted for comparison. From these figures it 
is apparent that the in-scan errors are reasonably well represented by Gaussian distributions. The cross- 
scan errors are less Gaussian, showing a more concentrated center and more extended tails. 

The mean and population standard deviation of the positional differences for all sources in the sam- 
ples described above are given in Table VII.C. 1. Because the SAO positions are more accurate than the 
IRAS positions, the listed discrepancies should be representative of the IRAS position errors for sources 
detected at the short wavelengths. Since the rms position errors in the Dressel and Condon catalog are 4" 
in each direction, i.e., approximately the same as the IRAS in-scan errors, it is necessary to correct the 
statistics for this additional uncertainty. The estimates obtained by subtracting the 4" errors in quadra- 
ture from the calculated standard deviations in both directions are given in parentheses in the table. No 
account has been made for any offset of the IRAS source from the optical nucleus of the galaxy. Any 
such effect would cause an overestimate of the IRAS position errors. 

The mean positions of the IRAS stars do not deviate significantly from the SAO positions. This is 
not surprising since the in-flight calibration of the IRAS focal plane geometry used SAO stars detected at 
12 pm to determine the geometric position of the infrared focal plane with respect to the visible star sen- 
sors (Section V.D.3). There is also no significant deviation of the cross-scan position of the IRAS galaxy 
sample based on the galaxy positions. Table VII.C. 1 does, however, show a small, but statistically 


VII-5 





Figure VII.C. 1 Position differences for the IRAS stellar sources associated with SAO stars in the in- 
scan direction. See text. 




1 CROSS-SCAN DIFFERENCE) (arc sec) 

Figure VII.C.2 Position differences for the IRAS stellar sources in the cross-scan direction. See text. 








50 



| IN-SCAN DIFFERENCE [(arc sec) 



| IN-SCAN DIFFERENCE! (arc sec) 

Figure VII.C.3 Position differences for the galaxies associated with Dressel and Condon galaxies in the 
in-scan direction. See text. 




Figure VII.C.4 Position differences for the galaxies associated with Dressel and Condon galaxies in the 
cross-scan direction. See text. 


VII-7 






significant discrepancy for the in-scan positions of the galaxies. This small error is consistent with the 
differences in the positions of the seconds-confirmed sightings in individual wavelength bands of sources 
seen at multiple wavelengths. While there is no significant discrepancy between the 12 pm and 25 pm 
positions, there is an 0.8" discrepancy between the in-scan positions of sources as measured at 12 pm and 
60 pm, and a 2.4" discrepancy between the in-scan positions of sources as measured at 12 pm and 100 
pm. A 0.2% error in the image scale of the telescope could account for this effect. Because the band- 
merging process used positional information from all detected bands for a source, multiband sources 
would suffer least from this error; sources detected only at 100 pm could suffer from the full 2.4" error. 

The standard deviations of the position errors, given in Table Vn.C.l, show that the absolute posi- 
tion errors are quite small. The in-scan position errors depend, as expected, on source brightness and 
wavelength. The cross-scan positions of the brighter sources are significantly better than those of the faint 
sources, in large part because these sources were detected at multiple wavelengths. The difference 
between the errors in the bright stars and the errors in the galaxies is consistent with the expected 
diffraction effects and detector sampling rates. 


C.l.b The Quoted Position Uncertainties 

The quality of the IRAS position uncertainties is shown in Fig. VII.C.5 for the in-scan and cross- 
scan directions. Separately plotted for each of the bright and faint source samples are the mean absolute 
position differences of those sources as a function of the quoted standard deviation. For comparison, the 
value of this same quantity for a Gaussian distribution of position errors would be 




— a iras for stars an d 

71 


V? 


— (o}ras+Gdc)' a for galaxies. 

jt 


(vn.c.i) 


These relations are shown as a solid line in each figure. The difference in form for the galaxies is due to 
the 4" uncertainty, Gdc, in the Dressel and Condon catalog. From the figures it is evident that the IRAS 
position uncertainties are accurate estimates of the position errors only for small values and a consider- 
able overestimate for large values. 

The overall quality of the quoted error estimates is measured by the x 2 parameter, defined as 


x 2 -X 


if 


(XlRAS~ x SA0) 2 

2 

&IRAS 


for stars 


* 2 -x 

N 


( x iras~ x dc ) 2 

(<5iras+Gdc) 


for galaxies 


(VII.C.2) 


VII-8 


MEAN ABSOLUTE CROSS-SCAN DIFF MEAN ABSOLUTE IN-SCAN DIFF MEAN ABSOLUTE CROSS-SCAN DIFF MEAN ABSOLUTE IN-SCAN DIFF 

(arc sec) (arc sec) (arc sec ) ( arc sec ) 



IRAS QUOTED UNCERTAINTY (ore sec) 





Figure VII.C.5 


Observed position differences vs. the quoted uncertainties. The top panels are for 
stars and the bottom are for galaxies. See text. 


VII-9 







Values of x 2 per degree of freedom are given in Table V1I.C.1 for the various samples. As can be seen 
from the table, the position uncertainties are overestimated for the stellar sources in both the in-scan and 
cross-scan directions, while the uncertainties are reasonably estimated for the galaxy population, as a 
whole. 

It should be borne in mind that the observed position errors are not due entirely to pointing recon- 
struction errors. Detection timing errors form a significant component in the in-scan direction, so that 
the IRAS positions of the brighter stars are determined more precisely than those of the faint stars 
because of the greater accuracy of the timing of the detection of high signal-to-noise objects. In the 
cross-scan direction the uniform component of the position uncertainty due to the rectangular aspect of 
the detectors often dominated the errors of position reconstruction. To minimize this effect in assessing 
the performance of the pointing reconstruction processor only triple (edge) detections were used. In these 
cases the image of a source traversed the focal plane in the narrow region of overlap between three detec- 
tors in a single band. 

C.2 Accuracy of Scan-by-Scan Pointing Reconstruction 

As a check on the quality of the pointing reconstruction for each scan, the positions of seconds- 
confirmed band-merged sources were compared with the positions of a preselected set of standard stars 
(Sections V.B and V.D.4). This set was composed primarily of K stars selected from the SAO Catalog 
excluding stars within 2.5“ of the Galactic plane or within 3' of another detectable star. 

It should be borne in mind that the observed position errors are not due entirely to reconstruction 
errors. Detection timing errors form a significant component in the in-scan direction, so that the IRAS 
positions of the brighter stars are determined more precisely than those of the faint stars because of the 
greater accuracy of the timing of the detection of high signal-to-noise objects. In the cross-scan direction 
the uniform component of the position uncertainty due to the rectangular aspect of the detectors often 
dominated the errors of position reconstruction. To minimize this effect in assessing the performance of 
the pointing reconstruction processor only triple (edge) detections were used. In these cases the image of 
a source traversed the focal plane in the narrow region of overlap between three detectors in a single 
band. 


The differences between the IRAS and SAO positions after seconds-confirmation are summarized in 
Table VII.C.2. The data are reported for three different periods during the mission. The first and longest 
period had the best pointing reconstruction and lasted from the beginning of the survey (SOP 29) through 
the end of the second hours-confirming coverage (SOP 425). The second period began with the start of 
the third coverage (SOP 426) when the survey strategy required large cross-scan slews. These resulted in 
larger thermal misalignments and limit-cycle bursts (see Section V.B) and the quality of the pointing 
reconstruction suffered in both the in-scan and cross-scan directions. A third period started at SOP 466 
when on-board attitude control was switched from the noisy z-axis gyro ZA to the quieter ZB gyro at 
SOP 466. The quality of the in-scan positions regained its earlier value, while the cross-scan accuracy, 
although improved, never returned to values seen earlier in the mission. 


VII- 10 


Table VII.C.2 IRAS-SAO Position Differences at Seconds-Confirmation 




(In-Scan) 




SOP Range 

No. 

Matches 

Bright Stars 
Mean 
<"> 

Population 

Sigma 

n 

No. 

Matches 

Faint Stars 
Mean 
(") 

Population 

Sigma 

(") 

29-425 

10332 

-0.9 

3.0 

59805 

-1.0 

4.0 

426-465 

783 

- 0.7 

5.9 

5114 

-0.6 

7.0 

466-600 

2446 

-0.2 

3.0 

15974 

-0.1 

4.0 

(Cross-Scan) 

SOP Range 

No. 

Matches 

Bright Stars 
Mean 

n 

Population 

Sigma 

(") 

No. 

Matches 

Faint Stars 
Mean 
(") 

Population 

Sigma 

('') 

29-425 

1806 

1.2 

7.1 

(4.1)* 

3342 

1.6 

7.3 

(4.5)* 

426-465 

138 

-2.0 

10.3 

(8.5)* 

306 

-2.4 

10.9 

(9.2)* 

466-600 

457 

1.3 

8.7 

(6.5)* 

1010 

0.9 

9.3 

(7.3)* 


* Sigma after removal of the 10" half-width uniform uncertainty associated with edge detections. 


D. Photometric Accuracy 

D.l Absolute Calibration Uncertainty Checks 

There are a number of checks which have been made to ensure that the absolute calibration, which 
was established with respect to the pointed observations, was carried consistently through the survey pro- 
cessing. Three specific checks include a comparison of the IRAS 12 and 25 pm observations of selected 
stars with the ground based observations of the same stars; a comparison of the quoted flux densities for 
the secondary standard NGC 6543 with those inserted into the processing; and a comparison of the flux 
density ratios between different wavelength bands with the relations assumed in Eq. (VI.C.2). 

Table VII.D. 1 lists the flux densities quoted in the catalog of a subset of bright stars observed by 
Rieke et al. (1984) and by Tokunaga (1984) at 10 and 20 pm. The selected stars showed no obvious 


VII- 11 




excess at the IRAS wavelengths, and thus the IRAS flux densities were color-corrected assuming their 
input energy distributions followed a 10,000 K blackbody. The flux densities are represented as magni- 
tudes following the formulation of Eq. (VI.C.l). As in Section VI.D, for the purposes of easier com- 
parison, the magnitudes of the ground-based observations were adjusted in their zero points to 
correspond to the IRAS magnitude and color conventions. It is seen that the agreement between the 
IRAS measurements and the ground-based measurements is excellent. 

The second check of the photometry, that of comparing the quoted catalog values of the flux of 
NGC 6543 with those which were inserted in the processing is important in order to verify that no sys- 
tematic errors were introduced in the complex processing procedure. In fact, the ratios of the quoted 
catalog flux densities to those entered in the processing are: 1.02 ±0.02, 1.04 ±0.01, 1.01 ±0.01, and 
0.97 ±0.01 in the four wavelength bands. The quoted flux densities were the result of 18 hours- 
confirmed measurements of the source and the uncertainties quoted above are in the mean value derived 
from these observations. Thus the processing did, in fact, introduce some systematic bias into the cali- 
bration procedure, but this bias is less than 5%. The origin of this bias is not understood. 

The final check on the photometric accuracy of the IRAS calibration is provided by Figs. VII.D.l, 
D.2 and D.3. In these figures, the 25, 60 and 100 pm flux densities are plotted versus the 12, 25 and 60 
pm flux densities for all point sources in the IRAS survey at Galactic latitudes 16 1 > 50°. In each of the 
plots, the locus followed by hot stars is apparent as the top envelope of the plots. At the longer 


■“ Table VII.D.l 

IRAS Survey 

versus Ground Based Magnitudes 




12 pm Band 



25 pm Band 

A R 

Star 

IRAS mag 1 

A? 2 

A* 3 

IRAS mag 

A t 

P And 

-2.11 

-0.09 

— 

-2.14 

- 0.10 

~ 

a Ari 

-0.70 

— 

+0.08 

-0.83 

-- 

-0.03 

a Tau 

-3.08 

-0.07 

-0.07 

-3.02 

+0.02 

+0.02 

a Aur 

-1.90 

+0.02 

+0.02 

-1.95 

-0.07 

-0.02 

a CMa 

-1.36 

+0.04 

— 

-1.38 

-0.07 


a CMi 

-0.71 

+0.03 

+0.01 

- 0.72 

-0.04 

-0.01 

P Gem 

-1.21 

+0.01 

+0.00 

-1.19 

-0.03 

+0.03 

a Hya 

-1.46 

— 

-0.04 

-1.37 

— 

+0.07 

\i UMa 

-0.98 

+0.03 

— 

-1.07 

-0.04 

— 

a Boo 

-3.22 

-0.07 

-0.07 

-3.09 

-0.01 

+0.07 

Y Dra 

-1.44 

— 

+0.04 

-1.50 

— 

+0.01 

y Aql 

-0.68 

+0.08 

+0.07 

-0.81 

-0.04 

-0.04 

Avg - 

— 

-0.00 

+0.00 

- 

-0.04 

+0.01 

Population 
Sigma - 

— 

0.06 

0.06 

- 

0.04 

0.04 


1 IRAS mag is the magnitude obtained from the survey using (Eq. VI.C.2). 

2 A t is the difference between the survey magnitudes and those of Tokunaga (1984). 

3 A R is the difference between the survey magnitudes and those of Rieke et al. (1984). 


VII- 12 



wavelengths the regions defining these stars becomes more spread out because of the presence of cooler 
envelopes around many of the stars. The color relations defined by Eq. (VI.C.2) and 

[100 pm] - [60 pm] - 0.06 mag (VTI.D.l) 

are included on the figures. Equation (VII.D.l) is based on the solar flux densities discussed in Section 
VI.C. The color relationships are seen to agree with the observed colors of the presumed hot stars, thus 
confirming the validity of the calibration and processing procedures at the 10% level. 

D.2 Relative Photometric Accuracy 

A global view of the quoted relative photometric errors of the catalog is shown in Table VTI.D.2 
including all sources of moderate or high quality flux densities (Section V.H.5). As with the analysis of 
the positional accuracy, the analysis of the photometric accuracy was done separately for bright sources, 
with flux densities greater than 1.2, 1.5, 1.9 and 4 Jy at 12, 25, 60 and 100 pm, and faint sources, with 
flux densities below these limits. In addition, the effects of excluding sources with a quoted probability of 
true variability greater than 50% at 12 and 25 pm (Section V.H.5) and sources in high source density 
regions (see Section V.H.6.a) are shown. In all cases, discrepant fluxes (i.e. sources with HCON-to- 
HCON flux ratios showing reduced X 2 greater than 9; see Section V.H.5) were excluded from the samples. 



Figure VII.D.l The 12 and 25 pm flux densities for cataloged sources are plotted. The line shows 
photosphenc models (see Section VI.C). 


VII- 13 



LOG f v (60 (l m) , Jy 


Figure V1I.D.2 The 25 and 60 flux densities for cataloged sources are plotted. The line shows 
photospheric models (see Section VI. C). 



LOG f„ (100 /jm), Jy 


Pi-ure VH.D.3 The 60 and 100 pm flux densities for cataloged sources are plotted. The line shows 
photospheric models (see Section VI. C). 



Table VII.D.2 

Point Source Catalog Relative Photometric Uncertainties 



Bright Sources 



Faint Sources 


wavelength 

Number 

Mean 1 

Mean 

Number 

Mean 

Mean 

(pm) 

Sources 

Unc. 

No. HCON 

Sources 

Unc. 

No. HCON 

(Moderate and High Quality 1 

Flux Densities) 

12 

61,653 

0.09 

3.2 

97,122 

0.10 

2.9 

25 

32,916 

0.11 

3.2 

67,652 

0.13 

3.2 

60 

23,335 

0.14 

3.1 

47,342 

0.11 

2.8 

100 

26,755 

0.15 

3.0 

43,562 

0.13 

2.7 

(High Source Density Regions Excluded) 

12 

39,029 

0.07 

3.2 

86,962 

0.10 

2.9 

25 

17,320 

0.09 

3.2 

56,682 

0.13 

3.2 

60 

9,437 

0.11 

3.0 

43,012 

0.11 

2.8 

100 

7,080 

0.11 

2.8 

34,593 

0.12 

2.6 

(High Source Density Regions and Probability of Variability >50% Excluded 

2 ) 

12 

30,228 

0.06 

3.1 

82,869 

0.10 

2.9 1 

25 

11,586 

0.07 

3.1 

48,609 

0.125 

3.2 

60 

8,055 

0.11 

3.0 

41,056 

0.11 

2.8 

100 

6,809 

: 

0.11 

2.7 

34,358 

0.12 

2.6 


1 Mean unc. is the quoted fractional uncertainty 

2 Sources deemed variable at 12 and 25 pm were not used at 60 and 100 pm 


Tests of the photometric accuracy were made by comparing the flux densities obtained for a given 
source on all possible pairs of HCONs. The ratio of two hours-confirmed flux densities belonging to a 
weeks-confirmed source is a random variable whose variance should be the sum of the variances in the 
two individual hours-confirmed flux densities. 

D.2.a Relative Photometric Accuracy of HCONs 

Figures VII.D.4 and VII.D.5 show histograms of the natural log of ratios of HCON flux densities 
for bright and faint weeks-confirmed sources. The samples were the bottom set described in Table 
Vn.D.2. The flux density chosen to be in the numerator of the ratio is the one with the higher flux status 
(see Table V.D.5) or the later time of observation in cases of equal flux status. 

The distributions in Fig. VII.D.4 and VII.D.5 were fitted to Gaussian distributions using only data 
above 25% of the peak value of the distribution. This choice fits the observations in the central region; 
the excess in the wings becomes quite visible. The significant non-Gaussian component is thought to 
come primarily from the effects of particle radiation, and is especially obvious for the brighter sources. In 
addition to particle radiation effects, there may also be some contribution due to intrinsic variability 
(which was significantly reduced in the samples discussed here), cross-scan variations in detector sensi- 
tivity particularly at 60 and 100 pm and the systematic selection of brighter fluxes at points in the pro- 
cessing where the measurements do not qualify for refinement (Section V.D). All of these effects are 
probably masked in fainter sources by their larger Gaussian noise. 


VII- 15 





NUMBER OF HCON PAIRS NUMBER OF HCON PAIRS 





NUMBER OF HCON PAIRS NUMBER OF HCON PAIRS 



LN (FLUX2/FLUX1) 




-0.8 -0.4 0 0.4 0.8 


LN (FLUX2/FLUX1) 


Figure V1I.D.5 Histograms of ratios of pairs of HCON measurements for bright sources in the four 
wavelength bands. 


VII- 17 






The exact nature of the particle radiation effects on the photometry is not well understood, but it is 
clear that some subtle phenomena are involved. Large spikes were suppressed by the deglitcher circuits, 
and were most likely to inhibit the detection of a real point source, although a detection with a large pho- 
tometric estimation error could result. Small spikes may easily have been lost in the other noise 
processes. Intermediate spikes tended to cause flux overestimation in cases which have been examined, 
but whether this led to setting the discrepant-flux flag depended on the actual brightness of the source 
relative to the spike as well as the uncertainties assigned to the fluxes. 

Table VII.D.3 lists the widths (la) of the Gaussian fits to the In (f v (2 )// v (1)) distributions shown 
in Fig. VII.D.4 and VII.D.5. If the variance of these distributions is twice the mean variance of a single 
HCON and if the catalog flux densities are generated from an average of three HCONs, the expected 
mean catalog relative photometric uncertainty would be 0.03, 0.03, 0.04 and 0.05 for bright sources in 
the 12, 25, 60 and 100 pm bands, respectively. This estimate of catalog flux density uncertainties is 
based on the intrinsic uncertainty of HCON pairs and is substantially smaller than the mean uncertain- 
ties listed in the catalog (Table VII. D. 2). For the 12, 60 and 100 pm bands the differences can be attri- 
buted to effects of the non-Gaussian wings, although in the 25 pm band the non-Gaussian wings are quite 
small. 

D.2.b The Quoted Flux Density Uncertainties 

The uncertainties in the quoted flux densities are based on a statistical evaluation of the consistency 
of HCON ratios. As discussed below, the IRAS survey produced highly repeatable flux densities and 
therefore small quoted uncertainties. Systematic effects on the photometry, such as flux-dependent 
nonlinearities, are discussed in Chapters IV and VI and almost certainly dominate the true photometric 
uncertainties, especially for the brighter sources and longer wavelengths. 


Table VII.D.3 Gaussian Fits to Distributions of Photometric Ratios 




Bright Sources 

Faint Sources 


wavelength 


Gaussian 1 


Gaussian 

(pm) 

Number 

a 

Number 


a 

ln(f v (2)/f v (\) ) 

12 

107,932 

0.074 

245,465 


0.15 

25 

44,551 

0.072 

145,069 


0.15 

60 

25,867 

0.096 

101,274 


0.15 

100 

17,234 

0.13 

72,003 


0.18 

/n(/ v (2)// v (l))/( HCON uncertainty) 2 

12 

107,932 

0.65 

245,465 


0.92 

25 

44,551 

0.51 

145,069 


0.71 

60 

25,867 

0.44 

101,274 


0.77 

100 

17,234 

0.60 

72,003 


0.87 


'The Gaussian a is derived from the fit to the upper 75% HCON pairs in Figs. VII.D.4, D.5, D.7, D.8 
2 The HCON uncertainty is the square root of the sum of the squares of the individual HCON uncertain- 
ties. 


The repeatability of the quoted flux densities was assessed by a comparison of the flux ratio of two 
hours-confirmed flux densities belonging to a weeks-confirmed source with the HCON uncertainty, 
defined as the square root of the sum of variances of the hours-confirmed flux densities. The ratio of two 
such hours-confirmed flux densities divided by the square root of the sum of the variances should be a 
random variable with unit mean and unit variance. 

Figure VILD.6 shows such a histogram for the entire HCON sample in the 12 pm band including 
the Gaussian curve which best fits all the observations in the histogram. It can be seen that the fit is 
somewhat crude. The wings of the distribution do not fall off rapidly enough to be Gaussian, and so the 
fit acquires too large a variance. In this figure, the standard deviation of the Gaussian curve is 1.07. 
The difference from unity is the residual error in the assignment of photometric uncertainties at seconds- 
confirmation and the error associated with the Gaussian assumption used in flux refinement at seconds- 
and hours-confirmation. 

Figures VII. D. 7 and VTI.D.8 show histograms of the HCON-to-HCON flux density ratios divided by 
the resultant HCON uncertainty defined above for the bright and faint sources in the selected sample. 
Included on each plot are the Gaussian fits using only data above 25% of the peak value, while Table 
VII.D.3 shows the widths (la) for the fits. The distributions typically show Gaussian variances 
significantly less than unity accompanied by substantial non-Gaussian wings. Together these two distri- 
butions produced the total photometric dispersion which led to the uncertainties quoted in the catalog. 



-4 -2 0 2 4 


LN (FLUX2/FLUX 1) /a 


Figure VII.D.6 Histograms of ratios of pairs of HCON measurements normalized to the resultant 
HCON uncertainty for bright 12 pm sources. A Gaussian fit to all of the data is 
shown. 


VII- 19 



NUMBER OF HCON PAIRS NUMBER OF HCON PAIRS 


— n r i 

12 Jim 

fy >1.2 Jy 

r 

i 

— i — i i 

‘ — ArfTTftllI--t 

' 1 

■1 


1 i iii 

60 

f i/ >1.9 Jy 

i 

— i — i — i i — 

s 

jfl 

i • — --rfifliiJtilli. 



-2 0 2 4 

LN (FLUX 2/FLUX 1)/<j 



Figure VII.D.7 


Histograms of ratios of pairs of HCON measurements normalized to the resultant 
HCON uncertainty for bright sources in the four wavelength bands. A Gaussian fit to 
the central portion of the data is shown. 


V1I-20 






NUMBER OF HCON PAIRS NUMBER OF HCON PAIRS 



-4-20 2 4 -4 -2 0 2 4 


LN (FLUX 2 /? LUX 1)/o [_N (FLUX 2/FLUX 1 )/a 


Figure VII.D.8 Histograms of ratios of pairs of HCON measurements normalized to the resultant 
HCON uncertainty for faint sources in the four wavelength bands. A Gaussian fit to 
the central portion of the data is shown. 


VTI-21 






The small variances of the fits imply that the quoted flux density uncertainties have been over-estimated, 
but this may well be an artifact of the selection criteria used to establish the sample. This suggestion is 
supported by the fact that during processing more reasonable x 2 test results were obtained. In the 12, 60 
and 100 pm bands the small variance may be due to fitting only the central portion of the distribution 
and excluding the broad wings. For the brightest objects, the small variance may be indicative of the true 
Gaussian dispersion which is superimposed on the broader distribution in the tails. 

The brightest and faintest 10% of the sources in the samples show the same effects as discussed 
above in an enhanced manner. In addition, the faintest sources show a statistically significant deviation 
in the mean of about 2% although the fitting process assumed a mean of zero. This would slightly inflate 
the derived standard deviation of the Gaussian fit. The brightest sources show a small non-zero mean in 
the opposite direction from that of the faint sources. These effects are not understood. 

D.3. Variable Sources 

The method by which a probability of variability as quoted in the catalog is calculated for a source 
with 12 and 25 pm fluxes is given in Section V.H.5. Figure VII.D.9 shows the distribution of the flux 
density variations between HCONs in units of the standard deviation for 25 versus 12 pm flux densities. 
Only catalog sources in unconfused areas of the sky are plotted. Sources in the shaded region (a) have a 
"probability of variability" between 0.5 and 0.99. Those sources in the outer regions (b) are all flagged 
with a probability of variability greater than 0.99. 

Table VII.D.4 gives the approximate relative change in the flux densities at 12 and 25 pm associ- 
ated with a given probability of variability. Two values are given, one for bright sources and one for faint 
sources. The table shows that there is a strong increase in the observed relative flux change as the quoted 
probability of variability increase from 90 to 99%. It is seen that variable sources whose fluxes change by 
20% or more at 12 and 25 pm can be identified reliably. A caution is necessary however. In regions of 
high source density, scan-to-scan variations in the baseline can affect the measurements enough to pro- 
duce spurious indications of variability. 


Probability of 
Variability 

Table VII.D.4 

Approximate Relative Flux Change at 1 2 and 25 pm 
Bright Sources 

Faint Sources 

0.50 

10% 

15% 

0.75 

15% 

20% 

0.90 

20% 

30% 

0.99 

30% 

40% 


iiimm mu Hill ill iiimin ii n !■■■■ n i m ilium n n iiiaiiin iiiwuhmi <h mu mu 



ORIGINAL PAGE IS 
OF POOR QUALITY 



-10 -8 -6 -4 -2 0 2 4 6 8 10 

LN (HCON1/HCON2) / a 12/im 


Figure VII.D.9 HCON-to-HCON flux variations at 12 and 25 pm, normalized with respect to the 
resultant HCON uncertainty. 


Figure VII.D. 10 shows the distribution in number of sources vs. probability of variability quoted in 
the catalog. Twenty-five percent of eligible sources have a probability of variability greater than 50%; 
13% have a probability of variability above 90%. Sources with a high probability of variability preferen- 
tially vary on a long time scale with roughly ten times as many sources varying over a time scale of 6 
months as on time scales of weeks. 


VTI-23 




Figure VTI.D.10 A histogram showing the number of sources with a given probability of variability. 


Figure VlI.Ap.16 shows the distribution on the sky of sources flagged with a probability of variabil- 
ity of 0.9 or greater. The plot clearly shows a population of sources occupying a region near the center of 
the Galaxy, probably due to a population of highly variable sources in the Galactic bulge. In the Galactic 
plane at Galactic longitude 320°, there is a diagonal edge showing a marked increase in the density of 
variable sources. This, as well as other more subtle structures, is probably a sampling effect since in this 
region the source density is high and the survey strategy produced extra HCONs at time intervals suitable 
for detecting variability. This same effect is also seen in the Large Magellanic Cloud, which because it 
was at the south ecliptic pole received many coverages. 

D.4 Discrepant Fluxes 

The distribution of sources with discrepant fluxes at 12 and 25 pm greatly resembles the distribu- 
tion of variable stars, suggesting that time variability is the dominant cause of such discrepant fluxes. At 

vn-24 


I II 




60 and 100 urn, the distribution is much different, since the sources cluster in two patches in the Galactic 
plane for reasons that are, at this time, unknown. Only about 1,000 sources have discrepant fluxes at 60 
and 100 pm. It is worth noting that the sources with discrepant fluxes are not predominantly found close 
to the ecliptic plane so that asteroids are not a significant cause of this problem (Section VTI.F). 

Roughly 50,000 sources are tagged with discrepant upper limits. These sources are mainly near the 
Galactic plane where the erratic behavior of the noise estimator causes this variation. 


E. Point Source Processing Considerations 
E. 1 The Nature of Rejected Sources 

Detections had to survive a number of stringent confirmation tests to become an HCON. The 
nature of those sources that became HCONs but did not go on to become cataloged objects is of interest 
to those trying to understand the completeness and reliability of the IRAS survey (see Chapter VIII). 

E.l.a Single HCONs 

Away from highly confused regions there were three main causes of single HCONs. Single HCONs 
can be due to inertially fixed sources faint enough to be below the completeness threshold, to moving 
sources such as asteroids and comets, and to wholly spurious sources generated by noise, radiation hits, 
diffraction spikes, and debris near the spacecraft. 

The number of true, fixed sources which appear only as single HCONs can be estimated from the 
completeness of the catalog in each flux range. A crude estimate based on the preliminary figures for 
completeness given in Section VIII. D suggests that half of all single HCON sources in the sky with two 
HCON coverages are real but incomplete; only one third of the single HCONs in the region covered with 
3 HCONs are predicted to be real. An estimate of the number of asteroids and comets is given in Sec- 
tion VII.F. 

Figures VTI. Ap. 1 7-20 show the distribution of single HCON sources detected in a given band in 
Galactic coordinates. These plots look much like the plots shown earlier for catalog sources, with the 
obvious addition of asteroids and comets in the ecliptic plane. 

E. 1 .b Rejected Weeks-Confirmed Sources 

Roughly 10,000 weeks-confirmed sources were rejected because they did not have consistent sight- 
ings of acceptable quality in at least one wavelength band. Most of these sources were caused by the ubi- 
quitous infrared cirrus. 

Regions with high source density were specially processed to generate sources that were relatively 
isolated, had repeatable fluxes and stood out prominently as point sources against the local background. 
High source density rules were applied to measurements in those wavelength bands for which the density 
of sources in a 1 sq. deg bin exceeded the confusion limit threshold. Sources were rejected from the cata- 
log only if the high source density processor rejected all the bands that it processed and if the bands that 
it did not process failed to meet the normal rules for inclusion in the catalog (Section V.H.5). 

The number of 1 sq. deg bins processed according to high source density rules was 690, 631, 1382 
and 6192 at 12, 25, 60 and 100 pm, respectively, or roughly one-seventh of the sky at the longest 


VII-25 


wavelength. In all, some 300,000 individual sources consisting of one or more hours-confirmed sightings 
were examined. Of these, approximately 170,000 were immediately rejected as having only one hours- 
confirmed sighting. Under no circumstances could any of these have reached the final catalog. About 
40,000 were rejected for failing to have two hours-confirmed sightings with at least two "perfect sightings 
(with FSTAT-7) in at least one band. 

Another 17,000 weeks-confirmed sources subsequently failed one or more of the tests described in 
Section V.H.6. Because these objects lacked high quality fluxes in any band, they were thus excluded 
from the catalog. 

Table VII.E. 1 gives the fraction of sources for which a particular reason was the cause of the rejec- 
tion of a measurement in a particular band. The tests were applied in the order listed in the table and 
represent an obstacle course that all bands being processed according to high source density rules had to 
survive. Since a band could be rejected at any point in the sequence, the reason that finally led to its 
rejection is counted in the table. 

It can be seen from the table that the dominant reasons for rejection vary with wavelength. In all 
cases the first test which demanded at least two sightings with a correlation coefficient greater than 0.97 
removed the bulk of the sources, particularly at 60 pm. The correlation coefficient test served to reject 
both low signal-to-noise sources and extended sources. The latter evidently dominate confused regions at 
60 pm. The effects of neighboring sources accounted for most of the rest of the rejections. 



Table VII.E.l Reasons for Rejection of a Band 

(Percent of Rejected Sources) 


Reason 

All Bands 

12 pm 

25 pm 

60 pm 

100 pm 

Correlation 

Coefficient 

87 

47 

57 

96 

75 

Confusion 

Status 

5 

6 

5 

2 

14 

Inconsistent 

Fluxes 

0 

0 

1 

0 

0 

Weaker 

Neighbor 

2 

18 

10 

0 

1 

Confused 

Neighbor 

6 

29 

26 

2 

10 

Very Near 
Neighbors 

0 

0 

1 

0 

0 


Figures VII.E. la-d show the effects of high source density processing on the number of sources 
within a 1 sq. deg bin in the four wavelength bands. In these histograms the open bars give the 
number of bins containing the specified number of sources, where in this case a source is counted in a 
band if it meets the criterion of having at least two hours-confirming sightings with FSTAT > 3. The 




sNia jo yjawnN SNia jo aaawriN 


VII-27 







confusion limit corresponding to 25 beams per source is marked. The solid bars show the results after 
the high source density criteria are applied. In this case a source was counted only if it had a high quality 
flux in a given band according to the high source density criteria. The figure shows a drastic reduction in 
the average source density after the more stringent criteria were applied. There are relatively few 1 sq. 
deg bins with more than the confusion limited number of sources. The striped bars indicate that the 
number of sources per sq. deg increases when high and moderate quality fluxes are included suggesting 
that confusion effects may be important for moderate quality fluxes in highly populated areas. 

E.2 Bright Source Problems 

The point-spread function of the telescope caused bright point sources to illuminate many more 
detectors in the focal plane than just those over which their image centers passed. These extremely bright 
sources caused special problems and received special handling during the data processing. During final 
source selection the WSDB was searched for spurious sources due to optical cross-talk from bright 
sources. As discussed below, 22 of these were deleted. 

Often, the extra detections of a bright source could not be combined into a single seconds- 
confirmed sighting and created a variety of problems: four detections from one source in one band were 
sometimes processed as two separate seconds-confirmations; sources with "too many" triple detections 
could fail to band-merge; leftover unconfirmed detections could take priority over the seconds-confirmed 
source to produce an incorrect measurement of the source in one or more bands, with enough extra 
detections both a primary and one or more false HCONs could be produced. The spurious sources 
could have either similar or much weaker fluxes than the primary source. 

Many of these problems were solved in the normal course of the data processing. Optical cross-talk 
processing eliminated weaker neighbors of bright sources (Section V ,D.2.c). Slot extensions for edge 

detections, priority for adjacent wavelength bands and priority for seconds-confirmed sources over non- 
seconds-confirmed objects produced more complete band-merging. Second-confirmed sources were given 
priority over non-seconds-confirmed sources in hours-confirmation. All of these changes had the effect of 
pulling more detections into a given confirmed source, leaving fewer odds and ends to confirm and to 
cause unreliable sources. 

One product of bright source processing not corrected in the basic processing was the optical cross- 
talk detection of the diffraction image of the secondary support spider. The spider produced cross-talk 
emission in six arms in the focal plane, two in the in-scan direction and four oriented at angles of ± 60° 
and ± 120° with respect to the in-scan direction. Figure VH.E.2 shows an example of optical cross-talk. 
The source is IRC+10216 observed directly at 60 pm on only two detectors, 14 and 33. All the detectors 
show some evidence of its passage. Detections on 14 and 33 were well over the signal-to-noise threshold 
of 300, so that optical cross-talk processing suppressed the detections on 9 and 37. The double detections 
on detectors 10, 13, et al. are characteristic of the "spider-arm" diffraction and were not suppressed by 
cross-talk processing because they were outside the in-scan search window (Section V.D.2.c). Many 
potential detections did not survive because of poor correlation coefficients or poor in-scan alignment 
with potential partners. However, the detections marked on detectors 8 and 35 not only seconds- 
confirmed but went on to hours- and weeks-confirm as well. 



ORIGINAL PAGE is 

OF POOR QUALITY 



Figure VILE.2 Strip-chart tracings of the detector outputs at 60 pm during a passage of 

IRC+10216 over the focal plane. The detector timing has been adjusted so that 
the detector samples correspond to the same in-scan positions. 


Figure VII. E. 3 shows the area around IRC+10216 in the WSDB. This figure shows the optical 
cross-talk problem at its worst. IRC+10216 generated spurious HCONs in three bands. Spurious sources 
line up in the typical spider arm pattern ±60 and ±120° from the in-scan direction. Because this source 
lies in the ecliptic plane, the scan angle was always along the local meridian so that there was no HCON- 
to-HCON variation in twist angle to prevent weeks-confirmation of the cross-talk sources. For sources at 
higher latitudes, however, the scan direction usually varied enough to prevent the weeks-confirmation of 
spurious HCONs. Eleven spurious weeks-confirmed sources surround IRC+10216, of which seven were 
good enough to pass initial catalog selection rules. These seven were deleted by hand from the final cata- 
log, along with 15 other such sources. 

The final catalog data base was searched for bright sources and their neighbors above Galactic lati- 
tude 5°. Catalog sources brighter than 450 Jy at 25 pm or brighter than 1000 Jy at 12, 60, or 100 pm 
were examined. These thresholds were selected after analysis of several hundred bright sources and are 
the levels at which cross-talk confirmations begin to survive final catalog screening. Any bright source 
neighbor, i.e., having a position within a 1000 window and a high quality flux that was band-compatible 


VII-29 




ECLIPTIC LONGITUDE ECLIPTIC LONGITUDE 

Figure VII.E. 3 The vicinity of IRC+10216 is shown before and after the weeks- confirmation 

requirement was imposed on sources in the Working Survey Data Base. Spurious 
sources due to cross-talk are shown as triangles. Those that survived weeks- 
confirmation were subsequently deleted. 

with the bright source, was identified and evaluated for evidence of optical cross-talk. Analysis of the 
raw data in strip chart form, e.g., Fig. VII.E.2, gave the surest evidence of cross-talk. Neighbors judged to 
be independent sources had detections that were free from influence of the nearby bright source: the 
detector plots showed local minima separating the two sources. Cross-talk neighbors, on the other hand, 
were not separated by a local minimum but were found on a plateau of emission that was a function of 
the intensity, direction, and distance from the central source. 

In addition to being found next to an extremely bright source, cross-talk neighbors also tended to 
have the following characteristics: detection in only a single band; only 2 HCONs even if more were 
expected; a location in one of the preferred spider-arm directions; and a brightness proportional to its dis- 
tance from the primary object. 

Twenty-two neighbors survived final catalog screening but were found to be due to cross-talk. 
These are summarized in Table VII.E.2. Notice that only the cross-talk source near a Sco was confirmed 
in two wavelength bands. All others are single band sources. Galactic latitudes within 5° of the plane 
were not examined. In addition, the Orion region (especially around Mon R2, OMC 1, NGC 2024, and 
NGC 2071) had too many neighbors to examine. In such crowded regions the high source density pro- 
cessor suppressed most remaining cross-talk sources, but users should be aware that significant numbers 
of cross-talk sources may remain. 


VII-30 





Table VII. E.2 

Parent 

Name Source 

Bright Source Neighbors Suppressed as Cross-talk 

Parent Flux Problem Flux Ratio Distance 

Density (Jy) Band (urn) (X-talk/Parent) (") 

Total 

HCONs 

04361-6208 

R Dor 

5549 

12 

CO 

X 

o 

1 

4* 

141 

2 

04399+3604 

AFGL618 

1107 

60 

9X10 -4 

345 

2 

04400+3559 

n 

• 

60 

9xl0“ 4 

353 

2 

07204-2542 

VYCMa 

6651 

25 

vn 

1 

© 

X 

427 

3 

09501+6956 

M 82 

1145 

100 

2xl0~ 3 

494 

2 

09509+7000 

II 

ii 

100 

4xl0~ 3 

411 

2 

09508+6955 

n 

1168 

60 

2xl0 -3 

283 

2 

09514+6958 

ii 

ii 

60 

3xl0- 3 

206 

2 

09443+1328 

IRC+10216 

5652 

60 

1 x 10 -4 

780 

2 

09446+1340 

» 

23069 

25 

3xl0 -5 

750 

2 

09446+1329 

n 

n 

25 

5xl0 -5 

550 

2 

09446+1328 

ii 

47525 

12 

3xl0~ 5 

531 

2 

09448+1336 

ii 

23069 

25 

6xl0~ 5 

508 

2 

09455+1327 

it 

ii 

25 

6X10- 4 

303 

2 

09457+1332 

ii 

ii 

25 

lxlO -4 

477 

2 

09461+1332 

n 

ii 

25 

2xl0“ 5 

812 

2 

09458+1320 

ii 

rt 

25 

U> 

X 

o 

1 

u* 

834 

2 

09431-2147 

IRC-20197 

495 

25 

2xl0~ 3 

147 

2 

10494-2101 

V Hya 

459 

25 

4x 10 -3 

253 

2 

13271-2301 

R Hya 

585 

25 

2xl0~ 3 

150 

2 

13271-2303 

ii 

n 

25 

2xl0“ 3 

158 

2 

16261-2617 

a Sco 

3198 

12 

2X10" 4 

138 

2 



690 

25 

3xl0~ 3 




E.3 Sources of Incompleteness 

The causes of lost HCONs are various and include: (i) missing detections due to radiation hits or 
noise spikes (which cause correlation coefficients to fall below threshold); (ii) band-merge failure caused 
by noise detection in another band; and (iii) missing detections because the detections fell below the 
signal-to-noise ratio correlation coefficient thresholds. Causes (i) and (ii) usually are significant only 
because a failed detector also removed a detection as well. There exist two additional causes of incom- 
pleteness at 100 pm. First, the instability of the noise estimation caused by the presence of cirrus some- 
times causes the noise to be erroneously estimated high by a factor of two. Second, cirrus itself often 
creates confusion and cross-scan position shifts, resulting in a failure to hours-confirm. Finally, as dis- 
cussed in more detail in Section VII.F, asteroids can also cause lost HCONs. One bright source lost the 
HCON because of a coincidence with asteroid Valentine. 


VII-31 



1 'I'll If 


E,4 Effects of Failed Detectors 

The data processing allowed for failed or degraded detectors by giving the status of non-seconds- 
confirmed due to a failed detector (NSCF) to any detection whose failure to seconds-confirm might have 
been caused by the potentially confirming detector being failed or excessively noisy. NSCFs enjoyed the 
same status as seconds-confirmed band-merged detections and had the effect of significantly increasing 
completeness at the cost of hurting reliability. 

Because detectors 17 and 20 in the 25 pm band were dead, their seconds-confirming partners 40, 
44, and 41 often produced NSCF detections as did detectors 9 and 13, opposite the dead detector 36 at 
60 pm. Additional detectors declared dead due to their degraded performance were: 28, 25, and 26 at 12 
pm and 42 at 25 pm. The designation of so many "dead" detectors produced a flood of NSCF detections 
with NSCFs outnumbering seconds-confirmations by about 5 to 1. The purpose of declaring active, but 
degraded, detectors as failed was to avoid penalizing the more sensitive, confirming detector. Bright 
sources could still confirm on both detectors, since seconds-confirmation was attempted for detections 
produced by the degraded detectors. Weaker sources had a chance of confirming in the nominal way 
on subsequent sightings with a more favorable combination of detectors. 

An unfortunate side effect of allowing NSCFs was that weak radiation hits which passed the correla- 
tion coefficient test could masquerade as valid detections. When this occurred on the twelve detectors 
subject to NSCF status the radiation hit could band-merge and provide an erroneous flux or position for 
a true source. The requirement of weeks-confirmation in each wavelength band prevented this from 
becoming a serious source of error for the catalog. 

Seconds-confirmation was still possible despite passage over a failed detector. Consider a bright 
source passing over detectors 20 (failed), and 44 and 40 (Fig. II.C.6). Although detections from 44 and 
40 could seconds-confirm, they generally did not because any displacement from the overlap region 
resulted in a failure of the flux test. More than half the time the weaker detection band-merged into the 
main source while the stronger detection became a separate source. The weaker detection was more suc- 
cessful at band merging because it had a greater in-scan uncertainty than the stronger detection. Thus, at 
hours-confirmation one could often find two versions of the source, e.g., one with a good 25 pm flux and 
nothing else, the other with good fluxes at 12, 60, and 100 pm but a low flux at 25 pm. This problem 
was corrected in the reprocessing of SOPs 29-446 by suppressing the weaker detection in these nearly 
overlapping cases. An error in the computation of the cross-scan uncertainty prevented band merging of 
the stronger detection in about one-third of these special cases. 

E.5 Setting the Seconds-Confirmation Threshold 

When a confirmation threshold is set optimally, practically all true matches should be accepted, and 
practically all false matches should be rejected. One way to obtain a feeling for whether an acceptable 
threshold setting is possible and has been obtained is to process the same data with different thresholds. 
If one begins with a very high value, then lowering the threshold should result in a significant increase in 
the number of confirmations. Continuing to lower the threshold further should eventually yield little 
increase in the number of confirmations. This should happen as the completeness approaches unity. If 
the threshold is lowered still more, the number of confirmations will begin to rise again at some point as 
more spurious sources are accepted, unless there is no noise in the process at all. 


1 III 


VII-32 



If a reasonable range of threshold values is found in which the number of confirmations is essen- 
tially constant, then that range is assumed to contain the optimal value. Such a range should occur at 
thresholds which seem consistent with the properties of the measurement error. Assuming that gross 
blunders in estimating the measurement errors have not been made, then if no such range is found, even 
for threshold values which are clearly unreasonably large, then the noise contamination has probably 
blended significantly into the signal well before the onset of completeness for true matches. 

The position test thresholds in the seconds-confirmation processing were varied as described above. 
The optimal range was found at 12, 60, and 100 pm. The search was not as successful in the 25 pm 
band. While the number of additional confirmations rolled off as the threshold was raised, the total 
number never reached a plateau such as those of the other bands. Time constraints prohibited a 
thorough study of this anomaly, and an attempt to identify the spatial distribution of the excess events 
was inconclusive. The flux distribution appeared to be more strongly concentrated to the fainter end 
than the overall distribution of objects detected. 

A cursory analysis of the rate at which these sightings survived hours-confirmation indicated that 
only very few succeeded. It is likely that none at all passed through weeks-confirmation, so that no 
impact on the completeness or reliability of the catalog is expected. 


F. Asteroids and Comets 

F. 1 Number Present in Catalog (Aster oid Source Density) 

The WSDB contained 326 hours-confirmed sources associated with known asteroids and comets; 
among these, 108 sources have two or more hours-confirmations. Close examination of the sources with 
two or more HCONs reveals that they are all chance positional coincidences with inertially fixed objects. 
There is no case of any known asteroid or comet passing the weeks- or months-confirmation tests by 
self-confirmation. 

To determine the impact of these chance encounters on the reliability of catalog sources and their 
stated fluxes, the selection criteria and flux averaging methods discussed in Section V.H. 1 must be 
applied. Seventy-four sources associated with known asteroids meet all the selection criteria and appear 
in the IRAS catalog. In the case of 17 of these IRAS failed to detect the asteroid. Of those detected, 13 
had no effect on the cataloged quantities. The remaining 44 sources all have contaminated fluxes. These 
sources are identified in Table VII.F.l and have the following characteristics: 

(1) Two-thirds of these sources had at least one high quality flux contaminated, although in no 
case did an asteroid association change the value of the flux quality flag. 

(2) In the absence of the asteroid, 80% of these sources would have been detected at only a single 
wavelength, with equal probability for each of the four bands. Of the remainder, 10% had 
both 12 and 25 pm fluxes and 10% had both 60 and 100 pm fluxes. 

(3) The flux discrepancy flag is only a weak indicator of the impact of the asteroid. All the 
correct flags are set in only 15% of the cases, but in 58% of the cases at least one flag was set 
correctly. In 37% of the cases no flag was set. 


VII-33 


(4) Seventy-five percent of the affected flux densities are less than 1 Jy. The frequency of flux 
contamination is highest in the 25 pm band, however, and is a function of HCON as shown 
in Table VII.F.2. 

(5) The ecliptic latitude distribution of the contaminated sources follows the density distribution 
of the numbered asteroids. All the contaminated sources lie within 25° of the ecliptic with 
90% of them within 15°. 

The true population of asteroids in the WSDB is underestimated by the numbered asteroids. A plot 
of log N vs. log / v , where N is the number of known asteroids detected with 25 pm flux density equal to 
or greater than / v , shows a steep rise (N — kf^ 2 ) to N— 100, a flatter region (N — kj \, 1 ) from N— 100 
to 1000, and a gradual roll-off to N-3266. If the gradual roll-off is due to incompleteness of the num- 
bered asteroids, extrapolation of the flat portion of the curve to the 25 pm flux limit estimates that the 
true population of single HCON asteroids in the WSDB is about 10,000. Since the frequency of chance 
encounters is proportional to the number of asteroids available, these numbers suggest that about 135 
catalog sources could be affected. In general these may be expected to have characteristics similar to 
those discussed above for the numbered asteroids. 

If the extrapolation of the asteroid population is valid, then there should be about 0.2x135 “ 27 
sources within 25° of the ecliptic plane observed only at 25 and 100 pm, resulting from the detection of 
infrared cirrus plus a faint asteroid. Of the 500 IRAS sources with measurements only at 25 and 100 
pm, 34 of these clearly qualify as asteroid-cirrus combinations. The IRAS names of these objects are 


Table YII.F.l IRAS Names of Sources Contaminated by Numbered Asteroids 


-NAME-- 

—NAME— 

—NAME— 

—NAME— 

—NAME— 

00260-1016 

02449+0249 

03341+1224 

03301+1820 

04369+0449* 

05536+1944 

06000+1644 

07093+2522 

07153+1128 

07294+2521 

08140+3940 

07575+1756 

08070+2503 

08090+1239 

08534+0850 

09285+0847 

10253+1101 

15113-1310 

15170-2739 

15301-3259 

15427-2604 

16137-2047 

16475-0930 

16378-3133 

16528-0808 

16515-1634 

17013-2451 

17082-2903 

17239-0259 

17311-2013 

17424-2331 

17451-2102 

17457-3623 

17486-1701 

17586-1724 

18236-2320 

18317-1646 

18477-0940 

18559-2045 

18478+0215 

19171-2047 

19420-1304 

19509-1925 

21303-1858 

23586-0116 i 

^Source deleted by high source density processor. 


Table VILF.2 

Frequency of Flux Contamination 




Frequency (%) 



# HCONs 

12 pm 

25 pm 

60 pm 

100 pm 

2 

8 

97 

81 

16 

3 

60 

53 

27 

7 

4 

25 

50 

25 

0 




given in Table VII.F.3. Considering the small size of the sample, the agreement between the observed 
and predicted numbers is satisfactory and gives confidence in the estimated influence of asteroids. 

While non-inertial sources of infrared radiation have no place in a catalog of fixed point sources, 
they are of great interest to studies of the Solar system. As a part of the normal data processing all detec- 
tions of sources with colors appropriate to solar system objects (30K < T < 400K) were written to off- 
line data files for further processing which attempts associations with a larger set of known asteroids. 


Table VII.F.3 Cirrus Sources Possibly Contaminated by Asteroids 

—NAME-- 

—NAME— 

—NAME— 

—NAME— 

—NAME— 

02428+0748 

02313+2729 

03105+0744 

03100+1049 

03408+0101 

03085+3945 

03484+1744 

04070+1128 

04063+1220 

04012+2044 

04055+1859 

04438+1643 

04376+2922 

04362+3347 

05063+2015 

05341+3420 

05369+2350 

05533+3147 

07113+1225 

08230+2535 

14200-1907 

15049-2606 

15246-2503 

15455-2931 

16447-1644 

16478-3129 

17052-2835 

17343-0358 

17397-1317 

17538-0740 

18285-2217 

18371-3148 

19160-1606 

21531-1009 



G. Associations 

As described in Section V.H.9, the associations of cataloged sources with IRAS sources were done 
purely on a positional basis, except that no association was allowed between a star and an IRAS source 
detected only at 100 pm. For the catalogs with small position errors, the search box was 45" x 8" (half- 
width), and the probability of a chance association with an IRAS source was less than 0.1%. For those 
catalogs where the search radius was significantly larger, e.g., galaxy catalogs and catalogs of diffuse 
objects, the probability of random associations was significantly greater. For the majority of the 
galaxy catalogs, where the search radius was 90", the probability of chance associations with IRAS 
sources was about 0.3% at high galactic latitude ( I ft I > 45°) and increased with increasing IRAS source 
surface density. A particular concern is the association of galaxies with infrared cirrus. The far infrared 
properties of an external galaxy can be very similar to those of cirrus. The user is urged to examine the 
cirrus flags before attaching significance to such associations. 

It must be emphasized that the associations are not identifications. To be certain of any specific 
identification of an IRAS source with a cataloged source, the user must investigate the IRAS source in 
some detail. 

About 28% of IRAS sources were associated with cataloged objects. Approximately 46,000 IRAS 
sources were associated with cataloged stars and approximately 10,000 IRAS sources were associated with 
cataloged galaxies. Other classes of objects accounted for 15,000 associations. The vast majority of these 
were galactic nebulae and galaxies in the ESO (B) atlas. A measure of the fraction of spurious associa- 
tions can be found in the number of IRAS sources associated with both stellar and galaxy-type objects. 
At high galactic latitudes ~ 0.6% of the IRAS sources had both galaxy and stellar associations. 


VII-35 


H. Meaning of Point Source Flags 
H.l Confusion Flags 

Confusion flags are essential in assessing the quality of a point source. Sources with all flags equal 
to zero can be accepted without question. Sources with non-zero flags may have some problems as dis- 
cussed below. 

The confusion flags include a count of neighbors around a given source that were found in either 
the point or small extended source WSDBs, as well as a flag which was set by the processors during 
seconds-confirmation, band-merging and hours-confirmations. This relatively large number of flags was 
required by the complexities of the processing and of the infrared sky. 

H. 1 .a Point Source Neighbors 

Point source neighbor flags warn of problems in the immediate vicinity of the source. The point 
source neighbors come from the cleaned-up WSDB. The number of neighbors is counted in a rectangle 
oriented along the largest position uncertainty direction. That rectangle usually reflects the original scan 
direction used to observe that source. The large dimension of the box is a half-width of 6 along the 
smallest position uncertainty axis, which corresponds to the maximum distance in which another source 
may have shadowed or stolen any detections of the source in question (Section V.C.7). The half-width is 
4.5' along the orthogonal axis, corresponding to a full detector width. Neighboring sources within this 
distance produced at least one detection that was the sum of the two sources. 

Weeks-confirmed and single hours-confirmed point source neighbors are counted separately in that 
box. Non-zero values for either count, but especially for PNEARW, the count of weeks-confirmed 
sources, should caution the user that one or more bands of the source may be reported incorrectly and to 
look in a small region around the source for confusing neighbors. Most sources with neighbors are found 
in regions of high source density. 

H. 1 -b Small Extended Source Neighbors 

Two types of small extended source counts are given. The SES1 flag counts separately in each band 
the total number of hours-confirmed small extended sources. Note that a weeks-confirmed small 
extended source containing three hours-confirmations would contribute three to this count, and a small 
extended source containing only one hours-confirmation would contribute one to this count. The SES2 
flag counts separately in each band only the number of weeks-confirmed small extended sources. Both 
flags use the box described above for point source neighbor counting. 

Both flags warn of the presence of structures larger than point sources, and hence cast doubt on 
whether the reported source is truly point-like or is only part of a larger complex. Two separate flags 
exist for the following reason. In confused regions containing only a single small extended source, SES1 
would equal the number of hours-confirmations times SES2, and thus the flags would be redundant. 
However, in complex regions SES1 may be unrelated to SES2. That is, because of the effects of cluster 
analysis (Section V.E) there may have been no acceptable small extended source, even though many 
small extended source detections occurred. The distribution of sources with high values of SES 1 and 
with non-zero values of SES2 is shown in Figs. VII.Ap.21 and Ap.22. 


VII-36 



H. 1 .c Confuse Flag 


A final flag related to near neighbor confusion is derived from the repeated occurrence of confusion 
during the processing of a source. In a given band, if more than one candidate per scan could confirm 
with the source at either seconds-confirmation (excluding edge detections), band-merging, or hours- 
confirmation, the processor set a flag. If any such flag was set in two separate hours-confirmations of a 
source, then the confuse flag was also set in that band. Thus if the confuse flag is set, something other 
than an isolated point source was present at or near the position of the source in question, and the user 
should again be cautious. The distribution of sources with the confuse flag set correlates well with high 
source density areas. 

H.l.d Summary 

The user examining an IRAS source can immediately rule out almost any problem if all of its 
neighbor flags are zero, and SES1 is 0 or 1. Sources that have all those flags set to zero are almost always 
completely clean point sources. There are about 63,000 such sources in the entire catalog, 25% of all 
sources. Above a Galactic latitude of 20°, the fraction is much higher, 30,000 out of 60,000 sources. 
Including sources with SES1 less than 2, which are still almost always clean sources, increases the percen- 
tage of clean sources to 38% in the entire sky, and to 70% in the sky above Galactic latitude 20°. 

Sources with non-zero flags may be confused with point source neighbors or may only be part of an 
extended complex, and unfortunately must be examined in more detail. 

H.2 Cirrus Flags 

Cirrus flags are set to warn the user that a catalog source may be adversely affected by infrared 
cirrus or may even be no more than a bump in the infrared cirrus. These flags and the 2' sky brightness 
images are the best possible substitutes for examining the raw detector data. If all cirrus flags are smaller 
than the limits discussed below, then one can infer that the source is unlikely to have been affected by 
cirrus. 

CIRRI 

CIRRI gives an estimate of the existence of cirrus on the point source scale. CIRRI is total 
number of sources detected only at 100 pm, including sources with single and multiple hours- 
confirmations, in a 1 sq. deg box centered on the source. The sky distribution of sources with CIRRI 
greater than three is displayed in Fig. VII.Ap.23, and shows the areas where cirrus is a problem. This 
figure is similar to the map of sources detected only at 100 pm shown earlier (Fig. Vn.Ap. 10). 

CIRR2 

CIRR2 is a logarithmic function of the ratio of the 100 pm point source flux to an estimate of the 
100 pm flux produced by cirrus, derived from the filtered 0.5° 100 pm data (Section V.H.4). Figure 
VII.Ap.24 shows that the distribution of sources with CIRR2 greater than 4 correlates well with other 
cirrus maps. 

CIRR3 

Another estimate of the importance of cirrus comes directly from the total 100 pm emission is a 
0.5° beam data. CIRR3 is equal to the intensity of the 100 pm emission in MJy sr" 1 . 


VII-37 



SES1 


The SES1 flag at 100 pm, discussed earlier, is also a good cirrus indicator. See the earlier discus- 
sion. 

Summary 

No clear-cut prescription can be given to a user to guarantee that cirrus has not adversely affected a 
given source. If the source has a strong 100 pm flux, no missing bands, and all cirrus flags are well 
below the cutoff values discussed above, then it is likely that few problems exist. If any of the above con- 
ditions are violated, a closer examination of nearby sources will usually indicate the importance of cirrus. 
However, in some cases, only an examination of images or, finally, the detector data will resolve all 
doubts. 

I. The Small Extended Source Catalog 

An analysis of the properties of this catalog will be presented in the final version of this Supplement 
after the catalog has been released in 1985. 

J. Extended Source Products 

J, 1 Zodiacal Emission Effects 

Since the detector signals were DC coupled, the extended source data products include emission 
from sources on all angular scales. In particular, emission from interplanetary dust, or zodiacal emission 
was a prominent large scale signal component in all survey bands. The contribution of the zodiacal 
emission to the observed intensity toward any direction on the celestial sphere depended upon the 
integrated emission from the interplanetary dust along the line of sight at the time of observation, which 
in turn depended upon the Earth’s orbital position within the dust cloud (time of year) and the spatial 
structure of the cloud. Hence, repeated measurements in a given celestial direction at substantially 
different times of year gave different results. For example, for elongations near 90°, the 1 ° daily motion 
of the Sun produced about a 2% change in the sky brightness near the ecliptic plane in all survey bands, 
with the brightness decreasing with increasing elongation. Because the dust symmetry plane near 1 AU 
from the Sun was inclined with respect to the ecliptic plane, the brightness at high ecliptic latitudes varied 
sinusoidally annually with a peak-to-peak variation of about 20% in the bands. 

For the purposes of preparing the extended source data products, the hours-confirming scans (typi- 
cally separated by about two hours) were treated as if they had been obtained simultaneously and the 
results were averaged together. The difference in zodiacal brightness between one orbit and the next was 
less than 0.1% of the ecliptic plane brightness; within the maximum 36-hour spacing of an hours- 
confirming coverage, the brightness difference was less than 3%. The data from the three hours- 
confirming surveys were reduced separately and presented as three distinct sets of images. The data from 
the first two surveys, which each covered about 95% of the sky, were interleaved during the first 7 months 
of the survey, with the two weeks-confirming surveys typically separated by about 10 days at any given 
celestial position. The third survey, which covered 72% of the sky, was obtained during the last four 
months. 


VII-38 



As an aid to modeling and extracting the zodiacal emission contribution in any of the extended 
source products, the time-ordered Zodiacal Observation History File (ZOHF) was created. This file con- 
tains, at 0.5° sample spacing, the time (UTC), celestial coordinates. Sun-referenced observing angles, and 
measured brightnesses in the survey bands for all observations included in the extended source image 
products. 

J.2 Effective Resolution 

As discussed above, the data for the sky plates and Galactic plane maps were smoothed in the time 
domain to a sample spacing corresponding to 2'. Additional smoothing was inherent in the process of 
projecting these data into image grids. It was estimated that the resulting effective resolution, or, ability 
to distinguish point sources in these images is about three pixels, or 6'. 

J.3 Tests of Extended Source Calibration Consistency 

The primary test of the extended source calibration consistency was an examination of the meas- 
ured brightness in the survey data to show that the absolute baseline was being properly controlled 
through the daily observations of the baseline photometric reference area (the TFPR) as described in Sec- 
tion VI.B.3. 

One check consisted of observing the time variation of the brightness of the TFPR in the Zodiacal 
Observation History File (ZOHF). The mean of the ZOHF data followed the model of the TFPR given 
in Table VLB. 1 to within 5% in all bands except for a short period of time about 20 days before the end 
of the mission where deviations approaching 10% brighter than the model occurred. The scatter of the 
measurements was non-Gaussian, most of the points scattered toward higher brightness, with 50% of the 
points lying within 7%, 5%, 12% and 12% of the model in the 12, 25, 60 and 100 pm, respectively. 

As a second check the sum of the north and south ecliptic pole brightness, as measured from 
selected single scans of the telescope, was compared with parameter B 0 of the TFPR model in Table 
VI.B.l. The geometrical basis of the TFPR brightness model predicts that the sum of the north and 
south brightness will be constant and equal to twice B o. This was found to be true to within 5% in all 
bands. 


Authors: 

C. Beichman, T. Chester, R. Benson, T.Conrow, J. Fowler, T.N. Gautier, F. Gillett, 

G. Helou, J. Houck, H. McCallon, G. Neugebauer, B.T. Soifer, and R. Walker. 

References: 

Dressel, L.L. and Condon, J.J. 1976, Ap.J. [Suppl.], 31, 187 
Rieke, G., Lebofsky, M. and Low, F.J. 1984, preprint. 

Smithsonian Astrophysical Observatory Star Catalog 1966, Smithsonian Institution, Washington, 

D. C. 


Tokunaga, A., 1984, A.J., 89, 172. 




ORIGINAL PAGE IS 
OF POOR QUALITY 


II 


VII-40 


Figure VII.Ap.l The distribution of sources detected at 12 |im only (spectral combination "1000") 
shown in an Aitoff equal area projection in Galactic coordinates. 


ORIGINAL PAGE IS 
OF POOR QUALITY 



1 


Figure VTI.Ap.2 The distribution of sources detected at 12 and 25 pm (spectral combination "1100") 
shown in an Aitoflf equal area projection in Galactic coordinates. 



ORIGINAL PAGE IS 
OF POOR QUALITY 


I I! 


Figure VII. Ap. 3 The distribution of sources detected at 12, 25 and 60 (im (spectral combination 

"1110") shown in an Aitoff equal area projection in Galactic coordinates. 




Figure VII.Ap.4 The distribution of sources detected at 12, 25, 60 and 100 pm (spectral combination 
"1111") shown in an Aitoff equal area projection in Galactic coordinates. 



Figure VH.Ap.5 The distribution of sources detected at 25 pm only (spectral combination "0100") 
shown in an Aitoff equal area projection in Galactic coordinates. 








Figure VII.Ap.9 The distribution of sources detected at 60 and 100 |xm (spectral combination "0011") 
shown in an Aitoff equal area projection in Galactic coordinates. 




1C c< 




.ORIGINAL PAGE IS 
QE jpoor Quality 

VII-50 


1 Tli 





Figure VII.Ap.13 The distribution of sources detected at 12, 25 and 100 jun (spectral combination 
"1101") shown in an Aitoff equal area projection in Galactic coordinates. 


1011 



VII-53 


Figure VTI.Ap.14 The distribution of sources detected at 12, 60 and 100 pm (spectral combination 
"1011") shown in an Aitoff equal area projection in Galactic coordinates. 



VII- 54 


Figure VII.Ap.15 The distribution of sources detected at 25 and 100 pm (spectral combination "0101") 
shown in an Aitoff equal area projection in Galactic coordinates. 




ORIGINAL PAGE IS 
OR POOR QUALITY 



point sources that had only a single hours-confirmed sight- 
in an equal area Aitoff projection in Galactic coordinates. 




Figure VII.Ap.20 The distribution of 100 pm point sources that had only a single hours-confirmed 
sighting (an HCON) is shown in an equal area Aitoff projection in Galactic coordi- 





>f sources that were tagged as having one or more nearby, weeks- 
ttended sources (SES2 > 0) is shown in an equal area Aitoff pro- 
coordinates. 



ORIGINAL PAGE IS 

OE foor quality 










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projection in Galactic coordinates. 






1 1 III 


VIII. SKY COVERAGE, CONFUSION, COMPLETENESS AND RELIABILITY 

A. Introduction 

The main results concerning the sky coverage, completeness and reliability of the IRAS survey can 
be summarized by noting that: 1) 96% of the celestial sphere was covered sufficiently to appear in the 
catalog; 2) at Galactic latitudes lbl>20° and at wavelengths of 12, 25 and 60 pm the catalog is essentially 
complete for sources brighter than about 1.5 Jy; and 3) fewer than one source in 1000 having two hours- 
confirmed sightings is a spurious object. 

At 100 pm the situation is more complicated due to the effects of "cirrus"-like emission (cf. Low et 
al. 1984) that often affects point source fluxes and completeness and reliability even at high Galactic lati- 
tudes. Whether a particular point source is an independent entity such as a galaxy or a dense condensa- 
tion in a cirrus wisp has to be carefully considered. It is advisable to check the values of the cirrus and 
confusion flags for 100 pm flux densities because these may indicate extended emission at the point 
source position (see Sections V.H.4, VII.H.2 and X.B.l). 

In regions of high point source density such as the inner parts of the Galactic plane and the Magel- 
lanic Clouds, the completeness of the catalog is likely to be poor. More stringent criteria for the selection 
of catalog sources were applied in these regions to ensure a high reliability of the sources. It is hard to 
arrive at good estimates of the completeness and reliability, particularly since these are also regions of 
strong extended emission. The completeness is very low in regions affected by the Galactic plane "sha- 
dow" caused by the lagging of the noise estimator. 

The catalog of small extended sources is considerably less complete and reliable than the point 
source catalog due to the cruder algorithms used and to the effects of confusion from larger extended 
structures. These results are discussed in Section VIII.E. 

B. Sky Coverage 

The survey strategy was chosen to achieve a high degree of homogeneity of coverage of the sky. 
One hours-confirming (HCON) coverage is defined as the sky coverage obtained by two or more scans 
covering a given area of the sky with a separation in time ranging from one orbit of the satellite, 103 
minutes, up to a maximum of 36 hours. Figures I.C. 1 and III.D.3 show the distribution of the number of 
HCON coverages in equatorial and ecliptic coordinates. More detailed plots of the sky coverage are 
shown in Chapter XIII. 

Ninety-six percent of the sky was covered with the two or more HCONs required for sources to be 
considered for inclusion in the catalog. A three HCON coverage was achieved over 72% of the sky. 
Roughly 15% of the sky received more than three HCON coverages. Areas receiving more than three 
HCON coverages arose mainly from the minisurvey (see below), from the overlap between lunes used in 
the survey strategy, from rescans scheduled to fill in coverage holes and from the tendency of scans to 
overlap at high ecliptic latitudes. 

From the standpoint of the IRAS survey itself, the area of the "minisurvey" was the most important 
area receiving extra coverage (Section III.D. 1 1). The particular area was selected because it was available 
immediately after the telescope cover was ejected. The region was surveyed intensively to check the 


VIII- 1 



source detection and confirmation algorithms and to verify that the chosen survey strategy provided the 
desired completeness and reliability; see Rowan-Robinson et al. (1984). The minisurvey consisted of two 
strips of the sky centered approximately at ecliptic longitudes 60° and 250' and contained areas both of 
high source density near the Galactic plane and of reasonably low source density. Parts of the minisurvey 
received as many as four HCON coverages during the minisurvey proper. The whole minisurvey area 
received three more HCON coverages during the main survey. Thus some regions received as many as 
seven HCONs. 

The part of the sky with no HCON coverage is entirely contained within two strips on opposite 
sides of the sky each roughly 5° wide and 60" long. This part of the sky was not observed because of the 
operational difficulties explained in Section III.D.5. At the borders of the gap there are strips 0.5° wide 
that received only one HCON coverage. In addition, there are various small areas of the sky which 
received only one HCON coverage. For the single HCON areas, comprising in total some 500 sq. deg, 
there are no entries in the catalog. The main cause of the small holes was the inability to take useful data 
when the spacecraft passed through the South Atlantic Anomaly (the SAA, see Section III.D.4). These 
areas are mainly concentrated in the south (see the detailed area coverage plots in Chapter XIII). Hours- 
confirmed detections made in the single HCON sky are contained in the catalog of rejected point sources. 

C. Point Source Confusion 

It is the aim of this section to show that confusion caused by background fluctuations due to the 
large numbers of sources below the detection threshold is insignificant in the IRAS point source catalog 
outside of the Galactic plane and other high source density areas. 

Suppose that the mean density N (/ v ) of sources on the sky with flux density greater than / v obeys a 
power law: 

N(fy v) — Kf v a (VIII.C.l) 

where a (<0) is a constant. In a small interval in flux density (f v -f v + d/ v ) and within a beam of solid 
angle Cl one expects on average Cl a N (f v ) / v 7 1 df v sources. The actual number will be a stochastic vari- 
able with a mean-squared deviation equal to the average value. The mean-squared deviation in the flux 
will be / v 2 x the deviation in the number. Thus the mean-squared flux deviation, (A/ v ) 2 , from all 
sources below some threshold /vo is given by 

(A/ v ) 2 - Cl (VIII.C.2) 

Equation (VTII.C.2) can be used to estimate the effect of confusion on the IRAS survey in two ways. 
First, the confusion noise can be calculated directly from the equation using the observed density and dis- 
tribution of sources. Differential sources counts, dN/d(log/ v ), are given in Fig. VIII.C.l a-d for sources in 
all wavelength bands with Ibl > 50". The plots demonstrate that at high Galactic latitudes, objects 
detected at 12 and 25 pm follow a power law distribution (Eq. (VIII.C.l)) with an index a equal to -1.0 
and that sources brightest at 60 pm, mostly external galaxies, have a equal to -1.5. Sources brightest at 
100 pm do not follow a simple power law due to the effects of cirrus and will be discussed separately in 

VIII-2 


1 III 



(NP) OOl (NP) OOI 


vm-3 


Figure VHI.C.1 Differential source counts as a function of flux density for the four wavelengths for 
sources with Galactic latitudes \b I >50°. 






Section Vm.D.5. The typical source density for sources at 12, 25 and 60 pm with flux densities brighter 
than 0.5 Jy is 0.6 source per sq. deg at high Galactic latitudes. For typical detector areas the noise level 
due to confusion as calculated using Eq. (VIII.C.2) is "15 mJy, far below the instrumental detection lim- 
its. 

The second way to interpret Eq. (VHI.C.2) is to note that the critical term in determining the 
importance of confusion is the product £2 N which is the total number of sources per beam area. At high 
Galactic latitudes, the observed value of Q N is less than 0.001 sources per beam for sources above the 
threshold of 0.5 Jy. The source densities over the whole sky for the four wavelength bands are shown in 
Figs. V.H. 1.1-1.4. At the transition between the white and grey areas of these figures, the source density 
is less than 1 per 50 beams; this means that the confusion noise at the transition is less than 0.2 of the 
completeness limit at all wavelengths. Confusion noise can therefore be ignored in the white areas of Fig. 
V.H. 1.1 -.1.4. At 100 pm the shape of the unconfused regions is irregular and artificial due to the pres- 
ence of infrared cirrus; the entire sky except at the highest Galactic latitudes can be confusion limited due 
to local infrared cirrus. In the darker areas of these figures confusion noise dominated the source recog- 
nition process, as will be discussed in Section VIII.D.5. 

A signal-to-noise ratio of five which corresponds to the IRAS completeness limits at 12, 25 and 60 
pm would, according to Eq. (VTII.C.2), be achieved at 75 beams per source for a — 1.5 and at 25 beams 
per source for a— 1.0. This latter value was used as the confusion limit for the high source density pro- 
cessor. 


D. Point Source Catalog Reliability and Completeness 
D. 1 Definitions. Assumptions and Limitations 

The completeness of the catalog above a given flux density at a specified wavelength is defined as 
the fraction of true sources above that flux density which are present in the catalog. The completeness 
can be quantified by statistical analysis of the detection history in an area with multiple HCON coverage 
under the assumption that two different detections of a real source are two statistically independent 
events, each with the same probability. 

The reliability of the catalog is defined as the probability that a catalog entry corresponds to a true, 
celestially fixed source of radiation. An unreliable catalog entry would be created whenever a spurious 
HCON is confirmed by another spurious HCON on any subsequent survey scan. The survey strategy 
was developed so that observations of one area were repeated several times, thus lowering the probability 
that a chance detection would be confirmed. 

Ideally, the completeness and reliability of sources with flux densities in a given range would be cal- 
culated separately for each wavelength. However, because band-merging occurred before hours- 
confirmation, completeness and reliability unavoidably depend on the spectrum of a source. In particu- 
lar, two very different definitions of reliability must be understood. The reliability of a source, which is 
primarily discussed in this chapter, refers to the probability that a source in the IRAS catalog is a true 
celestial object measured in at least one wavelength band. The reliability of a flux measurement, on the 
other hand, refers to the probability that the flux density quoted for a source in a given band is a true 
measurement and not one due to spurious effects such as noise, radiation hits or asteroids. The concept 




of high and medium quality flux densities is crucial in this discussion (Section V.H.5). A high quality flux 
density measurement is one which has met all the confirmation criteria on timescales of seconds, hours 
and weeks. The existence of a source is guaranteed and the quality of its flux density measurement in a 
band is assured according to the completeness and reliability values discussed in this chapter if, and only 
if, the source has a high quality flux density in that band. If a source has only a moderate quality flux 
density in a given band, then lower standards of reliability apply (Section VIII.D.4.c). This dual concept 
of reliability was necessitated by a desire to provide flux information in as many bands as possible for 
sources whose existence was guaranteed in at least one band. 

The estimates of completeness and reliability given here are valid only for sources that are truly 
point-like. Sources which are slightly extended may still achieve an acceptable correlation with the 
point-source template in the in-scan direction, but extend across more than three detectors in the cross- 
scan direction. The point-source processing software resolved such sources into a string of two or more 
point sources with a spacing of a few arc minutes in the cross-scan direction. Successive passes may have 
resolved the source in different ways, causing the completeness of such sources to be poor. This problem 
is a common occurrence with "cirrus" at 100 pm and in all bands in the Galactic plane near the Galactic 
center. In regions of high source-density such strings of sources were discriminated against by the 
"confused-neighbor" and "weaker-neighbor" rejection criteria (see Section V.H.6). 

In addition, the completeness estimates are valid only for point sources without close neighbors. 
Close neighbors can (i) cause the detections of both sources to be lost when the confused detections of 
both sources no longer match a point source template; (ii) cause detections from one source to be lost 
because of source shadowing (Section V.C); (iii) create cross-scan confusion and alter significantly the 
cross-scan position of the source. 

D.2 Formalism for the Determination of Completeness and Reliability 

D,2.a Completeness 

It is necessary to estimate the fraction of real sources on the sky, above some limiting flux density 
value, which is actually present in the catalog. Let p be the probability that a genuine source in a given 
flux range fails to generate an HCON. The probability P{N,M ) that a source is detected on N HCONs 
out of a possible M HCONs is given by: 

P(N,M) - M\-p fp M - N (VIII.D.l) 

where M C N is the binomial coefficient. The completeness C(N miD ,M) of a N m JM survey, which requires 
a source to have at least Amin HCONs out of M possible in order to be included in the catalog, is then 
given by: 

M 

C(N mia ,M) - £ P(N,M) (VIII.D.2) 

N-N miD 


For a survey where two HCONs are required out of a possible two HCONs, the completeness C(2,2) is 
given by: 

C( 2 , 2 ) = (l— p) 2 . (Vm.D.3) 


VIII-5 



Seventy two percent of the sky was surveyed three times and a source was accepted if it was seen at least 
twice. For this "two-out-of-three" strategy, the completeness C( 2,3) is given by: 

C( 2,3) - 3(1 -p) 2 p + (1-p) 3 

-(l+2p)(l— p) 2 (VIII.D.4) 

where the first term on the right-hand side of the first equation corresponds to having seen the source 
twice out of the three possible times while the second corresponds to seeing it all three times. The com- 
pleteness of the 2/3 survey is therefore a factor (1+2/?) greater than that of a 2/2 survey. This factor is 
appreciable since most sources are near the completeness limit where p approaches 1. In small areas of 
the sky the completeness will be better than suggested by Eq. (VIII.D.4) since there were as many as 23 
HCON coverages, yet only two HCONs were required for inclusion in the catalog (see Table VII.B. 1 ). 


D.2.b Reliability 

Let q be the probability, per unit area of the sky, that a false source is created on a single HCON 
coverage of the sky. Let A be the effective area for "weeks-confirmation" (see Sections V.D, V.H.2), i.e. 
the area within which a series of single HCONs must fall if they are to be considered as detections of the 
same source. Then qA is the probability that in a single HCON coverage a false source will be created at 
a given point on the sky with the possibility of being confirmed with another false source at the same 
position on a later HCON. The probability Q(N,M ) that a false source will be created in a given area N 
times out of M HCONs is therefore given by 

Q(N,M) “ mC n (qA) n (1 -qAf 1 - 1 * . (VIII.D.5) 

In fact, Eq. (VIII.D.5) is only true if qA « 1 since it does not allow for the possibility that more than 
one false source is created within A during a single HCON. This is a necessary assumption for the IRAS 
processing because, by its very nature, the processing can only produce one final source per HCON from 
an area A . The results obtained justify the approximation. 

Let u be the density of true sources in a given region of the sky. The probability of detecting a true 
source on N out of M HCONs in a given area of the sky is uA P(N,M) where P(N,M) is given by Eq. 
(VIII.D.2); it is necessary that uA « 1. Hence the total probability T{N,M) of finding a source, either 
false or true, on N out of M HCONs in a given area of the sky is given by 

T(N,M) ~ Q{N,M ) + uA P(N,M) (VIII.D.6) 

where cross terms have been neglected. The reliability R(N,M) of an N/M source is unity minus the 
probability that the source is false. The probability that a given source is false is the probability of 
finding a false source divided by the total probability of finding any source. Therefore, 


R(N m ~ 1 

[Q(N,M) + uA P(N,M)] 


(VIII.D.7) 



D.3 Estimation of Parameters 


The only data available for estimating the parameters are the observed densities n(N,M) of sources 
on the sky, observed N out of M possible times. There is no a priori information as to whether these 
sources are real or false. Let E{x) denote the expectation value of x. Then 

E{ n(N,M )} - \Q{N,M) + uA P(N,M )] x A~ l (VIII.D.8) 

where Q(N,M) and P{N,M) are given respectively by Eqs. (VIII.D.l) and (VTII.D.5) and contain the 
parameters p and q. It is clear that estimates u, p and q of u, p and q can be obtained by replacing the 
expectation value of n(N,M) in Eq. (VIII.D.8) by the observed values, bearing in mind that the value of 
n(0,M) is unobservable. The following equations, from which u, p and q can be estimated, are derived 
using Eqs. (VIII. 1,5 and 8) 

n(l,Af) - Mlq( 1 - qA^' 1 + u(l - p)p M ~ l i (VIII.D.9) 

£ n(N,M) - A~ l [ 1 - (1 - qA~f] + m(1 - p M ) (VIII.D.10) 

iV-1 

M 

£ n(N,M)N— M[q + «( 1 - p) ] (VUI.D.l 1) 

iV-l 

Since (qA) < < 1 in the case of this survey, Eq. (VIII.D.9) and (VIII.D.10) can be approximated by: 

n(l,M) = M[q + u{ 1 - p)p M ~ l ] (VIII.D.12) 

£ n{N,M) = Mq + u(\ - p M ) (VIII.D.l 3) 

AT-1 


The equations can be greatly simplified for sources of high signal-to-noise ratio when p « 1 if it is 
assumed that the true source density is not greatly in excess of the false density. Equation (VIII.D.12) 
becomes 

q = EL LM1 (VIII.D.l 4) 

M 


Equations (VUI.D.l 3) and (VIII.D. 14) give 

M 

u ^ 

A^-2 

while Eqs. (VIII.D.l 1), (VIII.D. 14) and (VIII.D. 15) give 

M 

£ n(N,M)N 

, . N-2 

P " 1 M 

£ n(N,M)M 

N-2 

M 

£ n(N,M)(M~N) 

N^2 

" M 

£ n(N,MW 

N-2 


(VIII.D. 15) 


(VIII.D. 16) 


vm-7 



Each of these equations, which are valid only for bright sources, has a simple interpretation. Equa- 
tion (Vm.D.14) says that the probability of generating false sources is given by the density of single 
HCONs divided by the number of chances of seeing them. The contribution of true sources which have 
lost all but one HCON is neglected. Equation (VIII.D.15) says that all sources with two or more HCONs 
are real. Finally, Eq. (VIII.D. 16) says that the probability of missing an HCON is the sum of all the 
missed HCONs divided by the total possible number of HCONs; again the contribution of true sources 
which have lost M - 1 HCONs is neglected. 

The parameters w, p and q were estimated using Eqs. (VIII.D. 11, 12 and 13). Only preliminary 
results are presented here because these parameters were directly affected by the last processing steps 
(clean-up, flux analysis, and high source density processing) used to generate the catalog. 

D.4 Completeness and Reliability Outside of the Galactic Plane 

D.4.a Estimates from Minisurvey Data 

A preliminary analysis of the 12 pm completeness and reliability is given here for a 70 sq. deg area 
of the minisurvey covered by seven HCONs. It was bounded by ecliptic longitudes 60.5° and 64.0° and 
by ecliptic latitudes -40° and +10° and was thus out of the region of extremely high source densities near 
the Galactic plane. Table VIII.D. 1 gives the number of sources in the seven HCON area, broken down by 
flux density range and number of observed HCONs, as well as the derived estimates of p, q, and 
u. The weeks-confirmation area, A, was taken to be 30"x90". 

The reliability of a two HCON source found in regions of the sky observed with two or three 
HCONs can be deduced from the values in the table and always exceeds 0.99997 in every range of flux 
density. The reliability of two HCON sources in regions with many more than two coverages, such as the 
minisurvey, is, however, considerably less than this value. A source seen only twice in a seven HCON 
area may have a reliability as low as 0.5. Sources with three or more HCONs are, however, almost com- 
pletely reliable at any depth of HCON coverage, given the values of p and q deduced above. 


Table VIII.D.l Completeness and Reliability Data in 7 HCON Area 


Range in 12 |im 
Flux Density 
(Jy) 

N— 1 

2 

Numbers of HCONs 
in 70 sq. deg 7 HCON Area 

3 4 5 

6 

7 

p' 

Parameters 

*2 

Q 

« 3 

0.25-0.32 

26 

5 

1 

0 

0 

0 

0 

0.911 

0.018 

0.699 

0.32-0.40 

31 

4 

4 

1 

4 

0 

0 

0.555 

0.060 

0.208 

0.40-0.50 

22 

0 

3 

2 

2 

5 

7 

0.201 

0.045 

0.270 

0.50-0.63 

9 

0 

0 

1 

2 

5 

13 

0.082 

0.018 

0.300 

0.63-0.79 

6 

0 

0 

0 

0 

3 

8 

0.039 

0.012 

0.157 

0.79-1.58 

10 

0 

0 

0 

0 

2 

18 

0.014 

0.020 

0.286 

>1.58 

0 

0 

0 

0 

0 

0 

25 

0.000 

0.200 

0.364 


1 p is the probability of failing to detect a genuine source in a single HCON. 

2 q is the probability, per sq. deg, that a false source is created on a single HCON coverage of the sky. 

3 U is the probability, per sq. deg, of there being a true source in a given region of the sky. 


I I! 


VTII-8 






The completeness is plotted as a function of flux density in Fig. VIILD. la for a two HCON survey 
and in Fig. VIII.D. lb for a three HCON survey. Note that these figures give the differential complete- 
ness, not the cumulative completeness above a given flux density. The completeness is essentially unity 
above 1.5 Jy for a two HCON survey and above 0.6 Jy for a three HCON survey. It begins to fall shar- 
ply at 0.5 Jy for a two HCON survey and at 0.4 Jy for a three HCON survey. Some of the physical rea- 
sons for missing HCONs are discussed in section VII.E. 3. 

The completeness and reliability values quoted here for 12 pm also apply to 25 and 60 pm sources 
with the sharp fall-off in the completeness occurring at 0.5 and 0.6 Jy at 25 and 60 pm for a three 
HCON survey. The reliability is decidedly worse for sources seen only at 25 pm. The 100 pm values 
cannot be obtained from the minisurvey due to extensive cirrus contamination. 

D,4.b Verification of the Completeness from Source Counts 

The completeness of the catalog can be verified by the differential source-counts shown in Fig. 
Vffl.C.l for the four wavelength bands for \b\ > 50°. The counts fall away sharply at 0.4, 0.5, 0.6, 1.0 
Jy at 12, 25, 60, 100 pm. The completeness at 100 pm is severely degraded by cirrus below these Galac- 
tic latitudes. 


ID 

ID 


5 

o 

u 


1.0 


0.5 


0.0 


(a) 



)2fjL m 

2 HCON 
SURVEY 



Figure VII.D.l a) Completeness at 12 pm calculated for a two HCON survey (top); 

and b) for a three HCON survey (bottom). See XII.A.4 for a revised 
estimate of the completeness. 


VIII-9 





D.4.c Reliability of Point Sources with Flux densities of Moderate Quality 

A moderate quality flux density is based on at least four detections (Section V.H.5), but the 
moderate status means that the source also failed to be detected on at least one occasion when it should 
have been seen. If there were no alibis from dead detectors, the source may have only been detected on 
four occasions out of a possible eight. If the source had a good quality flux density in another band, the 
reality of the source is established in that band. In this case there is a reasonable certainty that the source 
detection in the moderate quality band is real, but the quoted uncertainty in the quoted flux density is 
larger than that for good quality fluxes, typically by a factor of 1.41. Moderate quality flux densities 
should be used with caution. Sources with moderate quality flux densities in two adjacent bands are reli- 
able. 

D.5 Completeness and Reliability in the Galactic Plane 

The completeness and reliability in regions of high source density, defined in Section V.H.6, can be 
estimated by methods similar to those discussed earlier in this section, but such estimates should be 
regarded with great caution. As an example, about 200 12 pm sources in a 4.5 sq. deg area from the 
minisurvey near the Galactic plane were analyzed. The region was covered with seven HCONs and has 
an average source density of 39 per sq. deg. Table VIII. D.2 gives counts of cataloged 12 pm sources as a 
function of flux density and the number of HCONs in this area. 

If the parameters p, q and u are evaluated for the high density regions using Eqs. (VIII.D.14-.D. 16), 
i.e., assuming that all single HCON sources are false and all two or more HCON objects are real, then p 
< 0.2 and q < 1.5 (sq. deg) -1 HCON -1 for / v > 1 Jy. These values imply that for flux densities above 1 
Jy the completeness of the survey in the two HCON high density sky exceeds 63% and in the three 
HCON high density sky exceeds 88%. 

The reliability of a two or more HCON source can be estimated from the value of q and Eq. 
(VTII.D.7). For M equal to 2 or 3 and A equal to 25 beam areas, 0.02 sq. deg, the estimated reliability of 
a two HCON object brighter than 1 Jy at 12 pm exceeds 99.5%. The numbers in Table VIII.D.2 show 
that applying the formalism of Eqs. (VTII.D. 14 - .D. 16) to the high density region is not entirely valid. 
The fact that there is a minimum in the number of sources with a specified number of HCONs at a value 
of three HCONs suggests that sources with fewer than three HCONs are spurious. It should be noted, 
however, that including these false two and three HCON sources in the determination of q decreases the 
estimate of the net reliability of a source only slightly. 

The existence of minisurvey sources (Table VIII.D.2) observed six or seven times in regions of high 
source density indicates that the completeness of bright sources is significant even in these regions. If the 
minisurvey region is representative of the entire Galactic plane region, then the completeness for 
sources with flux densities above 10 Jy will be about 86% for the part of the sky covered with two 
HCONs and about 98% for the sky covered with three HCONs. These results apply to sources observed 
at 12 and probably 25 pm where extended sources are few. At 60 and 100 pm the meaning of the relia- 
bility of a source is less clear because of the large population of extended objects seen against a very com- 
plex background. No attempt has been made to assess either the completeness or reliability of sources at 
these wavelengths. 


VIII- 10 


TTTI : 



Table VIII.D.2 

Number of HCONs in a 7 HCON High Source Density Region 





Numbers of HCONs in High 



Range in 12 |Xm 



Source Density 7 HCON Area* 



Flux Density 








(Jy) 

N-l 

2 

3 

4 

5 

6 

7 

1-2 

49 

9 

1 

6 

13 

17 

30 

2-4 

22 

2 

1 

3 

5 

17 

26 

4-8 

6 

1 

1 

0 

1 

10 

14 

8-16 

0 

0 

0 

0 

1 

3 

6 

16-32 

0 

0 

0 

0 

0 

0 

2 

32-64 

0 

0 

0 

0 

0 

4 

0 

64-128 

0 

0 

0 

0 

0 

1 

3 


* the area consisted of two 1.5° x 1.5° regions with the following Galactic coordinates. 
334.5' < / < 336.0°; -2.0° < b < -0.5° 

336.0° < / < 337.5°; -0.5° < b < +1.0° 


D.6 Galactic Plane Shadow 

The strong lagging of the noise estimator at 100 and 60 pm after passage through the Galactic 
plane prevented the acceptance of a large number of sources because their calculated signal-to-noise ratio 
fell below the threshold. The effect is demonstrated in Fig. VIII.D.2 which shows counts of point sources 
in bins of 1° in latitude and 20° in longitude. The counts of sources seen only at 100 pm shows a 
decrease of about a factor of 10 just after the telescope scanned past the Galactic equator in the first two 
HCON coverages. This dearth of sources is called the Galactic plane shadow and takes roughly the form 
of two rectangular strips. The first extends from Galactic longitude / - 10° to 30°, and from Galactic lati- 
tude b — 0° to 5°, and the second extends from / - 330° to 350°, and b — 0° to -5°. 

During the third HCON survey the telescope passed the Galactic equator in the opposite direction. 
Although many 60 and 100 pm sources were found in the shadow area, they were seen only on one 
HCON and were thus excluded from the catalog. 

The noise shadow is far less important at the shorter wavelengths. The counts of the 12 pm sources 
shown in the figure are highly symmetric around b — 0°, but show a small 10% dip in the central two 
bins. The cause of this dip is probably the high source density clean up (see Section V.H.6) and not the 
noise lag. 


E. Completeness and Reliability of the Catalog of Small Extended Sources 

The completeness and reliability of this catalog will be discussed in the hard-bound version of the 
Supplement when it and the catalog of small extended sources are published in mid- 1985. 


VIII- 11 





GALACTIC LATITUDE (deg) 

Figure VHI.D.2 The wavelength dependent effects of the Galactic plane shadow discussed in the text 
are shown. The solid lines are counts of 12 pm sources as a function of Galactic lati- 
tude; the dotted lines are counts of 100 pm sources. 


Authors: 

M. Rowan-Robinson, P. Clegg, C. Beichman, T. Chester, T. Conrow, H. Habing, G. Helou, G. 
Neugebauer, T. Soifer and D. Walker 

References: 

Low, F. J. et al., 1984, Ap.J., 278, LI 9. 

Rowan-Robinson, M. et al, 1984, Ap.J., 278, L7. 


VIII- 12 


CIRRUS 


IX. THE LOW-RESOLUTION SPECTRA 


A. Instrumentation 
A. 1 Introduction 

The IRAS survey instrumentation included a low-resolution spectrometer which covered the 
wavelengths between 8 and 22 pm. The spectrometer operated during the entire survey, providing spec- 
tra of the brighter point sources. This section briefly describes the instrument, dealing primarily with 
those aspects of its optics and its signal-handling electronics that have a direct bearing on the interpreta- 
tion of the spectra. The spectra are presented in two forms -- a tape version, henceforth called "The Cata- 
log of Low Resolution IRAS Spectra”, or simply "The Catalog", and a hard-copy version, designated as 
"The Atlas of Low Resolution IRAS Spectra", containing graphical representations of the spectra. 

Because the survey function required a passive instrument, a slitless design was selected. This 
design is equivalent to an objective prism spectrograph, oriented in such a way that the dispersion was 
aligned with the scan direction. Obvious penalties of this design are a degraded resolution for extended 
objects, sensitivity to spatial confusion, and short integration times. 

The detectors were sampled continually during the mission, and the data were received on the 
ground together with the other survey data. The spectra were extracted during data processing on the 
basis of point source detections from the survey array. 

A.2 Optical Properties 

The spectrometer had a rectangular aperture mask in the focal plane that measured 6' in the disper- 
sion direction and 15' across. Although the cross-scan width of the focal plane aperture was only half 
that of the survey array width, the overlap between adjacent scans in the survey strategy ensured full sky 
coverage. 

Two overlapping wavelength ranges were scanned simultaneously, one extending from 7.7 to 
13.4 pm and the other from 11.0 to 22.6 pm. The scan length was nominally 6' but an internal 
misalignment obscured the extreme low end of the short wavelength range. The resolution was primarily 
determined by the exit slit width of 15", although diffraction at the telescope aperture and electronic 
filtering caused significant smoothing. Figure IX. A. 1 shows, respectively, the aperture location of a 
source, the effective monochromatic image size, and the resulting spectral resolution as functions of 
wavelength. 

Several detectors were used in each of the two wavelength ranges to reduce confusion problems. 
Three short wavelength detectors each covered 5 ' of the aperture width. At the longer wavelengths two 
detectors each covered 7 . 5 '. 

The optical layout of the spectrometer is shown in Fig. IX.A.2. A back-reflecting KBr prism with 
curved surfaces served the three functions of collimation, dispersion and refocusing in the long 
wavelength spectrometer. A field mirror imaged the telescope pupil onto the prism. The short 
wavelengths also passed through this system, but were given additional dispersion in a second spectrome- 
ter section with a curved NaCl prism, using a field mirror adjacent to the first exit slit. Field optics 
immediately behind the exit slits refocused the telescope pupil on each of the detectors. 



Figure IX.A.l. 


Figure IX.A.2. 




Relative location and width of the image of the exit slit in the field-of-view of the 
telescope, and the resulting spectral resolution are given as functions of wavelength. 
The width and resolution have been calculated taking into account the diffraction by 
the telescope and the electronic filtering. 



The optical layout of the spectrometer is shown schematically.SO is the entrance 
aperture of the spectrometer in the focal plane of the telescope. The field mirror M 1 , 
the curved prism PI and the exit slit SI comprise the primary spectrometer. M2, P2 
and S2 form a secondary system that provides most of the dispersion at the short 
wavelengths. 


IX-2 



A.3 Electronics 


Si:Ga photoconductive detectors for the short wavelengths and Si:As photoconductors for the long 
wavelengths were used in conjunction with trans-impedance pre-amplifiers. The electronics included 
spike suppression circuitry. The electronic bandwidth was 12 Hz and the sampling frequency was 32 Hz, 
corresponding to a sampling interval of 7.2". 

An important aspect of the data handling was the encoding of signals with a large dynamic range 
into an 8-bit format. The output voltages were digitally encoded on a logarithmic scale with increments 
of 3.5% limiting the signal-to-noise ratio for single samples to less than 100. To avoid loss of precision, 
the baselines were kept at low positive levels by AC-coupling and "zero-clamping". Occasionally, the 
zero-clamping affected estimates of the baseline, as discussed in the next section. Fortunately, the 
occurrence of zero-clamping was rare among the spectra selected for the catalog, since the rejection of 
confused spectra tended to eliminate those with zero-clamps. 

Further detail on the spectrometer electronics is contained in the instrument description by Wilde- 
man, Beintema and Wesselius (1983). 

A.4 Effects of the Zero-Clamp 

Figure IX. A. 3 shows a schematic representation of the circuit used for the AC-coupling and zero 
clamping. The voltage across the coupling capacitor equaled the difference between output and input sig- 
nals and thus constituted an error signal. As long as the output signal remained positive, the circuit 



Figure IX.A.3. The AC-coupling and zero clamping circuit. Zero clamping occurred when the input 
signal was decreasing, resulting in an output signal of zero. 


IX- 3 




behaved like a high-pass R-C filter with a time constant of 10 seconds (a spectral scan lasts 1.5 seconds). 
The fact that the error signal changed at a rate proportional to the output signal was easy to correct. 
However, as soon as the output tried to drop below zero, the capacitor was instantaneously recharged to 
prevent any further drop; the output signal was clamped to the zero level. The error signal was now the 
inverse of the input and could rise rapidly, while the output indicated no rate of change. The clamp 
remained active as long as the input signal continued to decrease. As soon as the input signal rose again, 
the clamping ended and linear behavior was restored. The output signal again tracked the rate at which 
the error signal changed. 

Normally, zero clamps occurred just after spectral scans or during negative noise excursions on 
"empty sky." However, a steeply decreasing background could continually clamp the output at zero, 
except during a spectral scan. In such a case, the presence of a sloping baseline would go unnoticed. 
Because of the direction in which the spectra were scanned, the result is then an underestimate of the sig- 
nal at the shorter wavelengths of both spectral ranges. A rising background would only raise the baseline 
in the output signal with no worse effect than a reduced precision of digitization. Figure IX.A.4 illus- 
trates the effects of background slopes. 

A. 5 Summary of Instrumental Characteristics 

The design of the instrument and the planning of the IRAS mission resulted in the following global 
characteristics of the spectra: 

- spectral coverage from 8 to 22 pm in two overlapping ranges; 

- a spectral resolution for unresolved sources varying from 20 to 60; 

- a sky coverage identical to the coverage of the IRAS survey; 

- typically 2 or 3 observations per spectrum, together providing about 0. 1 5 seconds of integration 
time per spectral element per source; 

- slit widths corresponding to 15" and instantaneous fields of view of 30’ square (short 
wavelengths) and 45' square (long wavelengths); 

- quantization steps of 3.5%, resulting in an rms digitization error of 1%; 

- potential problems from negative backgrounds gradients. 


B. Performance and Calibration 

B. 1 Detectors 

The global properties of the detectors in the spectrometer are listed in Table IX.B.l. The weekly 
averages of the responses to flashes of the internal reference source (Section II.C) stayed within 1% of the 
values given, although systematic variations of up to 15% were found on shorter timescales. These were 
associated with passages through the South Atlantic Anomaly (SAA) and across the Galactic plane. The 
spectra were corrected for these effects by using a flash response interpolated between the most recent 
flash and the next one. As described below, other correction factors (especially the cross-scan gain) 
introduced larger uncertainties in the flux scale of the spectra. 

The noise estimates given in Table IX.B.l are based on baseline readings, and refer to a single spec- 
tral resolution element. Estimates for individual spectra deviate from the mean value by typically 20%. 



Figure IX.A.4. The effects of AC-coupling and zero clamping on flat spectral scans with sloping 
backgrounds. From top to bottom: input signals, output signals and reconstructed 
input signals. Absolute levels are relevant only for the output signals. During recon- 
struction, the possibility of zero-clamping is ignored. Within a spectral scan, 
wavelength decreased with time. 



Table IX.B.l 

Detector Characteristics 


Detector 

Response to 

Noise level 

NEFD 

Wavelengths 

Number 

Internal Reference 





(mV) 

(mV) 

(Jy) 

(pm) 

i 

11.9 

0.10 

1.4 

8 - 13 

2 

4.4 

0.04 

1.6 

8 - 13 

3 

7.4 

0.06 

1.3 

8 - 13 

4 

22.8 

0.12 

3.0 

11-22 

5 

46.3 

0.20 

2.5 

11-22 


IX-5 





B.2 Wavelength Scale 

The conversion from in-scan position to wavelength was adjusted and monitored using the lines in 
the spectrum of the planetary nebula NGC 6543. The wavelength scale is accurate within 0. 1 pm. 

The zero point of the wavelength scale could be determined by the in-scan position of the source as 
determined by the survey instrument during the same scan. The uncertainty in this position caused a 
jitter in the zero point of up to 1 sample (corresponding to 7" or one-half of a spectral resolution ele- 
ment). This effect degraded the correction for wavelength-dependent gain and the averaging of individual 
spectra. A better estimate of the zero point was obtained using the four well-defined band edges of the 
spectrometer. A signal-to-noise ratio of 10 measured at one of the edges gave an accuracy of 0.3 sample 
( 2 "). 

B.3 Cross-Scan Responsivity * 

The cross-scan responsivity of the detectors was determined by finely spaced raster scans across the 
star R UMi (as shown in Fig. IX.B.l.) where the curves are normalized over the central half of the detec- 
tors. The sharp dip in response of detector 4 was also found in the laboratory tests. Although the curves 
are accurate within a few percent, the derived value of the cross-scan correction is much less accurate. 
Because the cross-scan position determined by the survey instrument during a scan had a typical 
uncertainty of 1 '. Except for sources passing over the very center of a detector this introduced an uncer- 
tainty in the cross-scan correction factor of 10-20%. 

B.4 Wavelength-Dependent Responsivity 

In order to determine the factor for conversion of sample values into flux densities, the observed 
spectra of a Tau were compared with a black body spectrum of 10,000 K. The resultant responsivity 
curves (Fig. IX.B.2) show the ratios of flux density to sample value as a function of wavelength. The 
curves have been normalized in the overlap region from 11 to 13 pm and are accurate to 2% at the shor- 
test and to 4% at the longest wavelengths. No significant differences were found between detectors 1 , 
2 and 3, or between detectors 4 and 5. All detectors had somewhat better responsivity in the center of 
their wavelength range than at the edges. Except near the very edges of the band, the responsivity 
correction did not exceed 30%. The feature seen between 9 and 1 1 pm in the short-wavelength band is 
a characteristic of the instrument that is found in all raw spectra and was already known from laboratory 
tests. 

B.5 Radiation Effects 

Spikes caused by energetic particle radiation hits were removed by the electronic deglitcher without 
significantly increasing the noise. Only near the maximum of the SAA, where no regular observations 
were taken, were residual glitches and increased noise found. The polar horns appear to have had no 
effect on the detectors. Outside the SAA, baseline drifts on a time scale of 1.5 sec, roughly the dwell time 
of a point source on a detector had an amplitude of less than twice the noise level. Inside the SAA, the 
drifts increased by about a factor of ten. 

B.6 Multiplexer Glitches 

Occasionally, the data streams of the detectors were scrambled by the multiplexer. In most cases 
this was evident in the spectra as single-sample spikes. The larger ones were easily recognized and 

IX-6 


I 1 1 : 



RELATIVE RESPONSIVITY RELATIVE RESPONSIVITY 




Figure IX.B.l. Cross-scan responsivity for each of the five spectrometer detectors. Detectors 1, 2 
and 3 cover the 8-13 |im wavelength range, and detectors 4 and 5 cover the 
1 1-22 pm wavelength range. Open circles indicate the R UMi results; the curves were 
used for the correction (Section IX.C.2.e). 


IX-7 




1.50 



13 15 17 

WAVELENGTH (^m) 


Figure IX.B.2. Wavelength dependent responsivity for the 8-13 and the 11-22 pm detectors, respec- 
tively, plotted on the standard samples (Section IX.C.2.c). The responsivities are 
normalized in the 11-13 pm region. 


removed (see Section IX.C). In some cases a series of samples was scrambled and the spectrum could not 
be repaired without running the risk of erasing real features. 

B.7 Confusion 

Since the slitless spectrometer had a field-of-view of 5' by 6' (short wavelengths) or 7.5' by 6' (long 
wavelengths), regions that were confused in the short wavelength survey bands were also confused for the 
spectrometer. A spectrum could be confused by a nearby point source, or by extended structure associ- 
ated with the source itself or in the background. In most cases such confusion was recognizable as a 
difference in the baseline level between the two sides of each spectrum-half. However, as discussed 
below, if the test of baseline asymmetry was made too restrictive, otherwise good spectra were lost. 

B.8 Photon Induced Responsivity Enhancement 

No evidence was found for photon induced responsivity enhancement (Section IV.A.8 and VI.B.4c), 
although it would have been hard to discover. In the first place, most spectrometer sources were stars or 
star-like sources with flux densities gradually decreasing towards longer wavelengths. Since the spectra 
were scanned starting at the long wavelengths, and thus at the lowest flux density, the deformation of the 
spectra was minimized. Secondly, close to the Galactic plane, where the effects were greatest, there were 
very few clean, isolated point sources. Confusion probably masked any effects of photon induced 
enhancement present. 


IX-8 




B.9 Memory Effects 

Changes in the detector responsivity on a time scale of minutes due to exposure to a bright source 
or to a prolonged exposure to a medium bright source like the Galactic plane were found to occur. The 
latter effect was compensated for to within 5% by using interpolated internal reference flash responses to 
correct the flux densities. 

Memory effects did not cause differences in response of the system between spectral lines and con- 
tinuous spectra. The planetary nebula NGC 6543 was scanned at the nominal survey rate and at 1/2, 
1/4, and 1/8 of that speed. The line intensities did not change noticeably as a function of the different 
scan speeds. 

B.10 Linearity Checks 

It is expected that any non-linearity effects were small, although no comprehensive linearity tests 
were made. This statement is based on the agreement between two different determinations of the wave- 
length dependent responsivity curves. One set of curves was determined using a Tau, a star with a 
spectrum that rises steeply towards shorter wavelengths. The second set of curves was determined using 
asteroids -- cool objects which are brightest at the longer wavelengths. The differences between the two 
sets were small, and the comparison was limited by the accuracy with which the asteroid curves could be 
determined. 

B. l 1 Overall Flux-Density Scale 

After processing a large sample of spectra with the correction techniques described in Section IX.C, 
the integrated fluxes in the spectra (after convolution with the survey pass-band) were compared with the 
fluxes measured by the survey array at 12 pm. Comparison with the 25 pm survey fluxes was not 
attempted because the spectrometer hardly overlaps the survey 25 pm pass-band. Systematic factors of 
0.75 and 1.00, respectively, were applied to the data to make the 8-13 and 1 1-22 pm spectrometer fluxes 
agree with the survey fluxes. On the average the spectrometer flux densities in the catalog are con- 
sistent with the survey flux densities to within 10-15%. 

C. Data Processing 
C. 1 The Database 

The spectrometer data consist of three types of data: (i) uncorrected spectra with header informa- 
tion, (ii) calibration tables, and (iii) administrative files. The spectra were extracted out of the data 
stream whenever an hours-confirmed source with a signal-to-noise ratio greater than 25 or a source 
specified as "known source" (Section V.D.4), crossed the spectrometer aperture. Extractions were also 
made for designated calibration sources. The spectra were linked to survey sources using index- 
association records produced by the hours- and weeks-confirmation processors (Section V.D.7). 

There were three types of calibration tables: (a) the responses of the five detectors to the internal 
reference source flashes and the intensities of the flashes as derived by the survey calibration processor; 
(b) correction tables for the relative responsivity as a function of position across the five detectors; (c) 
correction tables for the relative responsivity as a function of wavelength. The cross-scan and wavelength 
dependent responsivity tables were derived from special observations (see Section IX.B). 


IX-9 



C.2 Processing the Individual Spectra 

Processing of the data always started with the raw data. This allowed the correction procedures to 
be improved continuously up to the time of production of the catalog. The major processing steps are 
discussed below. 

C.2.a Despiking 

Single sample spikes with an amplitude greater than 8% caused by multiplexer errors (Section IX. B) 
were removed and replaced by an interpolated value using a simple algorithm operating on the raw data. 
Multiple-sample spikes were not removed, but their presence was noted so that the spectrum-half would 
be rejected later (see Section IX.C.3). 

C.2.b Conversion to a Linear Scale 

A lookup table of 256 entries was used to convert the raw data to voltages on a linear scale.A stan- 
dard reconstruction of the input-voltage to the high-pass filter from the measured output voltage was car- 
ried out. An offset correction was reset to zero whenever the sample voltage dropped below a specified 
threshold because of the effects of the zero clamping (Section IX.A.4). 

C.2.c Interpolation to a Standard Wavelength Grid 

For ease of processing, an interpolation was carried out to a standard regular grid of angular posi- 
tions in the dispersion direction. Because the dispersion of the spectrograph changed rapidly as a func- 
tion of the angular position (Fig. IX. A. 1), the wavelength values corresponding to the standard sample 
values were not equidistant. Before the interpolation, allowance was made for variations in the scan 
speed of the telescope. For the spectra of the brighter sources (signal-to-noise ratio greater than 10) the 
well-defined in-scan detector edges were used to center the spectrum. The centering correction reduced 
the in-scan errors to approximately 2 " , corresponding to approximately 0.03 pm in wavelength. 

C.2.d Correction for Wavelength-Dependent Responsivity 

The interpolated samples were multiplied by a responsivity table sampled at the same standard grid. 
There was a table for each of the five detectors derived from observations of a Tau (Fig. IX.B.2). 
Although the software allowed selecting a different table for each of 16 regularly spaced cross-scan posi- 
tions on the detector, the evidence for a cross-scan variation of the wavelength-dependent gain was too 
weak to justify using this option. 

C.2.e Correction for Cross-Scan-Dependent Responsivity 

Depending on the nominal cross-scan position of the source, a correction (Fig. IX.B.l) was applied 
for the decrease of responsivity towards the edges of the detectors. This correction was the weakest link 
in the calibration process because of the relatively large uncertainty in the cross-scan position. The 
correction applied is uncertain by up to 20%, although this uncertainty was decreased by the process of 
joining the two spectrum-halves together (see Section IX.C.2.g). 

C.2.f Overall Responsivity 

The overall responsivity depended on the individual detector, on the time and/or on the sky posi- 
tion. To account for these variations a correction was derived from the voltage responses to the two 
internal reference source flashes bracketing the time of observation. After applying the responsivity 


IX- 10 





correction to a large sample of spectra, the integrated fluxes in the spectra were compared to the fluxes 
measured by the survey array. Systematic factors of 0.75 and 1.00 were applied to the integrated spec- 
trometer fluxes to bring them in line with the survey observations. 

C.2,g Joining of the Two Spectrum Halves 

The two spectrum-halves (8-13 pm and 11-22 pm) were treated independently until this point. 
Because of uncertainties in the cross-scan position of the scan path over the detectors and therefore in the 
nominal cross-scan responsivity correction, the two spectrum-halves often differed by up to 15 or 20% 
after subtraction of a linear baseline. The overlapping portion of the spectrum-halves was used to deter- 
mine another correction factor. In doing so, the nominal relative cross-scan positions were used to deter- 
mine which half of the spectrum to change by the largest amount. If either half had been observed by 
the central part of a detector, it was considered reliably calibrated, and that portion was not changed by 
the joining process, and the half of the spectrum observed near the edge of a detector was shifted up or 
down towards the other half. If both halves were considered equally reliable, then each was scaled by the 
square root of the ratio between the overlapping sections. This joining process reduced the overall error 
in the responsivity correction to less than 10%. 

C.3 Averaging Spectra, Quality Checks 

Before spectra were averaged, a number of quality checks were performed on the individual meas- 
urements of the two halves of a source’s spectrum. First, all measurements made within 18" of the edge 
of any detector were flagged. Measurements were rejected: 

(a) if they contained multiple-sample spikes (Section IX.C.2); 

(b) if the join-factor obtained before (Section IX.C.2) was outside the range 0.30 to 3.3; 

(c) if the measurement was confused by neighboring sources; this was considered to occur when 
the measurement met one of two criteria: (i) the central portion was below the baseline deter- 
mined from signal-free parts of the spectrum-half; (ii) the baseline at the low wavelength end 
of the spectrum-half differed from that at the high wavelength end by more than 20% of the 
signal in the 8-13 pm band or by 10% in the 1 1-22 pm band. The lower limit of these thres- 
holds was 2.5 times the sample noise. 

(d) if the measurement did not correlate with the "reference measurement", defined as the meas- 
urement with the smallest number of check-flags. This choice gave preference to measure- 
ments passing over the central part of the detector. Any 8-13 pm measurement for which the 
correlation coefficient with the reference was below 60% or any 11-22 pm measurement for 
which it was below 50% was rejected. For line spectra without a continuum in the 8-13 pm 
region the first criterion was waived. 

At least 80% of the spectra in the catalog had correlation coefficients above 70 and 60% in the short 
and long wavelength halves, respectively. Some 40% correlated internally with coefficients better than 
80% in both spectrum-halves (see Section IX.C.4). 

The spectrum-halves passing through all of the above tests were averaged using the inverse of' the 
square of the noise as a weighting factor. At least two measurements in each of the two spectrum-halves 


IX- 11 



(8-13 and 1 1-22 pm) had to be accepted before the spectrum could be averaged and included in the spec- 
tral catalog. 

After averaging, the two spectrum-halves were rejoined, giving both halves equal weight (see Section 
IX.C. 1). Generally the join factors differed from 1 .00 by only a few percent. 

The averaged spectrum was convolved with the 12 pm survey passband. The integrated flux thus 
obtained, was compared to the average 12 pm survey flux of the source. The ratio between the two 
fluxes is given in the low resolution spectrometer catalog record and has a 1 ct dispersion around unity of 
about 15%. Exceptions to this rule will be spectra with sharp lines (classes 8 and 9; see Section IX.D.2) 
or small 1 2 pm fluxes. 

C.4 Final Selection of Spectra 

Three selection criteria were applied for inclusion in the Catalog of Low Resolution Spectra. 

(i) The source is contained in the IRAS point source catalog. 

(ii) The entire spectrum must have been observed at least twice and the individual measure- 
ments should be mutually consistent. Individual spectra must pass all the checks mentioned 
in Section IX.C.3 and must have a minimum correlation coefficient of 50% between any two 
measurements of the source. The large majority of spectra had, however, much higher corre- 
lation coefficients (see below). 

(iii) The source must pass a subjective visual inspection. About 2.5% of all sources were rejected 
by this process, mostly because they showed non-point source characteristics or confusion 
with other sources. 

Four samples of sources were selected for inclusion in the catalog. 

a. Sources whose 12 pm survey flux density was larger than 25 Jy or whose 25 pm survey flux density 
was larger than 50 Jy. Individual spectrum-halves were required to correlate with each other with a 
correlation coefficient of 80% in either spectrum-half. This sample contains about 2 1 50 sources. 

b. Sources with 12 pm flux densities larger than 1 Jy or 25 pm flux densities larger than 2 Jy but not 
contained in Sample a. The vast majority of these sources proved to be brighter than approxi- 
mately 5 Jy at 12 pm or 10 Jy at 25 pm. The correlation coefficients were required to exceed 70% 
in the 8-13 pm band and 60% in the 1 1-22 pm band. This sample contains about 2450 sources. 

c. Sources in the same flux density range as sample b but with lower correlation coefficients: between 
60 and 70% in the 1 1-22 pm spectrum-half or between 50 and 60% in the 1 1-22 pm half, respec- 
tively. There are roughly 850 sources in this sample. 

d. Sources with minimum survey flux densities of 1 Jy at 12 pm or 2 Jy at 25 pm. Of these sources 
only the 1 1-22 pm spectrum-half was required to have been measured consistently with a minimum 
correlation coefficient of 50% between individual measurements. Out of this sample only sources 
with specified spectral lines and not contained in sample a, b, c were selected for the catalog. The 
selection was carried out by the classification program: only classes 8 and 9 (see Section IX.D.2 and 
Table IX.D.l) were kept. There are approximately 40 spectra with lines in this sample. 

Samples a, b, and c contain continuum sources and approximately 75% of the line sources of the 
catalog; sample d contains the remainder of the line sources. 


IX- 12 



* 


D. Classification 
D. 1 Introduction 

Visual inspection of spectra shows that they fall naturally into a number of classes. A characteriza- 
tion program was used to return a two-digit code for each spectrum. The first digit characterizes the 
overall shape of the spectrum (main class) and the second digit gives quantitative information on the 
dominant feature in the spectrum (subclass). 

This classification code appears to be an adequate description of the spectra of more than 95% of 
the sources in the catalog. Although the classification is based on spectrometer data only, and therefore 
independent of other astrophysical classification schemes, most classes are dominated by a well-known, 
well-defined type of sources. 

D.2 Classification Scheme 

D.2.a Overview 

There are five general types of spectra distributed in nine main classes. The five types are. Blue 
continuous energy distributions, Red continuous energy distributions, spectra with spectral lines including 
the 11.3 pm line, spectra with lines but not the 11.3 pm line and "others" which fit into none of the 
preceding classes. 

As a first step blue and red spectra were separated using the long-wavelength part of the spectrum. 
For red sources the flux density per octave (kf\) rises from 14 to 22 pm, and for blue spectra it decreases. 

The second step was the determination of the relative excesses or deficiencies in narrow bands (0.5 
pm wide) centered on the two broad features often seen in the spectra: the 10 pm silicate band and the 
1 1 pm SiC band. The local "continuum" level was estimated by a logarithmic interpolation between nar- 
row bands just outside the features. Also determined were the excess fluxes in 7 narrow bands (0.5 to 1 
pm wide) centered on the emission lines at 9.0 pm [Ar III], 10.5 pm [S IV], 1 1.3 pm (unidentified), 12.8 
pm [Ne II], 14.5 pm [Ne V], 15.5 pm [Ne III], and 18.8 pm [S III]. Those were the only lines that were 
seen in the spectra with the possible exception of the [Ne VI] line at 7.5 pm, just on the edge of the 
pass-band. The local continuum was estimated by a linear interpolation between two similar bands on 

both sides of the line. 

Table IX.D. 1 summarizes the classification scheme described below and lists the number of sources 
of each type. Figures IX.D.2.1-3 give examples of some of the different kinds of spectra. 

D.2.b Sources with Blue Energy Distributions: Classes ln,2n,3n,4n 

Sources with blue energy distributions without the 11.3 pm or other narrow spectral lines were 
classified according to the relative strengths of the 10 and 1 1 pm emission or absorption features. Spectra 
with no emission or absorption features at all were assigned to class In, where the subclass n is equal to 
twice the absolute value of the spectral index between 8.0 and 13.0 pm. The spectral index, P, is defined 
according to 

(IX.D.l) 

Normal K stars have no emission or absorption and show a Rayleigh-Jeans tail at these wavelengths with 
p--4. Such stars are thus assigned code 18. 


IX- 13 



ORIGINAL PAGE B 
OF POOR QUALITY 


SUBCLRSS FOR 3 SUBCLASS FOR 2 



Figure IX.D.l. The classification scheme for blue spectra. (See text.) 


LO 

c 

□0 

n 

r~ 

ZD 

LO 

LO 

~n 

o 

UO 

4= 


Sources with blue energy distributions showing 10 pm silicate band emission were assigned to class 
2 n, sources with silicate absorption to class 3 n. For class 2 n the subclass n was defined according to the 
strength of the emission band with respect to the adjacent continuum following the equation: 

" - 10 x IlnA(9.8pm) - (0.589 ln/ x (7.9pm) + 0.41 1 ln/ x (13.3pm)] (IX.D.2) 

For class 3 n, the subclass n was derived on the basis of the strength of the absorption band: 

« - -5 x rinA(9.8pm) - (0.589 lnA(7.9pm) + 0.41 1 ln/ x ( 13.3pm)] (IX.D.3) 

Sources with blue spectra showing 1 1 pm SiC band emission were assigned to class 4 n with the sub- 
class n derived from the strength of the band according to 

« - 10 x [ln/x(l 1.4pm) - (0.506 ln/x(9.8pm) + 0.494 ln/ x (13.3pm)] (IX.D.4) 

Figure IX.D. 1 shows how the classification depends on the ratios of the total flux to the narrow 
band fluxes at 10.0 and 1 1.0 pm. In the figure the uncertainties in the positions in the points is less than 
0.01, approximately the size of the plotting symbols. The cluster of sources at the center of the figure 
shows objects with no significant excesses or deficits at either wavelength (Class In). The loci of the other 
blue main classes and subclasses are indicated on the figure. 

D.2.c Sources with Red Energy Distributions: Classes 5n.6n.7n 

Sources with red energy distributions (as defined above), but without the 1 1.3 pm line or other nar- 
row spectral features, were classified according to the presence or absence of 10 pm silicate emission or 
absorption. No 1 1 pm emission or absorption features due to SiC were seen in these red sources. 


IX-14 



Table IX.D.l Spectral Classification Scheme 

Class 

Characteristic 

Number 

Typical objects 

On 

other class 

-subclasses /z -0,2, 3, 4: 
see text 
-subclass «— 1: 
blue , low S/N 
-subclass n— 5: 
red, low S/N 

363 

unknown 

In 

blue, featureless 
-subclass: n= 2 times 
spectral index 

2246 

stars with spectral type 
earlier than M5 

2 n 

blue, 10 pm emission 
-subclass: n— 
band strength 

1738 

stars with not too thick 
oxygen-rich envelopes 

in 

blue, 10 pm absorption 
-subclass: n- 
band strength 

230 

stars with thick 
oxygen-rich envelopes 

4 n 

blue, 1 1 pm emission 
-subclass: n- 
band strength 

542 

stars with carbon-rich 
envelopes 

5/i 

red, no line, 
no 10 pm band 
-subclass: n - 2 times 
spectral index 

64 

unknown 

j 6/i 

red, 10 pm emission 
-subclass: n= 
band strength 

78 

stars with very thick 
oxygen-rich envelopes 

In 

red, 10 pm absorption 
-subclass: ri- 
band strength 

67 

stars with extremely 
thick oxygen-rich 
envelopes and hot spots 
in molecular clouds 

8/i 

1 1.3 pm emission line 
-subclass n- 0: 
no atomic line 
-other subclasses: 
strongest line 
[Ne II], 12.8 pm:n-l 
[S III], 18.8 pm:n-2 
[Ar III], 9.0 pm:«=3 
[S IV], 10.5 pm:/i-4 
[Ne III],15.5 pm:/z— 5 
[Ne V], 14.5 pm:n=6 

71 

compact HII regions and 
planetary nebulae 

9/i 

same as class 8, but 
without 1 1 .3 pm line 

50 

unknown 


V*- 


■ * >'*, 


IX- 15 



NORMAL STAR (18) ? c | CARBON RICH STAR (44) 



Figure IX.D.2.1 Representative low resolution spectra. Numbers refer to the spectral classes defined in the text 




STAR WITH SILICATE ABSORPTION (38) = J STAR WITH SILICATE ABSORPTION (79) 



Figure IX.D.2.2 Representative low resolution spectra. Numbers refer to the spectral classes defined in the text. 



PLANETARY NEBULA (95) s m PLANETARY NEBULA (96) 



Figure IX.D.2.3 Representative low resolution spectra. Numbers refer to the spectral classes defined in the text. 



Class 5 n objects showed no significant 10 pm excess or deficit. The subclass n is again equal to 
twice the spectral index between 14 and 22 pm. 

Sources with red spectra showing 10 pm silicate band emission were assigned to class 6 n, sources 
with silicate absorption were assigned to class In. In both cases the subclass n was derived from the 
strength of the emission or absorption with respect to the adjacent continuum according to (Eq. 
IX.D.2,3). 

P.2.d Spectra with Spectral Lines: Classes Sn,9n 

For sources with spectral lines the classification proceeded independently of the red or blue shape of 
the overall energy distribution. The classification processor searched down to the 4 o level for the seven 
spectral lines listed in Section IX.D.2.a. If the 1 1.3 pm line was among those found, then the source was 
assigned to class 8 n. If the 1 1.3 pm line was not present, then the source was assigned to class 9 n. The 
subclass n was based on which of the six remaining lines, in order of increasing excitation, was strongest 
(see Table IX.D.l). 

D.2.e Miscellaneous Spectra: Class On 

Spectra that did not fit into one of the above categories, or which were of lower quality, were 
assigned to class On. Class 01 spectra had such low signal-to-noise ratios that excesses or deficits greater 
than 0.5 (Eq. IX.D.2-4) were not significant; spectra more than 5 sigma away from the boundaries of the 
ln,2n,3n,4n spectra were given codes 00, 02, 03, and 04. These subclasses had only a few sources. Code 
05 was given to those objects with low signal-to-noise red spectra without significant lines or bands. 

D.3 Performance of the Classification Scheme 

About 90% of the sources in the catalog fall in the blue classes 1 n to 4 n. Figure IX.D.l shows the 
adequacy of the classification scheme, with most of the strong sources located in three well-defined areas. 

The featureless stars (K and early M spectral types) cluster in the center, and their spectral indices 
in both the short- and long-wavelength parts of the spectrum are close to -4, as expected. 

The class 2 n objects (oxygen-rich stars with 10 pm emission bands) lie in a narrow strip indicating 
that the shape of the 10 pm band is fairly constant. Moreover, the spectral indices increase systematically 
as the 10 pm band becomes stronger. The class 4 n objects (carbon-rich stars with 1 1 pm emission bands) 
are well separated from the class 2 n objects except in the lowest subclasses, where confusion between the 
two types occurs. 

The class 4 n objects occupy a less clearly defined area in Fig. IX.D.l because of variable structure 
shortward of the 1 1 pm band. A few objects with code 04 and in the class 4 n seem to be oxygen-nch 
stars where the 10 pm band is self-absorbed. For these sources, simple classification schemes break down. 

The class 3 n sources (oxygen-rich stars with 10 pm absorption) suffers from the problem that the 10 
pm band becomes so broad that it covers most of the short-wavelength half of the spectrum, leaving no 
part of the spectrum for an estimate of the continuum. The class 3 n crosses show a large scatter in Fig. 
IX.D.l. 

The red classification was tuned to detect as many emission-line sources as possible, and very liberal 
criteria were applied for significance. Consequently, the automatic classification program contaminates 


IX-19 


classes 8 n and 9 n with objects that belong to other categories. Deep 10 pm absorption bands combined 
with the instrumental short-wavelength band edge near 13.5 pm often produced spurious 12.8 pm lines; 
10 and II pm emission bands gave rise to features that resembled 11.3 pm lines. All line spectra were 
inspected by eye and the classification was changed where needed. 

For the strongest 300 sources, the individual spectra contributing to the coverage were classified 
independently by visual inspection. Apart from the expected problems around the boundaries between 
the various classes, the classifications were consistent within one subclass. 

Another test of the classification scheme was done on the line sources. Spectra were extracted for 
more than a hundred planetary nebulae. Roughly twenty had acceptable line spectra according to the cri- 
teria given in Section IX.C.4. The classification processor correctly described those sources. 

E. Some Characteristics of the Catalog 

E. 1 Completeness 

No sharp flux density limit can be specified above which completeness is strongly guaranteed, 
because of the stringent rejection criteria (see Section IX.C.3). In general, completeness is much better at 
high than at low Galactic latitudes since most of its selection criteria were meant to discriminate against 
confused sources. Figure IX.E. 1 shows plots of log (number) versus log (f v ) where / v is the flux density 
at 12 pm. The three most populated main classes have been plotted as well as the total sample. The 
total number follows a power law with a slope of *-1.3, while for sources with flux densities between 15 
and 100 Jy the slope of the power law is closer to -1.0. These slopes are consistent with a population of 
Galactic sources. The catalog is obviously incomplete below "15 Jy. The variation of slope with / v is 
rather large for classes 4 and 2, while the slope of the class 1 curve is steeper than any of the others and 
much more constant. These differences can be understood if the spectral features characterizing classes 2 
to 4 become unrecognizable to the classification program for fainter sources. Consequently, the fainter 
class 2-4 sources tend to migrate into class 1 at lower flux densities, thus causing an over-population of 
class 1 and an under-population of classes 2-4. 

E.2 Checks on the Shape of the Spectra 

The general shape of the spectra is largely determined by the wavelength- dependent responsivity 
correction, the deviation of which has been discussed in Section IX.B.2. If a considerable error were 
made in assuming that the infrared spectrum of a Tau was consistent with that of a 10000 K black body, 
this would show up as a systematic variation of the ratio of integrated spectral flux over 12 pm survey 
flux with spectral index. Some individual spectra observed from the ground have been compared with 
the IRAS spectra. The comparison showed satisfactory results. 

A word of caution is due with respect to line spectra. Relative line strength observed in different 
spectrum-halves (8-13 and 11-22 pm) of spectra with little or no continuum, may not be reliably cali- 
brated. This uncertainty is because of the uncertainty in the cross-scan dependent responsivity correction 
(see Sections IX.B.3 and C.2.e). Without the presence of a continuum in the overlap region of the two 
spectrum-halves, joining the spectrum-halves is uncertain and has generally not been done. Each 
spectrum-half may therefore have an uncertainty of up to 20%. 




Figure IX.E.l. Plot of log (number) versus log / v for sources seen with the spectrometer. The 
number plotted in the ordinate is the integrated number of sources with flux densities 
less than the designated flux density. The flux densities plotted in the absissca have 
been obtained from the main point source catalog. The bottom three curves 
correspond to individual spectral classes defined in Section IX.D. 


Authors: 

E.Raimond, D. A. Beintema, and F. M. Olnon. 


References 

Wildeman, K.J., Beintema, D.A. and Wesselius, P.R. 1983, Journal British Interplanetary Soc., 36, 21. 


IX-21 





X. THE FORMATS OF THE IRAS CATALOGS AND ATLASES 


A. Introduction 

This chapter describes the formats of the point source and small extended source catalogs, the total 
intensity images and the low-resolution spectra in their printed and machine-readable forms. A brief 
description is given of each entry in the catalogs; tables describe each column of the catalog in more 
detail and give, for the machine readable versions, the logical type of each variable and its length in bytes. 
Each tape containing IRAS catalog information has at least one header file containing the date and ver- 
sion number of the data on the tape (Table X.A. 1). 



Table X.A.l Format of Header Files 


Start 

Byte 

Name 

Description 

Length 

00 

NAME 

Name of IRAS 
data product 

30A1 

30 

DATE 

Date of production 

12A1 

42 

VERS 

Version Number 

5A1 

47 

COMMENT 

32 bytes of 
comment 

32 A 1 


IRAS images are presented in the FITS format (Wells, Greisen and Harten 1981). Sample headers 
are given for the sky plates, the Galactic plane map and the all-sky map. The format of the 1/2° beam 
zodiacal history file is described in this chapter. The format of the extragalactic specialty catalog is 
described at the beginning of that catalog. 

IRAS data products discussed in this chapter include: 

Point Sources: 

Machine-readable version of the Catalog (Section X.B.l) 

Printed version of the Catalog (Section X.B.2) 

Machine-readable version of the Working Survey Data Base (Section X.B.3) 

Small Extended Sources: 

Machine-readable version of the Catalog (Section X.C.l) 

Printed version of Catalog (Section X.C.2) 

Extended Emission Images: 

16.5° Images (photographic representation - Section X.D.3a) 

16.5° Images (machine readable format - Section X.D.3b) 


X-l 




Galactic Plane Images (all formats - Section X.D.4) 

Low-Resolution Maps (all formats - Section X.D.5) 

Zodiacal History File (machine readable - Section X.D.6) 

Low-Resolution IRAS Spectra: 

Catalog of spectra (machine readable - Section X.E) 

B. Point Sources 

The information about infrared point sources is presented in increasing detail, progressing from the 
printed volumes to the tape version of the catalog to the detailed description of the observational and 
processing history of each source in a file known as the Working Survey Data Base (WSDB) augmented 
by its ancillary file. The printed version (Section X.B.2) is intended for users at the telescope or at insti- 
tutions without computerized information retrieval systems. The catalog tape (Section X.B.l) is intended 
for astronomers desiring to make statistical studies and to search the catalog for large numbers of sources. 
The WSDB and ancillary file (Section X.B.3) are meant to give the sophisticated researcher all the avail- 
able data on any given source such as its brightness on each hours-confirmed sighting, the detectors 
involved and the details of the data reduction such as confusion with neighboring sources. 

Another catalog available only in machine readable form lists the WSDB entries for all sources that 
failed one or more of the confirmation and confusion criteria and were not, thus, included in the main 
catalog. This file of REJECTED sources includes spurious objects, including: processing failures, space 
debris, asteroids and comets, and celestial sources that, due to incompleteness at faint levels or to varia- 
bility, failed to achieve the minimum criterion of two hours-confirmed sightings. In regions of high 
source density the file includes sources rejected by the more severe criteria for reliability applied there 
(Section V.H.6). Caveat emptor. 

B.l The Machine Readable Version of the Point Source Catalog 

The point source catalog tape is divided into six individual files, covering the range 0 to 24 hr in 
right ascension in blocks of four hours each. Each file contains from 20,000 to 90,000 sources arranged 
in order of increasing right ascension. Each of the six catalog files is preceded by a file containing a single 
80-character ASCII record which lists the date and version number of the catalog. Thus to read the 
entire catalog one must read sequentially through 12 files, six containing the data and six containing 
dates and version numbers. 

Table X.B.l describes each entry in the catalog tape. Those columns that are also included in the 
printed version are marked. Each catalog entry requires 160 + NID x 40 bytes of ASCII data where 
NID is the number of positional associations for each source. In the table the column "Format" refers to 
the length and type of the (FORTRAN) character field used to read or write each entry. Figure X.B. 1 
describes the format of the printed version of the catalog. 

The tape is written with 80-character (ASCII) logical records and blocked with 256 logical records 
per physical record so that one can regard the tape as a sequence of card images. The entries are arranged 
so that the source data fits into two records. Association information requires an additional 40 characters 
per association and appears in subsequent records, two associations per record. If one assumes an average 


X-2 


I I n 



Table X.B.l Format of Point Source Catalog Tape 


Start 

Byte 


Name 


Description 


Units 


-new record- 


80 

RELUNC 1 

percent relative flux density 
uncertainties ( 1 value 
per band). 


92 

TSNR 

ten times the minimum signal- 
to-noise ratio in each band 


112 

CC 1 

point source correlation 
coefficient ( 1 value 
per band). 


116 

VAR 1 

percent Likelihood of 
Variability 


118 

DISC 

Discrepant Fluxes flag 
(one per band, hex-encoded) 


119 

CONFUSE 1 

Confusion flag (1 flag 
per band, hex-encoded) 


120 

PNEARH 1 

Number of nearby hours- 
confirmed point sources 


121 

PNEARW 1 

Number of nearby weeks- 
confirmed point sources 



Format 


00 

NAME 1 

Source Name 

— 

1 1 A 1 

11 

HOURS 

Right Ascension 1950. 

Hours 

12 

13 

MINUTE 

Right Ascension 1950. 

Minutes 

12 

15 

SECOND 1 

Right Ascension 1950. 

deci-Seconds 

13 

A1 

18 

DSIGN 1 

Declination Sign 

± 

19 

DECDEG 

Declination 1950. 

Arc Deg 

12 

21 

DECMIN 

Declination 1950. 

Arc Min 

12 

23 

DECSEC 1 

Declination 1950. 

Arc Sec 

12 
T 1 

25 

MAJOR 1 

Uncertainty ellipse 
major axis 

Arc Sec 

13 

28 

MINOR 1 

Uncertainty ellipse 
minor axis 

Arc Sec 

13 

31 

POSANG 1 

Uncertainty ellipse 
position angle 

Degree 
(East of 
North) 

13 

34 

NHCON 1 

Number of times observed 

— 

12 

36 

(<25) 

FLUX 1 

Averaged non-color corrected 
flux densities ( 1 value per 
band) 

Jansky 

(10 -26 W m -2 Hz -1 ) 

4E9.3 

411 

12 

2A1 

72 

FQUAL 1 

flux density quality. 
(1 value per band) 


76 

NLRS 

Number of significant LRS 
spectra 


78 

LRSCHAR 1 

Characterization of averaged 
LRS spectrum. 



413 

415 

4A1 

12 

A1 

A1 

II 

II 


X-3 



Table X.B.l Format of Point Source Catalog Tape (Continued) 


Start 


Byte 

Name 

Description 

Units 

Format 

122 

SESl 1 

Number of seconds-confirmed nearby 
small extended sources 
( 1 value per band) 

— 

411 

126 

SES2 1 

Number of nearby weeks-confirmed 
small extended sources 
(1 value per band) 


411 

130 

HSDFLAG 1 

Source is located in high 
source density bin 
( 1 value per band, hex-encoded) 


A 1 

131 

CIRRI' 

Number of nearby 100 pm 
only WSDB sources 

— 

ii 

132 

CIRR2 1 

Spatially filtered 100 pm 
sky brightness ratio to flux 
density of point source 
(see text) 


u 

133 

CIRR3 

Total 100 jim sky 
surface brightness 

MJy sr“ [ 

13 

136 

NID 1 

(<25) 

Number of positional 
associations 

— 

12 

138 

IDTYPE 

Type of Object 



11 

139 

SPARE* 

2 1 spare bytes* 

( new record ) 

— 

2 1 A 1 * 

160 

CATNO 1 

Catalog number 


12 

162 

SOURCE' 

Source ID 



15A 1 

177 

TYPE' 

Source Type/Spectral Class 



5AI 

182 

RADIUS' 

Radius Vector from IRAS 
Source to Association 

Arc Sec 

13 

185 

POS 

Position Angle from IRAS 
Source to Association 

Degree E 
of N 

13 

188 

FIELD 1 2 

object field #1 
(magnitude/other) 

catalog 

dependent 

14 

192 

FIELD2 2 

object field #2 
(magnitude/other) 

catalog 

dependent 

14 

196 

200 

-240 

etc. 

FIELD3 

object field #3 
(size/other) 

continuation of associations 3 
in blocks of 40 bytes 

catalog 

dependent 

14 


'Quantities listed in printed version of catalog. 

- FIELD 1 is listed in printed version of catalog, except for catalogs 2 and 19, where FIELD2 is list- 


CATNO, SOURCE, TYPE, RADIUS, POS, FIELD 1-3 are repeated in blocks of 40 bytes 2 per 
logical record, as necessary. The definition and formats of HELD 1-3 depend on the individual 
catalog in which the association is found. See Table X.B.4. 

*For Version 2.0 of the Catalog, see Table XII.A.3. 


X-4 


i i n 


iim 1 'i« ii i mu ii liffniiiiiipiiiiiiiiii vim fad'll 




ORIGINAL PAGE IS 
OE POOR QUALITY 


Right Ascension; 00 I ’22' T, 03' , -00 ,, 27 m 31 l 


Declination: A 60'- + 70” 


00220 + 6519 
00220 • 6556 
00222 + 6354 

00222 + 6952 

00223 + 6745 

00223 + 6026 

00224 ♦ 6421 
00224 + 6600 A 
00224 + 6101 
00227 + 6805 


Galactic Uncertainly 
a B Coords SMJ SMN 

to n i b r > n 


35 0 
59 16 
15 5 26 
17 5 14 
21 B 17 

23 2 12 

24 1 54 
26 6 11 
29.9 50 
439 19 


IWWIIUV 


120 + 3 

120 » 3 
120 + 1 

121 + 7 
120+ 5 
120 2 
120 ^ 2 
120 -* 4 



860000200C 

940020000C 

350130760C 

9500200036 14 

5500000008 

9600001418 

84001 1540C 

9500 100 IOC 

750000062C 

5601001008 


AsMK+alions 


0 T Name Tyr* 


VR063 00 01 
DO 23258 


11177 K2 
OCL 0291 


L MEANING OF FIELD VARIES 
WITH CATALOG (SEE TABLE 
X B.3). GENERALLY. 
OPTICAL BRIGHTNESS IN 

decimag 

- LENGTH OF VECTOR IN 
ARCSEC FROM IRA5 SOURCE 
TO ASSOCIATED SOURCE 

• NAME OF ASSOCIATED SOURCE AND 
CHARACTER OR SPECTRAL TYPE 


-CATALOG NUMBER (SEE TABLE V.H.1) 


ASSOCIATED 
‘ SOURCE 


L NUMBER of POSITIONAL ASSOCIATIONS IN OTHER ASTRONOMICAL 

CATALOGS OF WHICH ONE IS PRINTED. ASTERISK DENOTES SSC ASSOCIATION 

!_ SES2 FLAG INDICATES PRESENCE OF WEEKS - CONFIRMED SMALL EXTENDED SOURCES 
IN WINDOW (HEX ENCODED BY BAND) 


- LRSCHARACTERIZATION (SECTION IX. D) 


I L HIGH SOURCE DENSITY FLAGS (HEX ENCODED BY BAND) 

1 NUMBER OF SECONDS - CONFIRMED SMALL EXTENDED SOURCES 

IN WINDOW FOR 1 2,25,60. 100pm BANDS 

NUMBER OF WEEKS - CONFIRMED POINT SOURCES IN WINDOW 
NUMBER OF HOURS CONF IRMED POINT SOURCES IN WINDOW 
CONFUSION (HEX ENCODED BY BAND) 

CIRRUS 2 RATIO OF CIRRUS FLUX TO SOURCE FLUX (SEE V H.4) 
CIRRUS 1-NUMBER 100pm ONLY SOURCES IN WINDOW 


[ CONFUSION BLOCK 


VARIABILITY - (% PROBABI LITY SOURCE IS VAR1ABLEJ/10 
CORRELATION COEFFICIENTS OF SOURCE WITH TEMPLATE FOR 1 2,25,60, 100pm BANDS 
L-FLUX DENSITY UNCERTAINTIES WHERE (BLANK) 


FLUX DENSITY AND FLUX QUALITY FOR 25,60 AND E 

100pm BANDS (AS IN 12pm BAND, BELOW) < 

L i2wm FLUX QUALITY WHERE: (BLANK) = HIGH QUALITY j 

: - MODERATE QUALITY 1 

— 12pm FLUX DENSITY IN JANSKY5 L = UPPER LIMIT 

(NOT COLOR CORRECTED) S = SATURATED 

NUMBER OF HOURS CONFIRMED SIGHTINGS 

I SEM^WOR^AXt! (ARCSEC! I— CHARACTERISTICS OF POSITION UNCERTAINTY ELL.PSE 
— SEMIMAJOR AXIS (ARCSEC) 


NO FLUX UNCERTAINTY GIVEN (UPPER LIMIT FLUX) 
0% < UNC < 4% 

4 < UNC < 8 
8 < UNC <12 
12 < UNC <16 
16 < UNC <20 
UNC > 20 


-GALACTIC LATITUDE 
-GALACTIC LONGITUDE 


- POSITION OF SOURCE IN GALACTIC COORDINATES 


L DEC SECONDS 
RA SECONDS 


L LETTER TO DISTINGUISH DUPLICATE POSITIONAL NAMES 


- DEC MINUTES 

- OEC DEGREES 

- DEC SIGN 


| DECLINATION COORDINATE OF SOURCE 


- RA TENTHS OF MINUTES 
-RA MINUTES 

- RA HOURS 


-RIGHT ASCENSION COORDINATE OF SOURCE 


-IRAS SOURCE NAME BASED ON POSITION IN 1950 EQUATORIAL COORDINATES 


Figure X.B. 1 Explanation of format of printed version of point source catalog. 


X-5 



of two associations per source, then a file for one of the right ascension blocks will require about 10 
Mbytes. 

In general, for quantities that have a value in each wavelength band, subscripts or array indices 
range from 1 to 4 and refer, respectively, to 12, 25, 60 and 100 pm. A number of the flags discussed 
below have values in each of the four wavelength bands. For compactness these are encoded into a single 
base- 16 (Hex) digit (values 0-F) in the following manner (Table X.B.2). The four bits of that hex digit 
correspond to the four wavelength bands with bit 0 (Least Significant Bit) for 12 pm, bit 1 for 25 pm, bit 
2 for 60 pm and bit 3 for 100 pm. The presence of a flag in a band is denoted by setting its bit to 1. 
Thus a source with a particular flag, e.g. CONFUSE, set at 12 and 25 pm would have 
CONFUSE-0011-3(Hex) while one confused in 25,60 and 100 pm would have 
CONFUSE- 1 1 10-E(Hex). A flag encoded in this manner will be referred to as "hex-encoded by band". 

The remainder of this section discusses individual entries in the catalog. 


100 pm 
(Bit 3) 

Table X.B.2 Meaning of Hex Encoded Flags 

Flag Set in A Particular Band Resultant Value 

60 pm 25 pm 12 pm of Encoded Flag 

(Bit 2) (Bit 1) (Bit 0) xxxx - HEX-Decimal 

0 

0 

0 

0 

0000 

-0-0 

0 

0 

0 

1 

0001 

-1-1 

0 

0 

1 

0 

0010 

-2-2 

0 

0 

1 

1 

0011 

-3-3 

0 

1 

0 

0 

0100 

-4-4 

0 

1 

0 

1 

0101 

-5-5 

0 

1 

1 

0 

0110 

-6-6 

0 

1 

1 

1 

0111 

-7-7 

1 

0 

0 

0 

1000 

-8-8 

1 

0 

0 

1 

1001 

-9-9 

1 

0 

1 

0 

1010 

- A -10 

1 

0 

1 

1 

1011 

- B -11 

1 

1 

0 

0 

1100 

- C -12 

1 

1 

0 

1 

1 101 — D — 13 

1 

1 

1 

0 

1110 

- E -14 

1 

1 

1 

1 

mi 

- F -15 


Source Name: NAME 

The IRAS source name is derived from its position by combining the hours, minutes and tenths of 
minutes of right ascension and the sign, degrees and minutes of the declination. In obtaining the minutes 
of right ascension and declination for the name, the positions were truncated. The letters ’A\ ’B’, ’C’, 
etc. are appended to names of sources so close together that they would otherwise have had identical 
names. Names were uniquely assigned to both catalog and reject file sources, with catalog sources receiv- 
ing letters first. 


X-6 


ni i 



Potion- mOIJRS.MINUTE.SECOND.DSIGN.DECPEG,DE CMIN,DECSEQ 

Positions are given for the equinox 1950.0, and epoch 1983.5. Hours (HOURS) and minutes 
(MINUTE) of right ascension are given as integers while seconds (SECOND) are rounded to integer deci- 
seconds. The declination is given as a character sign (DSIGN) followed by integer values of degrees 
(DECDEG), minutes (DECMIN) and seconds (DECSEC). 

Position Uncertainty: MAJOR, M INOR, POSANG 

As discussed in Section VII.B.2, the uncertainty in the position for a source depends on its bright- 
ness in the various wavelength bands, its path across the focal plane and the number of sightings. The 
final uncertainty after position refinement is expressed as a 95% confidence uncertainty ellipse (see Sec- 
tion V.D.9) whose semi-major (MAJOR) and semi-minor (MINOR) axes are given in seconds of arc. 
The orientation (POSANG) of the ellipse on the sky is expressed in terms of the angle between the major 
axis of the ellipse and the local equatorial meridian. It is expressed in degrees east of north. 

Number of Sightings: NHCON 

The number of hours-confirmed sightings is given. This number of flux entries will be found in the 
WSDB. 

Flux Density: FLUX (4) 

Each of the four wavelengths has a non-color-corrected flux density in units of Janskys, (Uy - 
jq -26 w m -2 Hz -1 ). The quoted value is an average of all the hours-confirmed sightings as obtained by 
the prescription described in Section V.H.5. The quality of each flux density is designated by FQUAL 

(see below). ... fl 

The flux densities have been calculated assuming an intrinsic source energy distnbution such that the flux 

density / v is proportional to v -1 . Corrections to other spectral shapes can be made by consulting Section 
VI.C. 

The flux densities for sources so bright that they saturated the analog-to-digital converter on every 
sighting are lower limits based on the brightest value recorded. The uncertainties are given as ten times 
the quoted flux density and a flag is set indicating saturation has occurred (Table X.B.7a). 

Signal-to-Noise Ratio: TSNR(4) 

The signal-to-noise ratio given for an individual hours-confirmed sighting is the highest value of the 
detections comprising that sighting (Section V.C.2). The values quoted in the catalog are ten (10) times 
the minimum of the signal-to-noise ratios for the various sightings (HCONs) of the source. A value is 
given for each wavelength band with a high or moderate quality measurement and for those upper limits 
coming from a non-seconds-confirmed detection. Values of TSNR greater than 30,000 are given as 

30,000. 

Source Variability: VAR 

VAR is the percent probability (0-99) that a source is variable based on an analysis of the 12 and 25 
pm flux densities and their uncertainties (see Section V.H.5). The value M" indicates that the source was 
not examined for variability. 


X-7 


Discrepant Fluxes: DISQ4) 

The DISC flag indicates whether any one of the fluxes in a given band disagrees with others in that 
band, hex-encoded by band (Section V.H.5). 

Flux Density Quality: FQUAL(4) 

As described in Section V.H.5, a flux density measurement can be either high quality (FQUAL-3), 
moderate quality (FQUAL-2) or an upper limit (FQUAL=1). 

Low-Resolution Spectra: NLRS.LRSCHAR 

The Low-Resolution Spectrometer obtained 8-22 pm spectra of bright 12 and 25 pm sources 
(Chapter IX). NLRS gives the number of statistically meaningful spectra available for the source. 
LRSCHAR gives a short characterization of the nature of the spectrum (Table IX.D.l). All of the LRS 
spectra are available in tape and printed forms ( Astronomy and Astrophysics Supplement, 1985) as 
described below (Section X.E) and in Chapter IX. 

Flux Density Uncertainties: RELUNC (4) 

Each flux density measurement other than an upper limit has an associated uncertainty expressed as 
a 1 a value in units of 100 x 8 / Jf v . Uncertainties are discussed in Sections V.H.5 and VII.D.2. 

Point Source Correlation: CC (41 

As described in Section V.C.4, the point source correlation coefficient can have values between 87- 
100%. These are encoded as alphabetic characters with A=100, B-99...N-87, one value per band. The 
value quoted is for the highest correlation coefficient seen for that source on any sighting. 

Confusion: CONFUSE.PNEARH.PNEARW.HSDFI AO 

As described in Section V.D.2, a great deal of care went into trying to untangle instances of confu- 
sion between neighboring sources. In parts of the sky where the source density is low, confusion process- 
ing was sometimes able to separate sources that are quite close together. The CONFUSE flag is set if two 
or more sightings of the source in a given band had confusion status bits set indicating confusion in the 
seconds-confirmation or band-merging processes. This flag is hex-encoded by band. 

Other indicators of possible confusion are given by PNEARH and PNEARW which are, respec- 
tively, the numbers of hours-confirmed and weeks-confirmed point sources located within a 4.5' cross- 
scan and 6 in-scan (half-widths) window centered on the source. Values larger than 9 are given as 9. 

Regions of high source density received special processing to improve the reliability of the quoted 
sources (see Section V.H.6). The regions are band-dependent. If a particular band of a given source 
went through high source density processing, then the appropriate bit in HSDFLAG (Table X.B.7) is set. 
HSDFLAG is hex encoded by band. 

Small Extended Sources: SES1(4), SES2(4) 

SES 1 is the number of seconds-confirmed, small extended source detections in a given band found 
within a window centered on the source. The size of the window is 6' in-scan x 4.5' cross-scan (half- 
widths). As described in Sections V.H.3-4 and VII.H.l, values of SES1 greater than 1 should caution the 


X-8 


reader that significant extended structure may exist in the region and that the source in question may be 
a point-source like piece of a complex field. 

SES2 is the number of weeks-confirmed small extended sources in a given band, located within a 6 
in-scan x 4.5' cross-scan window (half-width) centered on the source. SES2 greater than 0 means that 
the point source flux measurement should be treated with caution as the source in question may, in fact, 
be extended. The flux quoted in the catalog of small extended sources may provide a better value for the 
source. 

Cirrus Indicators: CIRR 1 , CIRR2, CIRR3 

Over a large range of Galactic latitudes the infrared sky at 100 pm is characterized by emission 
from interstellar dust on a wide range of angular scales. The so-called "infrared cirrus" can seriously 
hamper efforts to extract point source detections from the data. To aid the user in interpreting the 
quoted 100 pm measurements three quantities have been established (Section V.H.4 and VTI.H). 

CIRRI gives the number of 100 pm-only WSDB sources located within a ±1/2° box in eclip- 
tic coordinates centered on the source. The sources included in this count are the weeks-confirmed 
sources prior to high source density region processing (if applicable) plus those sources hours-confirmed 
but not weeks-confirmed. Values of CIRRI greater than 3 may indicate contamination by cirrus 
with structure on the point source size scale. 

CIRR2 gives a cirrus indication on a larger scale than CIRRI and compares a "cirrus flux" with the 
source flux at 100 pm (see (Eq. V.H.2)). Values larger than 4*5 indicate the presence of considerable 
structure in the 100 pm emission on a 1/2° scale. A value of 0 indicates that no 1/2° data were available 
for the source in question. 

CIRR3 is the total surface brightness of the sky surrounding the source in a 1/2° beam at 100 pm, 
clipped to exclude values greater than 254 MJy sr _1 . Values of CIRR3 greater than 30 MJy sr" 1 indicate 
emission from dust with an appreciable column density. A value of CIRR3 “ -1 means that no data 
were available. 

Positional Associations: NID.IDTYPE.C ATNO.SOURCE,TYPE,RADIUS,POS, FIELD 1^ 

Much of the utility of the IRAS catalog comes from the association of infrared objects with sources 
known to exist from other astronomical catalogs. As described in Section V.H.9, a large number of cata- 
logs have been searched for positional matches. The total number of matches found is given by NID. 
Each match results in a forty-character description (2 per record). 

IDTYPE ranges from 1 to 4 and states whether an association was found in an extragalactic catalog 
(1), a stellar catalog (2), other catalogs (3), or matches in multiple types of catalogs (4). CATNO is the 
number of the catalog in which the match was found (Tables V.H.l, X.B.4). 

SOURCE is the name of the object in that catalog and TYPE its character or spectral type, if avail- 
able. A vector is drawn from the IRAS position to the associated object. RADIUS is the length of that 
vector in arc seconds. POS is the angle between the vector and the local equatorial meridian expressed in 
degrees east of north. Three fields (FIELD 1-3) have values depending on the catalog in question (Table 
X.B.4). Typically FIELD 1,2 are magnitudes (in decimag) and FIELD3 a size. 


X-9 


B.2 The Printed Version of the Point Source Catalog 

The printed version of the catalog is a strict subset of the information described in the preceding 
section. A number of fields have been abbreviated or deleted to make possible a single line entry for each 
source. The entries in the book are discussed below and shown as a figure in Fig. X.B. 1 

Name: (NAME) 

The full IRAS name is derived from the hours(HH), minutes(MM) and tenths of minutes (T) of 
right ascension and from the sign, degrees (DD) and minutes (MM) of declination. The right ascension 
and declination have been truncated. 

Position:(RA(s). DEC("b 

To conserve space, only the seconds of time for the right ascension and the arcseconds of declina- 
tion are given. To reconstruct the source position one must also take the hours and minutes of the right 
ascension and the degrees and minutes of declination from the source name. Because the source name 
was obtained by truncating rather than rounding the positions, values of RA(s) and DEC( ") as large as 
60 are possible. 

Galactic Coordinates 

Galactic coordinates (/ (//) , b l!n ) are given to a precision of 1°. 

Positional Uncertainties (SMAJ.SMIN.PA) 

Semi-major, semi-minor axes (”) and the position angle (°) of the uncertainty ellipse are given as 
described in the previous section. 

Number of Sig hti ngs (NH) 

The number of hours-confirmed sightings. 

Flux Densities 

Non-color-corrected flux densities are given in Janskys in the four bands. A single character follow- 
ing the measurement denotes the flux quality of the observation (Section V.H.5). A blank denotes a high 
quality measurement, a denotes a moderate quality measurement and an ’L’ denotes an upper limit. 
An S’ indicates that all measurements in that band were saturated and that the value listed is the largest 
of all the saturated values. 

Flux Density Uncertainties (FLUX/UNCS) 

The relative uncertainties are given in each band for each high or moderate quality measurement 
according to the following convention (where the uncertainty was first rounded to two significant figures): 


Symbol 

Uncertainty Range 

A 

0.00 < S/ v // v < 0.04 

B 

0.04 < 5/ v // v < 0.08 

C 

0.08 < 5/ v // v <0.12 

D 

0.12 < 5/ v // v < 0.16 

E 

0.16 < 5/ v // v < 0.20 

F 

5/y // v 2* 0.20 


X-10 




Correlation Coefficient (CORR/COEF) 

As described in the previous section, the correlation coefficient of the source with the point source 
template is given in each band according to the convention A— 100%, B— 99%, ...N— 87%. 

Variability (VAR) 

The probability (0-99%) that a source detected at 12 and 25 pm is variable is truncated to a single 
digit 0-9. A blank indicates that the source did not qualify for variability testing. 

The Confusion Block 

Ten flags or values each consisting of a single digit are combined into a block to denote the pres- 
ence of nearby sources of possible confusion. The flags are discussed in detail in the preceding section 
and in Section V.H and include: 

Cl - CIRRI - the number of 100 pm only point sources. 

C2 - CIRR2 - the ratio of 1/2° extended emission to the source flux. 

CF - CONFUSE - hex-encoded flag indicating bands in which confirmation processor found confu- 
sion. 

PH - PNEARH - number of nearby hours-confirmed sources. 

PW - PNEARW - number of nearby weeks-confirmed sources. 

S1-S4 - SES1 - number of nearby hours-confirmed small extended sources per band. 

HD - hex-encoded flag indicating which bands, if any, were processed according to high source 

density rules (Section V.H.6). 

Low-Resolution Spectra (LRS) 

The presence of a low-resolution spectrum is indicated by this two-digit classification of the spec- 
trum (Section IX. D). 

Small Extended Source (S2) 

The presence of one or more weeks-confirmed small extended sources is denoted by a hex-encoded 
flag denoting the bands in which a small extended source was found. 

The Associations Block (NID.CAT.NAME, TYPE, RAD, MAG ) 

When an IRAS source is found to have at least one positional association with objects in other 
astronomical catalogs, one such association is printed. Six pieces of information are given as described in 
the previous section and in Section V.H. 9. NID gives the total number of associations found in search- 
ing all catalogs. CAT is the number of the catalog (Tables V.H.l, X.B.4). The NAME and the TYPE 
(usually spectral or Hubble type) of the object are given. RAD is the distance from the IRAS source to 
the position of the associated object in arcsec. FIELD 1, which is usually a magnitude, is given in the 
MAG field, when available, for all catalogs except 2 and 19, for which FIELD2 is given. 


X-ll 


The association printed is chosen first by catalog within a catalog type (IDTYPE) as follows: 



Catalogs in Order of Printing Priority 


Printing Sequence 

IDTYPE 

Type 

i 

2 

3 

4 

5 

6 

7 

8 

9 

10 

11 

12 

13 

14 

15-31 

1 

extragalactic 

9 

6 

12 

10 

29 

25 

26 

27 

28 

30 

31 

- 

- 

- 


2 

stellar 

13 

4 

15 

2 

1 

7 

16 

17 

18 

19 

24 

- 

- 

- 

- 

3 

other 

14 

22 

21 

20 

23 

3 

5 

8 

11 

- 

- 

- 

- 

- 

- 

4 

multiple 

13 

9 

14 

1 

2 

3 

4 

1 

5 

6 

7 

8 

10 

i 

11 

12 

1 

15-31 


If more than one association was found in the catalog chosen by the priority scheme, the closest associ- 
ated source (smallest RAD) is printed. 

B.3 The Working Survey Data Base 

The most complete observational and data reduction history for point sources is contained within 
the Working Survey Data Base (WSDB). The WSDB is broken up into 20 files, one for each range of 
ecliptic longitudes called a Lune (see below). Each file is preceded by a header file containing a single 
80-character ASCII record giving the date and version of the WSDB (Table X.A.l). Except as noted 
below, the entries are similar to those listed above for the catalog tape. Table X.B.3a lists the entries, 
their variable type and length. A second file, called the Ancillary File (Table X.B.3b) contains additional 
flags and derived quantities obtained during final product generation. Most of the information in the 
Ancillary file is the same as that in the catalog tape described in Section X.B.l. Character variables are in 
ASCII format. All arithmetic quantities are in integer format with the high order byte first. 

It is important to note that the WSDB and Ancillary files are in binary format with variable length, 
blocked records (due to the variable number of hours-confirmed sightings or associations). Extra infor- 
mation has been added to the WSDB and Ancillary file tapes to make it possible to read in these variable 
length records. Each physical block on the tape begins with a Block Control Word of length 1*4. Its high 
order two bytes give the length of the physical block in bytes; this length includes the 4-byte length of the 
Block Control Word itself. In addition, each logical record within the physical block is preceded by a Seg- 
ment Control Word of length 1*4. Its high order two bytes give the length of the logical record in bytes; 
again, this length includes the 4-byte length of the Segment Control Word itself. The inclusion of this 
information should make it possible for computers than can work with Variable Block formats to read 
the tape easily. In no case does the information for a single source span two physical blocks. 

The remainder of this section describes WSDB and ancillary file entries not discussed above. 



SPAS Lune Number: LUNE 


In the data reduction the sky was divided into twenty "lunes" based on ecliptic 
coordinates. Lune 1 comprises that part of the sky with p>60°; Lune 2 comprises that sky with P<- 
60°. Lunes 3-20 comprise that part of the sky with ipi<60° and ecliptic longitudes, X, in 20° wide blocks. 
Lune 3 extends from 0°<X <20°, lune 4 from 20°<A,<40°. etc. 

Ecliptic Bin Number: BIN 

To aid the data reduction the entire sky was divided into some 40,000 1 sq. deg bins. The bin struc- 
ture is quite simple in ecliptic coordinates (Fig. X.Apl.l) and an algorithm for generating bin numbers is 
given in Appendix X. 1 . 

Ecliptic Coordinates: ELAT. ELONG 

Ecliptic coordinates are given in units of 10 -8 radians in the equinox 1950. 

Scan Angle: SCAN 

After weeks-confirmation SCAN gives the average scan angle of the focal plane with respect to the 
south-going local ecliptic meridian. 

Positional Uncertainty: SIGY, LZ, SIGZ 

As discussed in detail in Section V.D, the position refinement describes the source positional uncer- 
tainty in terms of 1 a in-scan and cross-scan gaussian uncertainties (SIGY and SIGZ) and a uniform 
uncertainty (of half-width LZ) whose size depends on the exact paths of the source across the focal plane. 

Number of LRS Extractions: LRSX 

Each time a source with a signal-to-noise ratio greater than 25 at either 12 or 25 pm transversed the 
focal plane, the data reduction software automatically triggered a request to extract a spectrum from the 
low-resolution spectrometer data (Chapter IX). The threshold was purposefully set quite low so that 
sources with weak continuua but strong lines could be detected. Many sources with LRSX>0 will fail to 
have meaningful spectra and will thus have NLRS=0. 

Known Source ID: KSID 

The positions and predicted brightnesses of some 32,000 point sources including SAO stars, IRC 
objects and asteroids were incorporated into the data reduction software to provide a check on the posi- 
tional and photometric accuracy of the IRAS sources. Table X.B.5 lists the range of KSID values 
assigned to sources of various types. 

The following values are given for each hours-confirmed sighting: 

Flux and Flux Uncertainty: FLUX(4), SIGF(4) 

Flux and flux uncertainty measurements are given in units of 10 -I6 W m -2 for each hours-confirmed 
sighting. Note that an instrumental flux, not the flux density, is given. In order to convert the instru- 
mental flux to flux density, one must divide the instrumental flux by 13.48, 5.16, 2.58, 1.00 x 10 12 Hz at 
12, 25, 60 and 100 pm respectively. The derivation of flux uncertainties and the averaging of the indivi- 
dual hours-confirmed fluxes to give the average value (see AVGFLUX, below) quoted in the catalog is 
discussed in Section V.H.5. 


X-13 



Signal-to-Noise Ratio: TSNR(4) 

If a source is detected (but not necessarily even seconds-confirmed) in a given band, then a value 
equal to ten times the maximum signal-to-noise ratio (SNR) observed on any detector in that band in the 
hours-confirmed sighting is retained. In relatively simple parts of the sky the noise estimator used to 
derive SNR gives reasonable values. In more complex regions near the Galactic plane (see Section V.C.2) 
the utility of SNR is very limited. 

Correlation Coefficient: CORR 

The maximum point source correlation coefficient (see Section V.C.4.) obtained during a hours- 
confirmed sighting is retained for each band. The values for each band, expressed as percentages up to a 
maximum of 100, are encoded into a single integer according to the following algorithm: 

CORR - CC(1) x 2 24 + CC(2) x 2 16 + CC(3) x 2 8 + CC(4) (X.B.l) 

Flux Status: FSTAT 

In each band there is a hierarchy of measurement quality depending on how many times a given 
source is observed within a given hours-confirmed sighting. FSTAT plays a crucial role in determining 
whether a source is included in the catalog at all, whether a flux is of high, medium or low quality and 
how the flux averaging was performed (Section V.H.5). For the meaning of each FSTAT value refer to 
Section V.D.8. 

The values of FSTAT for the four bands are compressed into a single integer according to the fol- 
lowing algorithm: 

FSTAT - FSTAT(l) x 2 12 + FSTAT {2) x 2 8 + (X.B.2) 

FSTAT( 3) x 2 4 + FSTAT( 4) 


Detector IP’s: DETID 

The identities of all detectors observing a given source on the sightings (up to a maximum of 3 
sightings) comprising a given hours-confirmed observation are recorded in the array DETID(I,J). The 
four bands run from 1-1 to 4 while the three sightings run from J— 1 to 3. 

Each value of DETID contains an integer that must be decoded according to the following algo- 
rithm to obtain the detectors observing the source: 

DETID(I,J ) = D 1 x2 10 + Z)2x2 5 + D3 (X.B.3) 

where detector numbers, Dl, D2 and D3, range from 1 to 16 within each band. Table X.B.6 lists the 
correspondence between detector number within a band to the true detector number. The order of detec- 
tor number is significant as described in V.D.5. It should be emphasized that when three detectors are 
named in a sighting, indicating the presence of an edge detection, the weakest of the edge detections may 
not have played any part in the assignment of flux or position. 


Detector Name: DNAM, TNAM 


Throughout the course of the data processing each hours-confirmed sighting is known by a combi- 
nation of the first detector measuring the object (DNAM) and the time (in deci-UTCS since 1981, Janu- 
ary 1, Oh UT) of that sighting (TNAM). 

Q rtnfnsinn Status: CSTAT. PNEARH. PNEARW 

At various stages in the reconstruction of a point source attempts are made to recognize (and 
remedy) the effects of confusion between nearby sources. The confusion status word CSTAT plays an 
important role in selecting sources to be treated by the "clean-up’ processor (Section V.H.2) and in deter- 
mining which sources to keep in regions of high source density (see Section V.H.6). For a detailed dis- 
cussion of confusion processing see Section V.D.2. Values of CSTAT and brief descriptions of their 
meaning are given in Section V.D.8. 

The values of CSTAT for each band are encoded into a single integer according to the following 
algorithm: 

CSTAT - CSTAT{\) x 2 24 + CSTAT(2) x 2 16 (X.B.4) 

+ STAT(3) x 2* + CSTAT(4) 

The number of hours- and weeks-confirmed point sources within 6' x 4.5' (half-widths) of the quoted 
source, PNEARH and PNEARW, are encoded into a single byte: 

PNEAR - PNEARW x 2 4 + PNEARH (X.B.5) 

Pin™ Flags: CIRRUS. CIRRI. CIRR2, CIRR3 

The three flags denoting the presence of extended 100 pm emission ("cirrus") as discussed in Section 
V.H.4, are encoded in two bytes according to the algorithm: 

CIRRUS - CIRR3 x2® + CIRR 1 x 2 4 + CIRRI (X.B.6) 

Values of CIRR2 - 0 and CIRR3 - -1 mean no data were available. 

Small Extended Source Flags: SES1, SES2 

The two small extended source flags for each band discussed in Section V.E.l are encoded into two 
integers according to the following algorithm: 

SES\ - SES 1(1) x 2 12 + SES\(2) x 2 8 

+ SES 1(3) x 2 4 + SES 1(4) 

and 

SES2 - SES2(\) x 2 12 + SES 2(2) x 2 8 
+ SES2(3) x 2 4 + SES 2(4) 


(X.B.7) 


(X.B.8) 


X-15 



Clean up Processors: CLEAN. BRIGHT, ACCEPT, HSDPROC MISC 


In creating the WSDB from the weeks-confirmed data, several processors were applied in order to 
fix various known problems (see Section V.H). These processors set various flags that allow the user to 
understand what processing occurred. The clean-up processor allowed the weeks-confirmation of various 
WSDB sources that did not previously have an opportunity to weeks-confirm for purely technical reasons 
or which were incorrectly split asunder during band merging. The flags set by that processor are 
described in Table X.B.7a., and include flags indicating saturated fluxes. Optical crosstalk from certain 
bright objects such as Saturn and IRC+10216 produced a few spurious sources. These were marked for 
deletion. The byte BRIGHT (Table X.B.7b) notes these cases and also contains flags denoting final 
acceptance (ACCEPT) or rejection of the source in each band. The flags from the High Source Density 
processor (HSDPROC, see Section V.H.6) are given in 2 bytes per band (Table X.B.7c). Miscellaneous 
flags are set in MISC (Table X.B.7d) and include the presence of flux discrepancies in each band, whether 
the in-scan positional uncertainties needed to be increased (Section VTI.C) and whether the source was 
accepted in the catalog. 


Start 

Byte 

Table X.B.3a. 

Name 

The Catalog Working Survey Data Base (WSDB) 

Description Units 

Type 

0* 

LUNE 

SDAS Lune number 



1*4 

4 

BIN 

Ecliptic Bin Number 

— 

1*4 

8 

ELONG 

Ecliptic longitude 1950. 

10 -8 rad 

1*4 

12 

ELAT 

Ecliptic latitude 1950. 

10 -8 rad 

1*4 

16 

SCAN 

average scan angle with 
respect to ecliptic 
meridian 

milli-rad 

1*2 

18 

SIGY 

In-scan Gaussian 
position uncertainty 

p-rad 

1*2 

20 

LZ 

Cross-scan Uniform 
position uncertainty 

ji-rad 

1*2 

22 

SIGZ 

Cross-scan Gaussian 
position uncertainty 

p-rad 

1*2 

24 

LRSX 

Total number of LRS 
extraction requests 

" 

1*2 

26 

KSID 

Known Source ID 

— 

1*2 

28 NHCON Number of Hours- 

(<25) confirmed sightings 

The following values repeat for each hours-confirmed sighting: 


1*4 

32 

FLUX 

FLUX(I)-In-band power 

1(T 16 W m -2 

41*4 

48 

SIGF 

SIGF(I)=U ncertainty 
in FLUX(I) 

1(T ,6 W m -2 

41*4 

64 

TSNR 

TSNR(I)= 1 0x(max SNR) 

— 

41*2 

72 

CORR 

Maximum correlation 
coefficient ( 1 value 
per band, Eq. X.B.l) 


1*4 

76 

FSTAT 

Flux status 
word ( 1 value per 
band, Eq.X.B,2) 


1*2 

78 

DETID 

Detector ID Array 
(4,3) for 4 bands 
(Eq.X.B.3) 


121*2 

102 

LRSXNO 

Number of LRS extraction 
requests 


1*1 

103 

DNAM 

detector part 
of source name 

~ 

1*1 

104 

TNAM 

deci-UTCS part 
of source name 

deci-sec 

1*4 

108 

112 - 

192- 

etc... 

CSTAT 

-191 

-271 

Confusion status flags 
(Eq.X.B.4) 

repeats bytes 32-111 
for subsequent hours 
confirmed sightings. 


1*4 


* See X.B.3 for discussion of block and segment control words 


X-17 



Start 

Byte 

Name 

Table X.B.3b. Ancillary WSDB File 

Description 

Units 

Type 

O 1 

PNEAR 

Nearby hours- and 

weeks-confirmed 

neighbors 

— 

1*1 

1 

CLEAN 

Clean up Processor 
flags 

— 

1*1 

2 

SES1 

Number of nearby 
unconfirmed SES 
(4 bands encoded) 


1*2 

4 

SES2 

Number of nearby 
weeks-confirmed SES 
(4 bands encoded) 


1*2 

6 

CIRRI, 

CIRR2 

CIRR3 

Number of 100 pm 
only WSDB sources, 
spatially filtered 
100 pm emission, 
value of 100 pm half- 
degree beam total 
intensity 

MJy sr _I 

1*2 

8 

AVGFLUX 

AVGFLUX(I) Averaged flux 
(in-band power) 
in band I. 

10~ 16 W m -2 

41*4 

24 

AVGUNC 

AVGUNC(I) Uncertainty in 
averaged (in-band) flux 
in band I. 

10" 16 W m" 2 

41*4 

40 

HSDPROC 

HSDPROC(I). Rags set by 
the high source density 
processor in band I 

” 

41*2 

48 

RA 

Right ascension 1950 

IQ-5" 

1*4 

52 

DEC 

Declination 

l 0 -5» 

1*4 

56 

NAME 

source name 

— 

A* 12 

68 

NLRS 

Number of meaningful 
LRS spectra 

— 

1*2 

70 

LRSCHAR 

characterization of 
LRS spectra 

— 

A*2 

72 

BRIGHT 

ACCEPT 

Bright Source Clean up 
Flag for catalog sources 
accepted in the catalog 


1*1 

73 

VAR 

Percent likelihood of 
Variability 

— 

1*1 

74 

FQUAL 

FLUX Quality flags 
(one per band) 

BIT 0-1, 12 pm 
BIT 2-3, 25 pm, etc. 


1*1 




Table X.B.3b. Ancillary WSDB File (Continued) 


Start 

Byte 


Name 


Description 


Units 


Type 


1*1 


75 MISC 


76 

LUNE 

80 

BIN 

84 

ELONG 

88 

ELAT 

92 

NID 2 


(<25) 

94 

IDTYPE 

96 

CATNO 

98 

SOURCE 

113 

TYPE 

118 

RADIUS 

120 

POS 

122 

FIELD 1 

124 

FIELD2 

126 

FTELD3 

128 


-160 



Miscellaneous Status 
bits, incl. DISC 
ACCEPT, SIGY 
See Table X.B.7d 
SDAS Lune number 
Ecliptic Bin Number 
Ecliptic longitude 
Ecliptic Latitude 
Number of Associations 
to follow 

Type of association 
Catalog Number 
Source ID 

Source Type/Spectral Class 
Position difference 
Position Angle 
object field # 1 
object field #2 
object field #3 
continuation of association 3 
in blocks of 32 bytes 


_ 1*4 

... 1*4 

10 _8 rad 1*4 

10" 8 rad 1*4 

1*2 

1*2 
1*2 
A* 15 
A* 5 

(") 1+2 

' E of N 1*2 

... 1*2 

... 1*2 

._ 1*2 


in blocks of 32 bytes, 2 per logical 
on the individual catalog in which 


I etc. 

1 See X.B.3 for discussion of block and segment control words 

2 If NID— 0, one blank association field of 32 bytes is written. 

3 CATNO, SOURCE, TYPE, RADIUS, POS, FIELD 1 -3 are repeated 
record, as necessary. The definition and formats of HELD 1-3 depend 
the association is found. See Table X.B.4. 


X-19 



Catalog 

Table X.B.4 Meaning of the Source Association Fields 

Field and Meaning* 

1 GCVS 

Type 

Blank 


Field 1 - 

Code gives meaning for Fields 2-3 


if Field 1 - 1 

Field2 and Field3 are B mag [decimag] at max, min 


- 2 

Field2 and Field3 are V mag [decimag] at max,min 


- 3 

Field2 and Field3 are photographic mag [decimag] at max, min 


- 4 

Field2 and Field3 are estimated V mag [decimag] at max, min 


- 5 

Field2 is 999 and Field3 is 0 

2 Dearborn 

Type 

Blank 

Obs. 


Field 1 

Code for Field2 (1,2) 


Field2 

if Field 1 is 1, Field2 is red magnitude [decimag] * 
if Field 1 is 2, Field2 is 999 


Field3 

0 

3 Revised 

Type 

Blank 

AFGL 


Field 1 

Magnitude at 4.2 jim [decimag] 


Field2 

Magnitude at 1 1 jim [decimag] 


Field3 

Magnitude at 27 jim [decimag] 

4 2 -jim Sky 

Type 

Blank 

Survey 


Field 1 

K magnitude [decimag] 


Field2 

I magnitude [decimag] 


Field3 

0 

5 Globules 

Type 

Blank 

(Wesselius) 


Field 1 

999 


Field2 

Minimum diameter [arcsec] 


Field3 

Maximum diameter [arcsec] 

6 R C 2 

Type 

Blank 


Field 1 

Harvard V magnitude [decimag] 


Field2 

Bj [decimag] 


Field3 

D 0 [arcsec] 

7 Stars with 

Type 

Blank 

em. lines 


Field 1 

V magnitude [decimag] 


Field2 

999 


Field3 

0 


X-20 




Table X.B.4 

Catalog 

Meaning of the Source Association Fields (Continued) 

Field and Meaning * 


8 Equatorial 

Type 

Blank 


IR Cat. 





Field 1 

Flux density |10 — 16 W cm -2 |im -1 ] at 2.7 pm 



Field2 

999 



Field3 

0 


9 UGC 

Type 

Blank 



Field 1 

Zwicky magnitude [decimag] 



Field2 

Minimum diameter [arcsec] in B 



Field3 

Maximum diameter [arcsec] in B 


10 MCG 

Type 

Blank 



Field 1 

999 



Field2 

Minimum diameter [arcsec] in B 



Field3 

Maximum diameter [arcsec] in B 


11 Strasbourg 

Type 

Blank 


Planetary Nebulae 





Field 1 

V magnitude of Nebula [decimag] 



Field2 

B magnitude of Central Star [decimag] 



Field3 

Minimum diameter of Nebula [arcsec] 


12 Zwicky 

Type 

Blank 



Field 1 

Zwicky magnitude [decimag] 



Field2 

999 



Field3 

0 


13 SAO 

Type 

Spectral Type 



Field 1 

V magnitude [decimag] 



Field2 

p g magnitude [decimag] 



Field3 

0 


14 ESO/ 

Type 

First 3 characters of object type 


UPPSALA 





Field 1 

B magnitude [decimag] 



Field2 

Maximum diameter [arcsec] 



Field3 

Minimum diameter [arcsec] 


15 Bright 

Type 

Spectral Type 


Stars 





Field 1 

V magnitude [decimag] 



Field2 

B-V [centimag] 



Field 3 

U-B [centimag] 



X-21 




Table X.B.4 Meaning of the Source Association Fields (Continued) 


Catalog Field and Meaning * 


16 Suspected 

Type 

Spectral Information 


Var. 

Field 1 

V magnitude at maximum [decimag] 



Field2 

999 



Field3 

0 


17 Carbon 

Type 

Spectral Type (May be truncated) 


Stars 

Field 1 

p g magnitude [decimag] 



Field2 

V magnitude [decimag] 



Field3 

I magnitude [decimag] 


18 Gliese 

Type 

Spectral Type (May be truncated) 



Field 1 

V magnitude [decimag] 



Field2 

B-V magnitude [millimag] 



Field3 

U-B magnitude [millimag] 


19 S Stars 

Type 

Blank 



Field 1 

p g magnitude [decimag] 



Field2 

V magnitude [decimag] 



Field3 

I magnitude [decimag] 


20 Parkes HII 

Type 

Blank 


Survey 

Field 1 

999 



Field2 

Minimum diameter [arc sec] 



Field3 

Maximum diameter [arcsec] 


21 Bonn HII 

Type 

Blank 


Survey 

Field 1 

Flux density at 4.875 GHz (Jy) 



Field2 

Diameter [arcsec] 



Field3 

0 


22 Blitz 

Type 

Blank 



Field 1 

Diameter [arcsec] 



Field2 

V co [Km/s] 



Field3 

Peak T a [°K] 


23 OSU 

Type 

Blank 



Field 1 

999 



Field2 

999 



Field3 

Diameter [arcsec] 




X-22 



Table X.B.4 Meaning of the Source Association Fields (Continued) 

Catalog 

Field and Meaning 

* 

24 IRC 

Type 

C if 2.2 |im sources are possibly confused, blank otherwise 

w/good pos. 




Field 1 

Right ascension difference (IRC-IRAS) [deciseconds of time] 


Field2 

Declination difference (IRC-IRAS) [seconds of arc] 


Field3 

0 

25 DDO 

Type 

Blank 


Field 1 

999 


Field2 

999 


Field3 

0 

26 Arp 

Type 

Blank 


Field 1 

999 


Field2 

999 


Field3 

0 

27 Markarian 

Type 

Blank 


Field 1 

999 


Field2 

999 


Field3 

0 

28 Strong 

Type 

Object type (GAL or QSO) 

5 GHz 




Field 1 

V magnitude [decimag] 


Field2 

5 GHz flux density [deciJyl 


Field3 

0 

29 Veron- 

Type 

Object classification 

Veron 




Field 1 

V magnitude [decimag] 


Field2 

Redshift x 1000 


Field3 

0 


X-23 




Table X.B.4 

Catalog 

Meaning of the Source Association Fields (Continued) 

Field and Meaning* 

30 Zwicky 

Type 

Blank 

8 Lists 


Field 1 

999 


Field2 

999 


Field3 

0 

31 VV 

Type 

Blank 


Field 1 
Field2 

Special flag, (see below) 999 otherwise 
999 


Field3 

0 

32 IRAS Small 

Type 

Blank 

Scale Structure 


Field 1 

Hex coded bands 


Field2 

Blank 


Field3 

Blank 

39^ OSU Radio 

Type 

Blank 


Field 1 

Frequency 


Field2 

Flux (deci Jy) 


Field3 

0 

40 Michigan 

Type 

Class** 

Spectral 


Field 1 

Mag (decimag) 


Field2 

HD number (low byte) 


Field3 

HD number (high byte) 

41 Serendipitous 

Type 

Blank 

1 Survey 


Field 1 

Hex coded SSC bands 


Field2 

First SSC flux (m Jy) 


Field3 

Second SSC flux 




In the printed version FIELD 1 is listed if present except for catalogs 2 and 19, where FIELD2 is 


**Source name and type fields were combined to hold spectral type and class. 
Catalog numbers 33-38 reserved for internal use. 

VV Catalog Flags (Catalog 31) 


FIELD 1 Explanation 

10 vv c 10 has the sam e coordinates as VV 29 in the W Atlas. The UGC was used to 
confirm that the coordinate is correct for W 29 and erroneous for VV 10. The UGC 
position for W 10 = UGC 10814 was adopted, 

1 1 The VV Position is substantially different (>400") from positions for the object in other 
catalogs. The VV position has been assumed to be in error because two or more other 
catalogs agree on a different position. The UGC position has been adopted. 

*2 Same as for 1 1, but the OSU position has been adopted. 

13 The Position in the VV Atlas, and the position listed for the W object in the OSU are in 
disagreement. The true position has been established to be close to that of the OSU bv 
the use of overlay transparencies on the POSS. The OSU position has been adopted. ' 

1 4 Same as for 1 3 > but the osu position is not very good either, so a new position has been 
measured (accurate to about 1 ). 


X-24 



Table X.B.5 Known Source ID’s 

Range 

Source Type 

1-465 

Selected AFGL catalog sources 

466-3157 

Two Micron Sky Survey sources 
without SAO counterparts 

3158-26979 

Selected SAO stars (mostly M 
and K stars). 

26980-27497 

Selected objects for Low-Resolution 
Spectrometer 

>30000 

Solar System Objects (asteroids, 
comets and outer planets) 


Table X.B.6 Detector Number 

In Band True In Band True 

Detector Detector Number Detector Detector Number 


Number 

12 pm 

25 pm 

60 pm 

100 pm 

Number 

12 pm 

25 pm 

60 pm 

100 pm 

1 

23 

16 

08 

01 

9 

47 

40 

31 

56 

2 

24 

17 

09 

02 

10 

48 

41 

32 

57 

3 

25 

18 

10 

03 

11 

49 

42 

33 

58 

4 

26 

19 

11 

04 

12 

50 

43 

34 

59 

5 

27 

20 

12 

05 

13 

51 

44 

35 

60 

6 

28 

21 

13 

06 

14 

52 

45 

36 

61 

7 

29 

22 

14 

07 

15 

53 

46 

37 

62 

8 

30 

39 

15 

55 

16 

54 

00 

38 

00 


X-25 






Table X.B.7a. CLEAN Bit Assignment 


Bit# 

Meaning 

0-3 

Bits 0-3 denote a flux limit derived 
solely from saturated detections. 
Bit 0—12 |rm. Bit 1—25 pm. 

4 

Source failed to weeks-confirm with 
another WSDB source in the mini-survey 
region. 

5 

Source resulted from weeks-confirming 
at least two WSDB sources in the mini-survey 


region. 

6 

Source failed to weeks-confirm with 
another WSDB entry observed within 
36 hours. 

7 

Source resulted from weeks-confirmation 
of two separate WSDB sources observed 
within 36 hours of each other. 


Bit# 

Table X.B.7b. BRIGHT/ACCEPT Bit Assignment 

Meaning 

0 

False source generated by nearby bright 


source 

1-4 

Source satisfies single 


band acceptance criteria 


BIT 1-12 |im, Bit 2-25 pm, etc. 

5-7 

Source satisfies adjacent 


band acceptance criteria 


BIT 5-12 and 25 pm 


BIT 6-25 and 60 pm 


BIT 7-60 and 100 pm 


X-26 


I II 






Table X.B.7c. HSDPROC High Source Density Processor Flags 

Bit # Meaning 


Byte 1 

0 and 1 00-0— not processed 

01 — 1 — low quality flux 

10- 2 — medium quality 

1 1- 3— high quality flux 

2 Band rejected 

3 Band accepted 

4-7 final reason for rejection 

0000-0- not rejected 
000 1 - 1 - not weeks-confirmed 

0010- 2- bad flux status 

001 1- 3- bad correlation coefficients 

0100- 4- bad confusion status 

0101- 5- inconsistent fluxes 

01 10- 6- weaker neighbor 

01 1 1- 7- confused neighbor 
1000-8- merging problems 

9-15- spare 

Byte 2 (Faults with Source) 

0 not weeks-confirmed 

1 bad flux status 

2 bad correlation coefficient 

3 bad confusion status 

4 inconsistent fluxes 

5 weaker neighbor 

6 confused neighbor 

7 merging problems 


X-27 




Bit # 

Table X.B.7d. MISC Bit Assignment 

Meaning 

0-3 

Discrepant Flux found in band. 


BIT 0, 12 pm, 


BIT 1 , 25 jim, etc. 

4,5 

SPARE 

6 

Oy fix (Section VII.C) 

7 

Source accepted in catalog 




C. The Small Extended Source Catalog 

The formats of the Small Extended Source Catalog (also known 
ture Catalog) and a description of the final generation of that catalog 
this series. 


as the Small Scale Slruc- 
are given in Volume 7 of 


X-29 



D. Extended Emission 
D. 1 Introductory Comments 

The extended emission data are presented as maps of the infrared sky at two different scales in all 
four wavelength bands. The entire sky surveyed is mapped with 2’ pixels at an effective resolution of 4’- 
6’ in 212 16.5° x 16.5° fields and also with an effective resolution of 1° in a single field. A special map of 
the region within ±10° of the Galactic plane was also made at 4’-6’ resolution. All of these maps are 
available in both digital and photographic formats. In addition, data from the survey averaged in a 1/2 ° 
x 1/2 ° beam is available in time ordered form. The details of the presentation of each of these products 
are described below. The methods used to produce the products were described in Section V.G. All the 
map projections used are described below. 

The survey covered most parts of the sky several times and the extended emission data have been 
separated for the reasons outlined in Section V.G.l. The data released in Nov. 1984 includes the 188 
fields in the 75% of the sky covered in the third sky coverage, along with the associated Galactic plane 
and low-resolution all-sky maps. Subsequent releases include a small part of the sky covered by the 
"minisurvey" done at the beginning of the mission (Section III.C.ll) and the first and second coverages, 
each of about 95% of the sky. 

D.2 Map Projections and Transformation Equations 

Three separate map projections were used. The descriptions below provide enough information for 
the user of the IRAS data. More information can be found Richardus and Adler (1972). 

D.2.a Gnomonic Projection - 1 6,5 ° Images 

The gnonomic projection used for 16.5° images produces a geometric projection of the celestial 
sphere onto a tangent plane from a center of projection at the center of the sphere. Each individual field 
has its own tangent projection plane with the tangent point at the center of the field. This projection is 
neither conformal (angle preserving) nor equivalent (equal area) but does have the property that all 
straight lines in the projection are great circles on the sphere. All projections were done so that the sky 
coordinate associated with a pixel refers to the position at the center of the pixel. For the 16.5° fields the 
maximum distortion of angles is 0.6° and the maximum distortion of area is 6%. The area distortion is 
approximately proportional to the inverse cube of the cosine of the angular displacement from the center 
of the field. The distortion is in the sense to make extended areas cover more sq. arcminute pixels than 
their true solid angles would require. This results in over estimating fluxes when integrating sources 
within fixed intensity contours. 

The transformation equations for conversion between line and sample number in the map and right 
ascension and declination on the sky are shown below. 

Forward: 

define 

scale - 30 pixels/degree 

A — cos(8) x cos(a — ao) 

F - scale x (180/jt)/[sin(8 0 ) x sin(8) + A x cos(8 0 )] (X.D.l) 


X-30 



then 


LINE - -F x [cos(5o) x sin(5) -Ax sin(5o)l 
SAMPLE F x cos(5) x sin(a - do) 


(X.D.2) 


Reverse: 

define 

X - SAMPLE /{scale x 180/n) 

Y - LINE /{scale x t 80 /tc) 

A — arctan [(X 2 + T 2 ) l/2 1 
P — arctan{—X/Y) 

XX - sin(8 0 ) x sin(A) x cos(P) + cos(5 0 ) x cos(A) 

YY - sin(A) x sin(P) 

then 

5 ~ arcsin [sin(8o) x cos(A) — cos(5o) x sin(A) x cos(P)l 

a - ao + arctan{YY/XX) ^ 

where S„ and no are the declination and right ascension of the held center. The arctangent functions for 
P and a must be four quadrant arctangents. 


D-2.h Equivalent Cylindrical Projection - Galactic Plan e Fields 

The Lambert normal equivalent cylindrical projection was used to provide an equal area projection 
of the sky within 10° of the Galactic plane. The projection cylinder is tangent to the celestial sphere at 
the Galactic equator and the projection proceeds by projecting radially outward from each point on the 
polar axis of the Galactic coordinate system in a plane parallel to the equatorial plane. The maximum 
angular distortion (deviation of bearing) is 0.9°. The equal area property of the transformation preserves 
photometric accuracy when integrating fluxes for an extended source. 

The transformation equations are: 

Forward: 

define 

scale - 30 pixels/degree 

then 

LINE scale x 180/x x sin {b 11 ) 

SAMPLE - -scale x (/" - Iff) (X.D.4) 


X-31 


I 


Reverse: 

l" - l(f - SAMPLE /scale 

b" - -arcsin [LINE /{scale x 180/7t)] (X.D.5) 

where b 11 and / 11 are Galactic latitude and longitude and subscript zero denotes the field center. 

D.2.c Aitoff Projection - Low-Resolution All-Skv Maos 

The Aitoff equal area projection was used to provide a photometrically correct map of the entire 
celestial sphere. Galactic coordinates were chosen as a convenient and natural coordinate system. The 
transformation equations are: 

Forward: 

define 

scale - 2 pixels/degree 

p * arccoslcos(b u ) x cos({/^ - 1$}/ 2)] 

0 - arcsin \cos(b H ) x sin({/" - /^}/2)/sin(p)] (X.D.6) 

then 

SAMPLE - -4 x scale x 180 /tt x sin(p/2) x sin(0) 

LINE “ ±2 x scale x 1 80 /tt x sin(p/2) x cos(0) (X D 7) 

where the ’+’ applies to b 11 <0 and the to b 11 >= 0. 

Reverse: 

define 

y - -LINE /(2 x scale x 1 80/7t) 

X - -SAMPLE /( 2 x scale x 1 80/rt) 

A - (4 - X 2 - 4 x Y 2 )' h 

then 


b" ” 180/7t x arcsin (A x Y) 

t 11 ~ Iff + 2 x 180/7t x arcsin [A x X /(2 x cost// 7 })] (X.D.8) 

where l" and b" are Galactic longitude and latitude and subscript zero denotes the field center. 

D.3 16.5° Images 

The 16.5° images the high-resolution presentation of the IRAS sky survey in image form. Two 
hundred twelve (212) 16.5 x 16.5° fields cover the whole sky with field centers spaced by approximately 
15 . The three sky coverages of the full mission are presented as three separate sets of 212 maps, with 


i 1 1 


X-32 


ii iii> h mill mu ji 


some maps not included in the third coverage (Table X.D.l). A list of plate numbers vs. plate centers 
comprises Appendix X.2 and a map of plate locations is given in Fig. X.D.l. An individual map 
consists of 499x499 array of 2'x2' pixels into which the IRAS survey data were mapped using a 
gnomonic projection (described in Section X.D.2.a). The 16.5* images are available as photographic 
prints and as digital magnetic tape. The formats of these two forms are described below. Detailed 
descriptions of the procedures used to produce the maps are given in Section V.G. 


Table X.D.l Plates Missing From the Third Sky Coverage 


47 95 130 

60 96 131 

70 97 132 

71 107 143 

72 108 144 

83 118 154 

84 119 155 

04 120 168 


D.3.a Prints of 16.5° Images 

Photographic black and white negative transparencies were produced from the digital map data with 
a film recorder. All four bands of each field were reproduced side by side in a rectangular format approx- 
imately 5 inches sq. intended for enlargement to 16 inches by 20 inches. Two hundred fifty-six (256) 
brightness levels were available with the film recorder and the brightness range in each band of each field 
was individually compressed, clipped and scaled to fit within these 256 levels. 

The compression, clipping and scaling were accomplished by first extracting the fifth root of the sur- 
face brightness to compress the dynamic range of the data. A histogram of the fifth root map was made 
and the pixel values were shifted and scaled into the 0-255 range so as to saturate the lower and upper 
one percent of the histogram. The approximate surface brightness value of any pixel can be recovered by 
first comparing the density of the pixel of interest to the grey scales which show every 17th pixel value 
from zero to 255. The shift and scale are removed using the 0 DN - X Jy sr -1 and 255 DN - Y Jy sr 1 
information in the label (see the sample label below). 

The complete formula is: 

surface brightness (Jy sr -1 ) = [- — ■— (T 1/5 — X l/5 ) + T 1/5 ] 5 (X.D.9) 

where D is the pixel value determined by comparison with the grey scale and Y and X are obtained from 
the label. 

Final calibration factors were not applied to the third sky coverage data before production of the 
photo products. Therefore, the intensities obtained using the procedure above on the third coverage may 
not agree with intensities found on the magnetic tape versions of the maps. The magnetic tape version 
uses the final IRAS catalog calibration and should be used in case of discrepancy. Intensities obtained 


X-33 


ORIGINAL PAGE IS 
OF POOR QUALITY 



i l r 


from the third coverage photo products can be approximately corrected to the final calibration by multi- 
plying the intensity derived from a photograph by the following factors: 


12 pm 

0.84 

25 pm 

0.80 

60 pm 

0.80 

100 pm 

0.69 


The first, second and mini-survey coverage photos were corrected to the final calibration. 

The coordinate system of each map is arranged so that when viewed with the printing in the label 
right side up north is up and east is to the left at the field center, which is adopted as pixel (0,0). In this 
orientation the horizontal rows of pixels are by convention called lines and the pixels within each line are 
called samples. Line numbers increase from top to bottom and sample numbers increase from left to 
right. The line numbers of the top and bottom lines are given in the label as TOP and BOTTOM, 
respectively. Similarly the left and right extreme sample numbers are given as LEFT and RIGHT. With 
this information and the tic marks along the sides of the image area the line and sample coordinates of 
any pixel can be determined for application of the inverse map projection formulae given above in Sec- 
tion X.D.2. The tic marks also allow alignment of the co-ordinate overlay grids as described in Sec- 
tion X.D.7. 

Color composite negative transparencies in 4x5 inch format of each sky plate field have been pro- 
duced by recording the 100 pm map in red (positive), the 60 pm map in green and the 12 pm map in 
blue. These color versions of the data are not intended for quantitative analysis. The shift and scale 
information in the label is difficult to read and no attempt has been made to produce a consistent color 
balance among the plates. In one plate a particular hue will indicate one ratio to 100 pm to 60 pm to 12 
pm brightness and in another plate the ratio for that hue will be somewhat different. 


Sample Label from 16.5° Image Photograph 


IRAS SKYFLUX HCON: 3 HELD: 153 DEC-30 RA: 10:00 25 MICRON 

R.A. & DEC GNOMONIC PIXEL: 2.00 ARCMIN JD:2445654.25-2445660.25 
TOP:-249 BOTTOM: 249 LEFT:-249 RIGHT: 249 1 TIC - 5 PIXELS 

0DN- 4.28E+ 2 JY/SR 255DN- 5.38E+ 2 MJY/SR 5TH ROOT DATE:84/08/10 


First line: 
SKYFLUX: 

HCON: 

HELD: 

DEC,RA: 

MICRON: 


refers to the 212 maps which cover the whole sky with 2' pixels. Can also be ALL SKY 
for the 1/2° pixel all-sky maps. 

serial number of the sky coverage 

IRAS field number. See maps in Fig. X.D.l. Can also be Galactic Plane field. 

coordinates of field center. DEC in degrees, RA in hours and minutes. Can also be LON 
and LAT for Galactic coordinate ( l n ,b n ) projections. 

wavelength band of the map. 


X-35 




Second line: 

RA & DEC GNOMONIC: Projection type. 

Can also be EQUIVALENT CYLINDRICAL or LON & LAT AITOFF 

PIXEL: Pixel size. 

JD: Julian dates of the earliest and latest data in the map, accurate to 1/2 day. 

Third line: 

TOP,BOTTOM, LEFT, RIGHT: Line and sample numbers at the edges of the map. 

1 TIC: Spacing between the tic marks around the edge in pixels 

Fourth line: 

ODN, 255DN: Surface brightness values of the minimum and maximum pixels in the map. 

5TH ROOT: Indicates that the fifth root of the actual surface brightness is shown. Can also be 

LINEAR in which case the true, uncompressed, surface brightness is shown. 

DATE: Year/Month/Day on which the map was assembled. 

D.3.b Tapes of 16.5° Images 

The magnetic tape form of the 16.5° images contains the calibrated surface brightness data in 
499x499 arrays of 2 x2' pixels recorded in the FITS format. The article by Wells et al. (1981) in con- 
junction with the label records of each tape file gives a detailed description of the format of each map 
image. A brief description of the format follows and a listing of a sample FITS label can be found in 
Appendix X.3. 

One sky coverage consists of 27 tapes of 6250 bpi (bits per inch) density. The third coverage has 
only 24 tapes with a total of 188 plates. Each plate consists of four surface brightness maps and four sta- 
tistical weight maps, one of each for each wavelength band. The plates are ordered on the tapes by plate 
number and within a plate the image files are ordered: 12 pm brightness, 12 pm weight, 25 pm bright- 
ness, 25 pm weight, 60 pm brightness, 60 pm weight, 100 pm brightness and 100 pm weight. The first 
two records of each file contain the label, then the image appears as a stream of pixel values divided into 
2880 byte records without regard for line length. The stream begins with the smallest line and smallest 
sample number and the sample number increases fastest. The last record is padded to 2880 bytes with 
zeros. Four-byte integers are used for brightness image data numbers and 2-byte integers for weight 
images, high order byte first. All this and other information necessary to successfully regenerate a map is 
contained in the FITS label records described in Appendix X.3. 

D.4 Galactic Plane Maps 

For convenience in dealing with the Galactic plane the survey data within 10° of the Galactic plane 
were remapped from the into a set of images in Galactic coordinates to cover the full circle of the 
Galaxy. This remapping from the 16.5° images resulted in a slight degradation in resolution even though 
the pixel size was the same in both sets of maps. Twenty-four 16.7° x 20° fields cover the Galactic plane 
with field centers at integral multiples of 15° Galactic longitude. The three sky coverages of the survey 
were separated into three sets of maps. The image format is 499 lines of 599 samples each, projected 
from Galactic coordinates with an equal area cylindrical projection (see Section X.D.2.b). Galactic plane 
maps are available in both photographic and FITS tape formats. Two 6250 bpi tapes of 12 maps each 
hold the 24 Galactic plane maps. The differences in the FITS label between the 16.5° images and the 



Galactic plane maps are noted in the description and listing of the FITS label in Appendix X.3. No sta- 
tistical weight images are included for the Galactic plane maps. Statistical weight information may be 
obtained from the 16.5° maps. Coordinate overlays described in Section X.D.7 are available for the 
Galactic plane maps. 

D.5 Low-Resolution All-Sky Maps 

The 1/2° x 1/2° beam data contained in the Zodiacal History file described below was split into the 
three separate sky coverages and assembled into three all-sky maps with an Aitoff equal area projection in 
Galactic coordinates. Two fields of each sky coverage were produced; one centered on the Galactic 
center and one centered on the Galactic anti-center. The pixel arrays consist of 325 lines of 649 samples 
each. Galactic north is in the direction of decreasing line number (up) and Galactic east in the direction 
of decreasing sample number (left). The all-sky maps are available in both photographic and FITS tape 
forms. The formats of the photographic and tape forms of all-sky maps are very similar to those of the 
16.5° images (see Section X.D.3) with the differences described in the labels of the photographs and tape 
files. Coordinate overlays described in Section X.D.7 are available for the all-sky maps. 

D.6 Zodiacal Observation History File 

For convenience in the analysis and treatment of background emission from interplanetary dust 
(zodiacal emission) and other extremely large scale emission features, the survey data were averaged to 
1/2° x 1/2° beam size and along with pointing information was preserved as a time ordered data set con- 
taining all three sky coverages of the survey. This file is available on magnetic tape written with the for- 
mat described in Appendix X.4. 

D.7 Coordinate Overlays 

A set of coordinate overlays for the photographs is available as photographic negative transparencies 
in the 5 inch sq. format. The scale is identical to the corresponding map product so the overlays will be 
the correct size if enlarged by the same factor as the map. One overlay is provided for each declination 
zone from -30° to +30° where the overlays for zones of opposite sign are obtained by rotating the grids 
through 180°. Five overlays are provided for each declination zone between 45° and 75° to accommodate 
the fact that integer hour meridians cross the plates in these zones in five different configurations. Again 
the overlay for the zone of opposite sign is obtained by rotating the grid 180°. All integer hour meridians 
are labeled 00M. The hour of right ascension should be determined from the position of the plate center 
given in the label on the photograph. The plate numbers to which a particular overlay pertains are 
printed in the lower right comer of the overlay. The overlays are aligned by matching the two triangular 
fiducials along each edge of the overlay with the two large map border tics which straddle the center of 
each side of the map. The two overlays for the polar regions are similarly aligned with the fiducials and 
tics. The correct orientation makes the lettering on the overlay read the same way as the lettering in the 
plate label. 

One overlay is used for all Galactic plane maps. It is aligned with the same method as the sky 
plates. One orientation of the overlay is used for even numbered fields and has 0° as the center longi- 
tude; the other orientation, used for all fields, has 5° as the center longitude. The tens digit of the true 
longitude should be obtained from label of the picture. 


X-37 



The overlays for the low-resolution all-sky maps come in only two varieties, l n - 0.0° in the center 
and \ u - 180.0° in the center. Alignment of the overlays is similar to that for the other maps. 

E. Low-Resolution Spectra 

Two files give the data for the low-resolution spectra (LRS) which, as described in Chapter IX, con- 
sist of two bands (8-13 pm and 1 1-22 pm). The first file on the tape consists of a single 80-character 
ASCII record giving the date and version of the LRS data. The next file lists only the wavelengths 
corresponding to each sample of the spectra (Table X.E.la). The last file contains the spectra and associ- 
ated header information (Table X.E.lb). The spectra are in order of increasing right ascension. Both 
files have 80-character ASCII logical records and 256 logical records per physical record. An atlas of the 
spectra will be published as an Astronomy and Astrophysics Supplement. 

E. 1 Catalog Header File 

Distance from Slit. ANGLE (100) 

As described in detail in Chapter IX, each sample in the spectrum corresponds to a certain in-scan 
distance of the source from the centerline of the spectrometer entrance aperture. ANGLE lists these 
angular distances. 

Wavelength Calibration: LAMBDA 1 , LAMBDA2 

There is a non-linear relation between the displacement of the source from the centerline of the 
spectrometer and the sampled wavelength. The wavelengths corresponding to each sample (or ANGLE) 
are given in LAMBDA 1 and LAMBDA2 for the two wavelength bands. The beginning and end of 
each spectrum contain measurements that lie outside of the wavelength coverage of the instrument but 
which can be used for electronic baseline determination. The wavelength values corresponding to these 
values of sample number (or ANGLE) are set to 0. 

E.2 Spectra Records 

The items starting at bytes 0 through 25 are identical to the items with the same names in the point 
source catalog (see Section X.B.l). This is also the case for the items starting at bytes 71, 73, 75 and 960. 

Number of Spectra: NSPECTRA, NACCEPT 

NSPECTRA is the number of spectra of this source observed; NACCEPT are the numbers of 8-13 
pm and the 1 1-22 pm spectrum halves ultimately averaged to make the entry in the catalog. 

Characterization of the Spectrum: LRSCHAR 

A description of the method of characterization is given in Section IX.D. Table X.E.2 lists the 
spectral classes used to characterize the spectrum. 

Quality of the Spectrum Halves: SPQUAL 

Depending on the signal-to-noise ratio of the 8-13 pm and 11-22 pm halves of the spectra, the 
number of accepted spectra halves and the difference in level of the baselines on either side of the spec- 
trum halves, a quality digit is assigned to each half of the spectrum; 1 indicates good quality, 2 moderate 
quality, and 3 barely acceptable. 


X-38 


'1 II 


Scale Factor for All Spectrum Flux Densities: SCALE 

Multiplying the integers BASELINE, NOISE, and SPECTRUM by the factor SCALE converts the 
values into units of W m -2 pm -1 . 

Baseline of Spectrum Halves: BASELINE 

These four values give the average value of samples 1 through 20 (short wavelength end) and 8 1 
through 100 (long wavelength end) of both the 8-13 pm and the 1 1-22 pm halves of the spectrum. 

RMS Noise: NOISE 

Using the twenty samples on the long wavelength end used for BASELINE, the rms noise per spec- 
trum half is determined. 

Signal-to-Noise Ratio: SNR 

The average value of the samples in the wavelength ranges 8-13 pm and 1 1-22 pm, respectively, are 
divided by the NOISE values determined. 

Baseline Asymmetry: ASYMM 

This value indicates by what fraction of the average signal the baselines on the short and long 
wavelength sides of the spectrum halves differ. A large baseline asymmetry indicates that a confusing 
source may have contaminated the spectrum. The baseline asymmetry is usually large near the Galactic 
plane. 

LRS/Survev Flux Ratio: SRATIO 

The ratio of the integrated flux (after convolution of the spectrum flux densities with the 12 pm 
band pass of the survey instrument) in the spectrum and the 12 pm survey flux is given in this item. 
Normally this value should be close to unity (Section IX. C). As the 1 1-22 pm part of the LRS spec- 
trum hardly overlaps with the 25 pm survey band, a ratio of LRS/survey for this band is not significant. 

The Spectrum: SPECTRUM 

The integer values of the spectrum must be multiplied by SCALE for conversion to W 
m -2 pm -1 . The wavelengths corresponding to the 100 samples given for the 8-13 pm and the 11-22 pm 
halves of the spectrum are given in the catalog header file. Values for non-significant wavelengths are set 
to zero. For baseline interpolation, either the sample numbers or the ANGLE (distance in arcmin from 
the spectrometer center line) can be used. 


X-39 



Table X.E.la 

Header Information for Catalog of Spectra 


Start 





Byte 

Name 

Description 

Units 

Format 

0 

ANGLE 

Angle from center of 
slit (one number per 
sample) 

arc min 

100F8.4 

800 

LAMBDA 1 

Wavelength in 8- 1 3 pm 
corresponding to each 
sample 

pm 

100F8.4 

1600 

LAMBDA2 

Wavelength in 1 1-22 pm 
corresponding to each 
sample 

pm 

100F8.4 





Table X.E.lb. Format of Spectra in Catalog 


Start 

Byte Name Description Units Format 


0 

NAME 

Source name 

— 

1 1A1 

11 

HOURS 

RA 1950 

hrs 

12 

13 

MINUTE 

RA 1950 

min 

12 

15 

SECOND 

RA 1950 

deci-sec 

13 

18 

DSIGN 

Declination sign 

± 

A1 

19 

DECDEG 

DEC 1950 

arc deg 

12 

21 

DECMIN 

DEC 1950 

arc min 

12 

23 

DECSEC 

DEC 1950 

arc sec 

12 

25 

FLUX 

Averaged non-color 
corrected flux 
densities ( 1 value per 
band) 

Jy 

4E9.3 

61 

NSPECTRA 

No. of observed spectra 

— 

12 

63 

NACCEPT 

No. of accepted spectrum 
halves 

_ 

212 

67 

LRSCHAR 

Characterization of 
spectrum 

_ 

211 

69 

SPQUAL 

Quality of 8-13 pm and 
1 1-22 pm parts of 
spectra 


211 

71 

VAR 

Percent of variability 
likelihood (from catalog) 

" 

12 

73 

NID 

(<25) 

No. of associations 

” 

12 

75 

IDTYPE 

Type of association 

— 

11 

76 

SPARE 

4 spare bytes 

( new record 

--) 

4A1 

80 

BASELINE 

Average of outer 20 
samples of spectra. 

Short and long wavelength 
end for each spectrum half 

scaled by SCALE 

414 

96 

NOISE 

RMS noise per sample 
(one value per spectrum half) 

scaled by SCALE 

214 

104 

SNR 

Signal-to-noise ratio 
(average signal in 
spectrum part divided 
by noise, one value per 
spectrum half) 


2E10.3 


X-41 



Start 

Byte 

Table X.E.lb Format of Spectra in Catalog (Continued) 

Name Description Units 

Format 

124 

ASYMM 

Relative baseline 
asymmetry (difference 
of left and right 
baselines divided by 
average signal, one value 
per band) 

— 

2E10.3 

144 

SRATIO 

Ratio of integrated LRS 
flux to 1 2 pm survey 
flux 


F5.2 

149 

SCALE 

Scale factor for all 
flux densities 

( new record 

W m“ 2 pm -1 

-) 

El 1.5 

160 

SPECTRUM 

100 samples for each 
of the two bands 
(8-13 pm, 1 1-22 pm), 
scaled by SCALE 

( new record 

scaled by SCALE 
-) 

20014 

960 

1000 

ID#1 

ID#2 

Association field from 
main catalog (see 
Section X.B.l) 
additional records 
as required. 


40AI 





Authors: 

C. Beichman, T. Chester, T.N. Gautier, G. Helou, C. Oken, E. Raimond, B.T. Soifer and D. 
Walker. 


References: 

Richardus, Peter and Ron K.. Adler, 1972, Map Projections for Geodesists, Cartographers and Geo- 
graphers, North-Holland Publishing Company, Amsterdam-London and American Elsevier Publish- 
ing Company, New York. 

Wells, D., Greisen, E.R. and Harten, R. 1981, Astron. and Astrophys. [Suppl . ], 44, 363. 


X-43 


i 


i 


< 

i 


i 


i 

j 

I 

i 

i 

j 

i 

i 


HI Hi 



Appendix X.l Regions of High Source Density 


As discussed in section V.H.6, regions with more sources per sq. deg than the confusion limit in a 
given band were processed according to more stringent rules to insure the reliability of the data presented 
in the catalog. The regions so processed were selected on a band by band basis depending on the number 
of sources with high or moderate quality fluxes located within a 1 sq. deg bin in ecliptic coordinates. 

Users of the catalog who want to know the identity of those regions may use a machine readable 
file that lists the bin number (see below), a flag called HSD which indicates which bands in that bin 
were processed according to high source density rules, the ecliptic coordinates of the center of the bin, 
and its length in ecliptic longitude. All bins have a height of 1° in ecliptic latitude. The file consists of 
80-character ASCII records with 256 records per physical block. The information for two bins fits within 
a single 80-character logical record. This file is preceded on the tape by a short file which contains a sin- 
gle 80-character ASCII record which lists the date and version number of the information. 

The flag HSD is hex-encoded by band, as described in Table X.B.2. In this notation each band 
corresponds to one of the four bits of a hex digit. 12 pm corresponds to bit 0 (Least Significant Bit) and 
100 pm to bit 3, etc. If a band went through high source density processing in that band, then the 
appropriate bit in the hex digit was set. Thus, if high source density rules were invoked at 25, 60 and 100 
um, HSD would have the value 1110 (binary) - E (hex). 

Ecliptic bins start at the ecliptic north pole and step around the sky in bands of constant ecliptic 
latitude, stepping 1° southward after completing each band. The length of the bin, in ecliptic longitude, 
was adjusted for the cosine of the ecliptic latitude to maintain an approximately constant area. FOR- 
TRAN programs are given below to convert from bin number to ecliptic coordinates and vice versa. 
There is a known bug in the computation of bins at ecliptic latitude 60°, causing bin 2842 to be skipped. 
The following program, while incorrect in this way, will give results consistent with the bin numbers used 
in the data processing. To avoid problems with roundoff errors within a few arcseconds of bin boun- 
daries, all arithmetic should be calculated in double precision. 


PRECEDING PAGE BLANK NOT FILMED 


pas t mm 


X-45 



n n n on n nn 


FORTRAN Program to Convert between Ecliptic Position and Bin Number 


CONVERT AN ECLIPTIC POSITION TO BIN NUMBER. 
SUBROUTINE P2BIN(LAM,BET,BIN) 

IMPLICIT REAL*8 A-H, P-Z 

REAL*8 LAM,BET,LSIZE,R2D/57. 2957795/ 

INTEGERS BIN,I, J,N,MAXBIN( 1 82) 

COMMON /LATBND/MAXBIN 

I - 90. - BET*R2D + 2.5 
N - MAXBIN(I) - MAXBIN(I-l) 

LSIZE - 360./FLOAT(N) 

J - LAM*R2D/LSIZE + 1 
BIN - MAXBIN(I-l) + J 
RETURN 
END 


CONVERT A BIN NUMBER TO A CENTER POSITION AND LONGITUDE LENGTH. 
SUBROUTINE BIN2P(BIN, CLON, CLAT, LWIDTH) 

IMPLICIT REAL*8 A-H, P-Z 
INTEGERS BIN,I,J,MAXBIN( 1 82), FIRST 
REAL*8 CLON,CLAT,LWIDTH,R2D/57. 2957795/, NBINS 
COMMON /LATBND/MAXBIN 

DO 1001-1,182 

IF(MAXBIN(I) .GE. BIN) GOTO 150 
00 CONTINUE 
50 I - I - 1 

FIRST - MAXBIN(I) + 1 
NBINS - MAXBIN(I-H) - FIRST + 1 
L WIDTH - 360./NBINS 
J - BIN - FIRST 

CLON - DFLOAT(J)*LWIDTH + LWIDTH/2. 

CLAT - 90. -DFLOAT(I-l) 

LWIDTH - LWIDTH/R2D 
CLON - CLON/R2D 
CLAT - CLAT/R2D 
RETURN 
END 

DATA SUBPROGRAM. 

BLOCK DATA 
INTEGERS MAXBIN(182) 

DATA MAXBIN 

* /0, 1 ,7, 1 9,37,62,93, 1 30, 1 73,223,279,34 1 ,409,483,563, 

* 650,743,842,947, 1 058, 1 1 75, 1 298, 1 427, 1 56 1 , 1 70 1 , 1 847, 1 999, 

* 2 1 56,23 1 9,2488,2662,284 1 ,3026,32 1 6,34 1 2,36 1 3,38 19,4030, 


X-46 


* 4246,4467,4693,4924,5160,5400,5645,5895,6149,6407,6670, 

* 6937,7208,7483,7762,8045,8332,8623,8917,9215,9516,9821, 

* 101 29, 10440, 10754, 1 107 1 , 1 1 39 1 , 1 1 7 1 4, 1 2040, 1 2368, 1 2699, 

* 1 3032, 1 3368, 1 3706, 1 4046, 1 4388, 14732, 1 5078, 1 5425, 1 5774, 

* 1 6 1 24, 1 6476, 1 6829, 17183,17538,1 7894, 18251,1 8609, 1 8967, 

* 19326,19685,20044,20403,20763,21 122,21481,21840,22199, 

* 22557,229 1 5,23272,23628,23983,24337,24690,25042,25392, 

* 2574 1 ,26088,26434,26778,27 1 20,27460,27798,28 1 34,28467, 

* 28798,29126,29452,29775,30095,30412,30726,31037,31345, 

* 31650,31951,32249,32543,32834,33121,33404,33683,33958, 

* 34229,34496,34759,35017,35271,35521,35766,36006,36242, 

* 36473,36699,36920,37 1 36,37347,37553,37754,37950,38 140, 

* 38325,38505,38679,38848,3901 1,39168,39320,39466,39606, 

* 39740,39869,39992,40109,40220,40325,40424,40517,40604, 

* 40684,40758,40826,40888,40944,40994,41037,41074,41 105, 

* 41130,41148,41160,41166,41167/ 

COMMON /LATBND/MAXBIN 
END 


Table X.Apl.l Format of File of High Source Density Bins 

Start 


Byte 

Name 

Description 

Units 

Format 

00 

BINNUM 

SDAS Bin Number 

— 

16 

06 

HSD 

Bands processed 
for high source 
density, hex 
encoded by band. 


A2 

08 

LAMBDA 

Ecliptic longitude 
of bin center 

deg 

10.5 

18 

BETA 

Ecliptic latitude 
of bin center 

deg 

10.5 

28 

LENGTH 

Length of bin in 
ecliptic longitude 

deg 

10.5 

38 

40-79 

SPARE 

2 spare bytes 

bytes 00-39 are 
repeated for 
next bin 


2A1 


X-47 



N 



+ 1.5° 
+0.5° 
-0.5° 
-1.5° 


Figure X.Ap. 1 Scheme for obtaining 1 sq. deg bins in ecliptic coordinates. 


X-48 










Appendix X.2 Location of 16.5° Image Fields 


PLATE 

DEC 

Table X.Ap2.1 16.5° Field Centers (Equinox 1950) 

RA PLATE DEC RA PLATE 

DEC 

RA 

1 

+90 

OH OOM* 

43 +45 

19H 12M 

87 

+ 15 

16H OOM 




44 

20H 24M 

88 


17H OOM 

2 

+75 

OH OOM 

45 

21H 36M 

89 


18H OOM 

3 


2H 24M 

46 

22H 48M 

90 


19H OOM 

4 


4H 48M 

47 +30 

OH OOM 

91 


20H OOM 

5 


7H 12M 

48 

1H OOM 

92 


21H OOM 

6 


9H 36M 

49 

2H OOM 

93 


22H OOM 

7 


12H OOM 

50 

3H OOM 

94 


23H OOM 

8 


14H 24M 

51 

4H OOM 

95 

+0 

OH OOM 

9 


16H 48M 

52 

5H OOM 

96 


1H OOM 

10 


19H 12M 

53 

6H00M 

97 


2H OOM 

11 


21H 36M 

54 

7H OOM 

98 


3H OOM 




55 

8H OOM 

99 


4H OOM 

12 

+60 

OH OOM 

56 

9H OOM 

100 


5H OOM 

13 


1H 36M 

57 

10H OOM 

101 


6H OOM 

14 


3H 12M 

58 

1 1H OOM 

102 


7H OOM 

15 


4H 48M 

59 

12H00M 

103 


8H OOM 

16 


6H 24M 

60 

13H OOM 

104 


9H OOM 

17 


8H OOM 

61 

14H OOM 

105 


10H OOM 

18 


9H 36M 

62 

15H OOM 

106 


11H OOM 

19 


11H 12M 

63 

16H OOM 

107 


12H00M 

20 


12H 48M 

64 

17H OOM 

108 


13H OOM 

21 


14H 24M 

65 

18H00M 

109 


14H OOM 

22 


16H OOM 

66 

19H OOM 

110 


15H OOM 

23 


17H 36M 

67 

20H OOM 

111 


16H OOM 

24 


19H 12M 

68 

21H00M 

112 


17H OOM 

25 


20H 48M 

69 

22H OOM 

113 


18H OOM 

26 


22H 24M 

70 

23H OOM 

114 


19H OOM 






115 


20H OOM 

27 

+45 

OH OOM 

71 +15 

OH OON 

116 


21H00M 

28 


1H 12M 

72 

1H OOM 

117 


22H OOM 

29 


2H 24M 

73 

2H OOM 

118 


23H OOM 

30 


3H 36M 

74 

3H OOM 

119 

-15 

OH OOM 

31 


4H 48M 

75 

4H OOM 

120 


1H OOM 

32 


6H OOM 

76 

5H OOM 

121 


2H OOM 

33 


7H 12M 

77 

6H OOM 

122 


3H OOM 

34 


8H 24M 

78 

7H OOM 

123 


4H OOM 

35 


9H 36M 

79 

8H OOM 

124 


5H OOM 

36 


10H 48M 

80 

9H OOM 

125 


6H OOM 

37 


12H OOM 

81 

10H OOM 

126 


7H OOM 

38 


13H 12M 

82 

1 1H OOM 

127 


8H OOM 

39 


14H 24M 

83 

12H OOM 

128 


9H OOM 

40 


15H 36M 

84 

13H OOM 

129 


10H OOM 

41 


16H 48M 

85 

14H OOM 

130 


1 1H OOM 

42 


18H OOM 

86 

15H OOM 

131 


12H OOM 



X-49 



PLATE 

DEC 

Table X.Ap2.1 
RA PLATE 

16.5° Field Centers (Continued) 

DEC RA PLATE 

DEC 

RA 

132 

-15 

13H OOM 

171 

-45 

4H 48M 2 1 1 

-75 

21H 36M 

133 


14H OOM 

172 


6H00M 212 

-90 

OH OOM* 

134 


15H OOM 

173 


7H 12M 



135 


16H OOM 

174 


8H 24M 



136 


17H OOM 

175 


9H 36M 



137 


18H OOM 

176 


10H 48M 



138 


19H00M 

177 


12H OOM 



139 


20H OOM 

178 


13H 12M 



140 


21H00M 

179 


14H 24M 



141 


22H OOM 

180 


15H 36M 



142 


23H OOM 

181 


16H 48M 






182 


18H00M 



143 

-30 

OH OOM 

183 


19H 12M 



144 


1H OOM 

184 


20H 24M 



145 


2H OOM 

185 


21H 36M 



146 


3H OOM 

186 


22H 48M 



147 


4H OOM 






148 


5H OOM 

187 

-60 

OH OOM 



148 


5H OOM 

188 


1H 36M 



149 


6H OOM 

189 


3H 12M 



150 


7H00M 

190 


4H48M 



151 


8H OOM 

191 


6H 24M 



152 


9H OOM 

192 


8H OOM 



153 


1 OH OOM 

193 


9H 36M 



154 


1 1H OOM 

194 


11H 12M 



155 


12H OOM 

195 


12H 48M 



156 


13H OOM 

196 


14H 24M 



157 


14H OOM 

197 


16H OOM 



158 


15H OOM 

198 


17H 36M 



159 


16H OOM 

199 


19H 12M 



160 


17H OOM 

200 


20H 48M 



161 


18H OOM 

201 


22H 24M 



162 


19H OOM 






163 


20H OOM 

202 

-75 

OH OOM 



164 


21H00M 

203 


2H 24M 



165 


22H OOM 

204 


4H 48M 



166 


23H OOM 

205 


7H 12M 






206 


9H 36M 



167 

-45 

OH OOM 

207 


12H OOM 



168 


1H 12M 

208 


14H 24M 



169 


2H 24M 

209 


16H 48M 



170 


3H 36M 

210 


19H 12M 




These RA’s produce the correct orientations 
when used in the transformation formulae. 





Appendix X.3 Sample FITS Headers 
Header for Sky Plate Intensity 


SIMPLE - 
BITPIX - 
NAXIS - 
NAXIS1 - 
NAXIS2 - 
NAXIS3 - 
BSCALE - 
BZERO - 
BUNIT - 
BLANK - 
CRVAL1 - 
CRPIX1 - 
CTYPEl - 
COMMENT 
CDELT1 - 
COMMENT 
CRVAL2 - 
CRPIX2 - 
CTYPE2 - 
COMMENT 
CDELT2 - 
COMMENT 
CRVAL3 - 
CRPIX3 - 
CTYPE3 - 
CDELT3 - 
DATAMAX - 
DAT AMIN - 
EPOCH - 
DATE-MAP - 
DATE - 
ORIGIN - 
TELESCOP - 
INSTRUME - 
OBJECT - 
PROJTYPE - 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 



T 

/ STANDARD FITS FORMAT 


32 

/ 4 BYTE TWOS COMPL INTEGERS 


3 

/ NUMBER OF AXES 


nn 

/ # SAMPLES PER LINE (FASTEST VARY INDEX) 


mm 

/ # LINES OF DATA IN IMAGE FILE 


1 

/ # WAVELENGTHS 


S.ssssssE-ee 

/ TRUE-TAPE*SCALE+BZERO 


0.0 

/ 

’JY/SR’ 


/ INTENSITY 


-2000000000 

/ TAPE VALUE FOR EMPTY PIXEL 


DD.dd 

/ RA AT ORIGIN (DEGREES) 


XXX. 

/ SAMPLE AXIS ORIGIN (PIXEL) 

’RA — TAN’ 


/ DECREASES IN VALUE AS SAMPLE INDEX 
INCREASES (GNOMONIC PROJECTION) 


-3.333333E-02 

/ COORD VALUE INCREMENT DEG/PIXEL 
AT ORIGIN ON SAMPLE AXIS 


DD.dd 

/ DEC AT ORIGIN (DEGREES) 


yyy- 

/ LINE AXIS ORIGIN (PIXEL) 

’DEC-TAN’ 


/ DECREASES IN VALUE AS LINE INDEX 



INCREASES (GNOMONIC PROJECTION) 


-3.333333E-02 

/ COORD VALUE INCREMENT DEG/PIXEL 
AT ORIGIN ON LINE AXIS 


zzzE-6 

1. 

0. 

D.ddddddE ee 

/ WAVELENGTH IN METERS 

’LAMBDA ’ 



/JY/SR (TRUE VALUE) 


D.ddddddE ee 

/JY/SR (TRUE VALUE) 


1950. 

/ EME50 

’01/11/84’ 


/ MAP RELEASE DATE 

’DD/MM/YY’ 


/ DATE THIS TAPE WRITTEN 

’JPL-IRAS’ 

’IRAS 


/ INSTITUTION 

’SKYPLATE’ 


/ IRAS SKY PLATE 

’PLnnn Hn’ 


/ PLATE NUMBER / HCON 

’GNOMONIC’ 


/ PROJECTION TYPE 

MINSOP - MMM; MAXSOP - NNN 
LOGTAG - VSFLOGK 7.6) 

GEOMTAG - GEOM( 7.5) 


PROJECTION FORMULAE: 

FORWARD FORMULA; RAO AND DECO ARE THE PLATE CENTER 
R2D - 45. / ATAN(1.) 

PIX - 30. 

A - COS(DEC) * COS(RAO - RA) 

F - PIX * R2D / (SIN(DECO) * SIN(DEC) + A * COS(DECO)) 

SAMPLE - -F * COS(DEC) * SIN(RA-RAO) 

XLINE - -F * (COS(DECO) * SIN(DEC) - A * SIN(DECO)) 

INVERSE FORMULA; REQUIRES ARCSINE 


X-51 



COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

END 


X - SAMPLE / (PIX * R2D) 

Y - LINE / (PIX * R2D) 

DELTA - ATAN(SQRT(X*X + Y*Y)) 

BETA - ATAN2(-X,Y) 

DEC - ASIN(SIN(DECO)*COS(DELTA)-COS(DECO)*SIN(DELTA)*COS(BETA)) 
XX - SIN(DECO)*SIN(DELTA)*COS(BETA)+COS(DECO)*COS(DELTA) 

YY - SIN(DELTA)*SIN(BETA) 

RA - RAO + ATAN2(YY,XX) 

REFERENCES: 

IRAS SDAS SOFTWARE INTERFACE SPECIFICATION(SIS) #623-94/NO. SF05 
ASTRON. ASTROPHYS. SUPPL. SER. 44,(1981) 363-370 (RE:FITS) 
RECONCILIATION OF FITS PARMS W/ SIS SFOS PARMS: 

NAXIS1 - (ES - SS + 1); NAXIS2 - (EL - SL + 1); 

CRPIX1 - (1 - SS); CRPIX2 - (1 - SL) 



Header for 16.5 deg Image Weight 


SIMPLE - 
BITPIX - 
NAXIS - 
NAXIS1 - 
NAXIS2 - 
NAXIS3 - 
BSCALE - 
BZERO - 
BUNIT - 
BLANK - 
CRVAL1 - 
CRPIX1 - 
CT Y PEI - 
COMMENT 
CDELT1 0- 
COMMENT 
CRVAL2 - 
CRPIX2 - 
CTYPE2 - 
COMMENT 
CDELT2 - 
COMMENT 
CRVAL3 - 
CRPIX3 - 
CTYPE3 - 
CDELT3 - 
DATAMAX - 
DATA MIN - 
EPOCH - 
DATE-MAP - 
DATE - 
ORIGIN - 
TELESCOP - 
INSTRUME - 
OBJECT - 
PROJTYPE - 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 


’RA~TAN’ 


’DEC-TAN’ 


’LAMBDA ’ 


’01/11/84’ 

’DD/MM/YY’ 

’JPL-IRAS’ 

’IRAS 

’SKYPLATE’ 
’PLnnn Hn’ 
’GNOMONIC’ 


T 

16 

3 

nn 

mm 

1 

1. 

0.0 

0 

DD.dd 

xxx. 


-3.333333E-02 

DD.dd 

yyy- 

-3.333333E-02 


z.zzE-6 

1 . 


/ STANDARD FITS FORMAT 
/ 2 BYTE TWOS COMPL INTEGERS 
/ NUMBER OF AXES 

/ # SAMPLES PER LINE (FASTEST VARY NDEX 
/ # LINES OF DATA IN IMAGE RLE 
/ # WAVELENGTHS 
/ TRUE— TAPE*SCALE+BZERO 

/ 

/ STATISTICAL WEIGHT 
/ TAPE VALUE FOR EMPTY PIXEL 
/ RA AT ORIGIN (DEGREES) 

/ SAMPLE AXIS ORIGIN (PIXEL) 

/ DECREASES IN VALUE AS SAMPLE INDEX 
INCREASES (GNOMONIC PROJECTION) 

/ COORD VALUE INCREMENT DEG/PIXEL 
AT ORIGIN ON SAMPLE AXIS 
/ DEC AT ORIGIN (DEGREES) 

/ LINE AXIS ORIGIN (PIXEL) 

/ DECREASES IN VALUE AS LINE INDEX 
INCREASES (GNOMONIC PROJECTION) 

/ COORD VALUE INCREMENT DEG/PIXEL 
AT ORIGIN ON LINE AXIS 
/ WAVELENGTH IN METERS 


0 . 

D.ddddddE ee / DIMENSIONLESS 

D.ddddddE ee / DIMENSIONLESS 

1950. / EME50 

/ MAP RELEASE DATE 
/ DATE THIS TAPE WRITTEN 
/ INSTITUTION 

/ IRAS SKY PLATE 
/ PLATE NUMBER / HCON 
/ PROJECTION TYPE 

MINSOP - MMM; MAXSOP - NNN 
LOGTAG - VSFLOG (7.8) 

GEOMTAG - GEOM (7.5) 


PROJECTION FORMULAE: 

FORWARD FORMULA; RAO AND DECO ARE THE PLATE CENTER 
R2D - 45. / ATAN(1.) 

PIX - 30. 

A - COS(DEC) * COS(RA0 - RA) 

F - PIX * R2D / (SIN(DECO) * SIN(DEC) + A * COS(DECO)) 

SAMPLE - -F * COS(DEC) * SIN(RA-RA0) 

XLINE - -F * (COS(DECO) * SIN(DEC) - A * SIN(DECO)) 

INVERSE FORMULA; REQUIRES ARCSINE 
X - SAMPLE / (PIX * R2D) 

Y - LINE / (PIX * R2D) 

DELTA - AT AN(SQRT(X*X + Y*Y)) 


X-53 



COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

END 


BETA - ATAN2(-X,Y) 

DEC - ASIN(SIN(DECO)*COS(DELTA)-COS(DECO)*SIN(DELTA)*COS(BETA)) 
XX - SIN(DECO)*SIN(DELTA)*COS(BETA)+COS(DECO)*COS(DELTA) 

YY - SIN(DELT A)*SIN(BET A) 

RA - RAO 4- ATAN2(YY,XX) 

REFERENCES: 

IRAS SDAS SOFTWARE INTERFACE SPECIHCATION(SIS) #623-94/NO. SF05 
ASTRON. ASTROPHYS. SUPPL. SER. 44,(1981) 363-370 (RE:FITS) 
RECONCILIATION OF RTS PARMS W/ SIS SFOS PARMS: 

NAXIS1 - (ES - SS + 1); NAXIS2 - (EL - SL + 1); 

CRPIX 1 - (1 - SS); CRPIX2 - (1 - SL) 


Header for Galactic Plane Fields 


SIMPLE - 
BITPIX - 
NAXIS - 
NAXIS1 - 
NAXIS2 - 
NAXIS3 - 
BSCALE - 
BZERO - 
BUNIT - 
BLANK - 
CRVAL1 - 
CRPIX1 - 
CTYPE1 - 
COMMENT 
COMMENT 
CDELT1 - 
COMMENT 
CRVAL2 - 
CRPIX2 - 
CTYPE2 - 
COMMENT 
COMMENT 
CDELT2 - 
COMMENT 
CRVAL3 - 
CRPIX3 - 
CTYPE3 - 
CDELT3 - 
DATAMAX - 
DATAMIN - 
EPOCH - 
DATE-MAP - 
DATE - 
ORIGIN - 
TELESCOP - 
INSTRUME - 
OBJECT - 
PROJTYPE - 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 


’JY/SR’ 


’LON 


’LAT-SIN 1 


’LAMBDA ’ 


’01/11/84’ 

’DD/MM/YY’ 

’JPL-IRAS’ 

’IRAS 

’GALPLANE’ 
’GPLnn Hn’ 
’LAMBECYL’ 


T 

32 

3 

nn 

mm 

1 

S.SSSSSSE-ee 

0.0 

-2000000000 

DD.dd 

xxx. 


-3.333333E-02 

0.0 

yyy- 


-3.333333E-02 


zzzE-6 

1 . 


/ STANDARD FITS FORMAT 
/ 4 BYTE TWOS COMPL INTEGERS 
/ NUMBER OF AXES 

/ # SAMPLES PER LINE (FASTEST VARY NDEX) 

/ # LINES OF DATA IN IMAGE FILE 
/ # WAVELENGTHS 
/ TRUE-T APE*SCALE+BZERO 

/ 

/ INTENSITY 

/ TAPE VALUE FOR EMPTY PIXEL 
/ GALACTIC LONGITUDE AT ORIGIN (DEGREES) 
/ SAMPLE AXIS ORIGIN (PIXEL) 

/ DECREASES IN VALUE AS SAMPLE INDEX 
INCREASES (LAMBERT EQUIVALENT 
CYLINDRICAL PROJECTION) 

/ COORD VALUE INCREMENT DEG/PIXEL 
AT ORIGIN ON SAMPLE AXIS 
/ GALACTIC LATITUDE AT ORIGIN (DEGREES) 

/ LINE AXIS ORIGIN (PIXEL) 

/ DECREASES IN VALUE AS LINE INDEX 
INCREASES (LAMBERT EQUIVALENT 
CYLINDRICAL PROJECTION) 

/ COORD VALUE INCREMENT DEG/PIXEL 
AT ORIGIN ON LINE AXIS 
/ WAVELENGTH IN METERS 


0 . 

D.ddddddE ee / JY/SR (TRUE VALUE) 

D.ddddddE ee / JY/SR (TRUE VALUE) 

1950. / EME50 

/ MAP RELEASE DATE 
/ DATE THIS TAPE WRITTEN 
/ INSTITUTION 

/ IRAS GALACTIC PLANE 
/ GALACTIC PLANE PLATE NUMBER / HCON 
/ PROJECTION TYPE 


PROJECTION FORMULAE: 

FORWARD FORMULA; XLONO IS THE LONGITUDE OF PLATE CENTER 
R2D-45. / ATAN(1.) 

PIX - 30. 

XLINE - -SIN(XLAT) * PIX * R2D 
SAMPLE - -(XLON - XLONO) * PIX 

INVERSE FORMULA; ARCSINE IS REQUIRED 
XLON - XLONO - SAMPLE / PIX 
XLAT - -ASIN(XLINE / (PIX * R2D)) 

REFERENCES: 

IRAS SDAS SOFTWARE INTERFACE SPECIFICATION(SIS) #623-94/NO. SF05 
ASTRON. ASTROPHYS. SUPPL. SER. 44,(1981) 363-370 (RE:FITS) 


X-55 



COMMENT 

COMMENT 

COMMENT 


RECONCILIATION OF FITS PARMS W/ SIS SF05 PARMS: 
NAXISI - (ES - SS + 1); NAXIS2 - (EL - SL + 1); 

CRPIX1 - (I - SS) CRPIX2 - (1 - SL) 


X-56 



Header for Low-Resolution All-Sky Image (Galactic Center) 


SIMPLE - 

T 

/ STANDARD FITS FORMAT 

BITPIX - 

32 

/ 4 BYTE TWOS COMPL INTEGERS 

NAXIS - 

3 

/ NUMBER OF AXES 

NAXIS1 - 

nn 

/ # SAMPLES PER LINE (FASTEST VARY NDEX) 

NAXIS2 - 

mm 

/ # LINES OF DATA IN IMAGE FILE 

NAXIS3 - 

1 

/ # WAVELENGTHS 

BSCALE - 

S.ssssssE-ee 

/ TRUE«TAPE*SCALE+BZERO 

BZERO - 

0.0 

/ 

BUNIT - 

’JY/SR’ 

/ INTENSITY 

BLANK - 

-2000000000 

/ TAPE VALUE FOR EMPTY PIXEL 

CRVAL1 - 

DD.dd 

/ GALACTIC LONGITUDE AT ORIGIN (DEGREES) 

CRPIX1 - 

XXX. 

/ SAMPLE AXIS ORIGIN (PIXEL) 

CTYPE1 - 

’LON-ATF ’ 

/ DECREASES IN VALUE AS SAMPLE INDEX 

COMMENT 


INCREASES (AITOFF PROJECTION) 

CDELT1 - 

-0.5 

/ COORD VALUE INCREMENT DEG/PIXEL 

COMMENT 


AT ORIGIN ON SAMPLE AXIS 

CRVAL2 - 

0.0 

/ GALACTIC LATITUDE AT ORIGIN (DEGREES) 

CRPIX2 - 

yyy- 

/ LINE AXIS ORIGIN (PIXEL) 

CTYPE2 - 

’LAT-ATF ’ 

/ DECREASES IN VALUE AS LINE INDEX 

COMMENT 


INCREASES (AITOFF PROJECTION) 

CDELT2 - 

-0.5 

/ COORD VALUE INCREMENT DEG/PIXEL 

COMMENT 


AT ORIGIN ON LINE AXIS 

CRVAL3 - 

z.zzE-6 

/ WAVELENGTH IN METERS 

CRPIX3 - 

1. 


CTYPE3 - 

’LAMBDA ’ 


CDELT3 - 

0. 


DATAMAX - 

D.ddddddE ee 

/ JY/SR (TRUE VALUE) 

DAT AMIN - 

D.ddddddE ee 

/ JY/SR (TRUE VALUE) 

EPOCH - 

1950. 

/ EME50 

DATE-MAP - 

’01/11/84’ 

/ MAP RELEASE DATE 

DATE - 

’DD/MM/YY’ 

/ DATE THIS TAPE WRITTEN 

ORIGIN - 

’JPL-IRAS’ 

/ INSTITUTION 

INSTRUME - 

’ALL SKY ’ 

/ IRAS LOW RES ALL-SKY 

OBJECT - 

’CENTER n’ 

/ ALL-SKY GALACTIC CENTER / HCON 

PROJTYPE - 

’AITOFF ’ 

/ PROJECTION TYPE 

COMMENT 


MINSOP - MMM; MAXSOP - NNN 

COMMENT 



COMMENT 

PROJECTION FORMULAE; 


COMMENT 

FORWARD FORMULA; XLONO IS THE CENTER LONGITUDE OF THE 

COMMENT 

MAP. ARC-SINE AND ARC-COSINE FUNCTIONS ARE REQUIRED. 

COMMENT 

R2D - 45. / ATAN(1.) 


COMMENT 

PIX - 2. 


COMMENT 

RHO - ACOS( COS(XLAT) * COS((XLON-XLONO)/2.) ) 

COMMENT 

THETA - ASIN( COS(XLAT) * SIN((XLON-XLONO)/2.) / SIN(RHO) ) 

COMMENT 

F - 2. * PIX * R2D * SIN(RHO/2.) 


COMMENT 

SAMPLE - -2. * F * SIN(THETA) 


COMMENT 

XLINE - -F * COS(THETA) 


COMMENT 

IF(XLAT .LT. 0.) XLINE - -XLINE 


COMMENT 



COMMENT 

REVERSE FORMULA; XLONO IS THE CENTER LONGITUDE OF THE MAP. 

COMMENT 

ARC-SINE AND ARC-COSINE FUNCTIONS NEEDED. 

COMMENT 

R2D - 45. / ATAN( 1 .) 



X-57 



COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

END 


PIX - 2. 

Y - -XLINE / (PIX * 2. * R2D) 

X - -SAMPLE / (PIX * 2. * R2D) 

A - SQRT(4.-X*X-4.*Y*Y) 

XLAT - R2D * ASIN(A*Y) 

XLON - XLONO + 2. * R2D * ASIN(A*X/(2 *COS(XLAT))) 

REFERENCES: 

IRAS SDAS SOFTWARE INTERFACE SPECIFTCATION(SIS) #623-94/NO. SF05 
ASTRON, ASTROPHYS. SUPPL. SER. 44,(1981) 363-370 (RE:FTTS) 
RECONCILIATION OF FITS PARMS W/ SIS SF05 PARMS: 

NAXIS1 - (ES - SS + 1); NAXIS2 - (EL - SL + 1); 

CRPIX1 - (1 - SS); CRPIX2 - (1 - SL) 



SIMPLE - 
BITPIX - 
NAXIS - 
NAXISi - 
NAXIS2 - 
NAXIS3 - 
BSCALE - 
BZERO - 
BUNIT - 
BLANK - 
CRVAL1 - 
CRPIX1 - 
CTYPE1 - 
COMMENT 
CDELT1 - 
COMMENT 
CRVAL2 - 
CRPIX2 - 
CTYPE2 - 
COMMENT 
CDELT2 - 
COMMENT 
CRVAL3 - 
CRPIX3 - 
CTYPE3 - 
CDELT3 - 
DATAMAX - 
DAT AMIN - 
EPOCH - 
DATE-MAP - 
DATE - 
ORIGIN - 
TELESCOP - 
INSTRUME - 
OBJECT - 
PROJTYPE - 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 
COMMENT 


Header for Low-Resolution All-Sky Intensity Galactic Anti-Center 


’JY/SR’ 


’LON-ATF ’ 


’LAT-ATF ’ 


’LAMBDA 


’01/11/84’ 

’DD/MM/YY’ 

’JPL-IRAS’ 

’IRAS 

’ALL SKY ’ 
’ACENTERn’ 
’AITOFF ’ 


T 

32 

3 

nn 

mm 

1 

S.ssssssE-ee 

0.0 

-2000000000 

180.0 

XXX. 


-0.5 


0.0 

yyy- 


-0.5 


z.zzE-6 

1 . 

0 . 

D.ddddddE ee 
D.ddddddE ee 
1950. 


/ STANDARD FITS FORMAT 
/ 4 BYTE TWOS COMPL INTEGERS 
/ NUMBER OF AXES 

/ # SAMPLES PER LINE (FASTEST VARY NDEX 
/ # LINES OF DATA IN IMAGE FILE 
/ # WAVELENGTHS 
/ TRUE-TAPE*SCALE+BZERO 

/ 

/ INTENSITY 

/ TAPE VALUE FOR EMPTY PIXEL 
/ GALACTIC LONGITUDE AT ORIGIN (DEGREES) 
/ SAMPLE AXIS ORIGIN (PIXEL) 

/ DECREASES IN VALUE AS SAMPLE INDEX 
INCREASES (AITOFF PROJECTION) 

/ COORD VALUE INCREMENT DEG/PIXEL 
AT ORIGIN ON SAMPLE AXIS 
/ GALACTIC LATITUDE AT ORIGIN (DEGREES) 

/ LINE AXIS ORIGIN (PIXEL) 

/ DECREASES IN VALUE AS LINE INDEX 
INCREASES (AITOFF PROJECTION) 

/ COORD VALUE INCREMENT DEG/PIXEL 
AT ORIGIN ON LINE AXIS 
/ WAVELENGTH IN METERS 


/ JY/SR (TRUE VALUE) 

/ JY/SR (TRUE VALUE) 

/ EME50 

/ MAP RELEASE DATE 
/ DATE THIS TAPE WRITTEN 
/ INSTITUTION 

/ IRAS LOW RES ALL-SKY 
/ ALL-SKY GALACTIC ANTI-CENTER / HCON 
/ PROJECTION TYPE 
MINSOP - MMM; MAXSOP - NNN 


PROJECTION FORMULAE: 

FORWARD FORMULA; XLONO IS THE CENTER LONGITUDE OF THE 
MAP. ARC-SINE AND ARC-COSINE FUNCTIONS ARE REQUIRED. 
R2D-45. / ATAN(1.) 

PIX - 2. 

RHO - ACOS( COS(XLAT) * COS((XLON-XLONO)/2.) ) 

THETA - ASIN( COS(XLAT) * SIN((XLON-XLONO)/2.) / SIN(RHO) ) 

F - 2. * PIX * R2D * SIN(RHO/2.) 

SAMPLE - -2. * F * SIN(THETA) 

XLINE - -F * COS(THETA) 

IF(XLAT .LT. 0.) XLINE - -XLINE 

REVERSE FORMULA; XLONO IS THE CENTER LONGITUDE OF THE MAP. 
ARC-SINE AND ARC-COSINE FUNCTIONS NEEDED. 


X-59 



COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

COMMENT 

END 


R2D-45. /ATAN(1.) 

PIX - 2. 

Y - -XLINE / (PIX * 2. * R2D) 

X - -SAMPLE / (PIX * 2. * R2D) 

A - SQRT(4.-X*X-4.*Y*Y) 

XLAT - R2D * ASIN(A*Y) 

XLON - XLONO + 2. * R2D * ASIN(A*X/(2 *COS(XLAT))) 

REFERENCES: 

IRAS SDAS SOFTWARE INTERFACE SPECIFICATION'S) #623-94/NO. SF05 
ASTRON. ASTROPHYS. SUPPL. SER. 44,(1981) 363-370 (REiFITS) 
RECONCILIATION OF FITS PARMS W/ SIS SF05 PARMS: 

NAXIS1 - (ES - SS + 1); NAXIS2 - (EL - SL + 1); 

CRPIX1 - (1 - SS); CRPIX2 - (1 - SL) 


X-60 


I 



Appendix X.Ap.4 Zodiacal Observation History File (ZOHF) Format 

The ZOHF is a low-resolution version of the entire IRAS sky survey. The average of all detectors 
within a band produced a 1/2° wide fan beam which was then boxcar averaged for eight seconds to pro- 
duce a 30" x 32" beam. Pointing information was similarly averaged over eight seconds to accompany 
the infrared data. The data base is time ordered and retains the subdivision of the original survey into 
semi-daily Satellite Operation Programs (SOP) and individual continuous survey scan observations (OBS). 

Data Base Organization: 

The ZOHF is organized into 572 files, with two consecutive file marks after the last file on each 
tape. There is one SOP per file starting with SOP 29, i.e. File 1 corresponds to SOP 29, and File 572 to 
SOP 600. Within a file, each 80-byte logical record contains the timing, pointing and infrared data for a 
single eight-second average. Each logical record is identified with an OBS, but OBSs are not otherwise 
separated from one another within a SOP file. SOPs for which no survey data exist contain one dummy 
record which has a valid SOP number but zeros in all other fields. 

The ZOHF is available only as an unlabelled, ASCII formatted magnetic tape. The logical record 
length is 80 bytes with a physical block size of 200 logical records or 16000 bytes. The approximate size 
of the data base is 96 Mbytes. All of the fields contain data. There are no header or trailer records. 


X-61 



start 

byte 

name 

Table X.Ap4.1 Format of ZOHF 

description 

units 

type 

1 

NSOP 

SOP number 



13 

4 

NOBS 

OBS number 

- 

13 

7 

NUTCS 1 

Time UTCS 

sec 

110 

17 

INCL 

Inclination 

degrees 

F6.2 

23 

ELONG 2 

solar Elongation 

degrees 

F6.2 

29 

BETA 

Ecliptic Latitude 

degrees 

F6.2 

35 

LAMBDA 

Ecliptic Longitude 

degrees 

F6.2 

41 

/v, 

12 pm Brightness Density 3 

Jy sr -1 

El 0.4 

51 

/v 2 

25 pm Brightness Density 

Jy sr -1 

El 0.4 

61 

/v, 

60 pm Brightness Density 

Jy sr -1 

E10.4 

71 

'v. 

100 pm Brightness Density 

Jy sr -1 

El 0.4 


'UTCS is elapsed time in seconds since 0 hours UT, 1 Jan 1981 

2 Elongation is the angle between the line of sight and the Sun. Inclination is the angle between the eclip- 
tic plane and the plane containing the Earth, Sun, and observation direction (i.e., the azimuth angle 
about the Earth-Sun axis). Elongation was fixed in each OBS, and inclination increased at a constant 
rate. These angles are related to geocentric ecliptic coordinates by the following expressions: 

cos(ELONG) - cos(p) cos(A.— X q) 

sin(/A r CL) - sin(p)/sin(EL0M7) 

where is the ecliptic longitude of the Sun. 

3 Conversion from in-band brightness to brightness density assumes a source with an energy distribution 
which is flat in flux per unit octave as explained in Section VI.C. 


Table X.Ap4.2 Table of Missing SOPs 1 


053 

258 

264 

054 

259 

442 

055 

260 

594 

056 

261 

595 

058 

262 

596 

200 

263 

597 


‘A dummy data record consists of a valid SOP number and zeros in all other data fields. 


X-62 





XI. KNOWN PROCESSING ANOMALIES 


In any undertaking of the magnitude of the IRAS data reduction, there are bound to be errors or 
software design deficiencies, which were either generic in the processing, discovered too late in the pro- 
cessing schedule to be fixed for all the observations or considered minor enough not to warrant fixing. In 
this chapter, all those errors which were recognized at the time of publication of the catalog (November 
1984) are listed. It is realized that this list does not encompass all the actual errors in the catalog. 

A. Processing of Extended ("Cirrus") Sources as Point Sources 

As discussed in Section V.C, the point source recognizer can be triggered by sources that are of 
unlimited extent in the cross-scan direction as long as they are less than about 1 in the in-scan direction. 
While this characteristic was recognized early in the design of the software, it was felt that this situation 
would arise much less frequently than the case where several true point sources triggered many cross-scan 
detectors simultaneously. The presence of "infrared cirrus” and the extended sources in the Galactic 
plane means that at 100 pm in particular the sky is dominated by extended structures rather than by 
clusters of point sources. The detection and confirmation software thus created strings of pseudo-point 
sources spread along the cross-scan direction. 

In retrospect, it would have been better to adopt a different approach. For example, whenever four 
or more cross-scan detectors were triggered, all such detections should have been discarded from the 
point source catalog or placed in a special data base to compare different scans to verify that these detec- 
tions were in fact due to several true point sources. The consequences of the present approach are clear. 
At 100 pm the point source catalog is unavoidably contaminated by the presence of point-like condensa- 
tions within the extended emission and by the effects of cirrus on discrete point sources that include miss- 
ing fluxes at shorter wavelengths caused by band-merging difficulties. 


B. Instability and Lag of the Noise Estimator 

The noise estimator (Section V.C.2) was picked for its computational speed and its performance in 
quiet areas of the sky. It performed poorly in regions of rapidly changing noise amplitude. As a result, 
the noise estimator near regions of high source density was a major reason for incompleteness for 
sources that would properly have had a signal-to-noise ratio of three or more because the estimator 
erroneously increased the noise level so that the computed signal-to-noise ratio fell below the threshold. 
Furthermore, the noise estimator exhibited a lag which resulted in the almost complete excision of 
sources within a degree or two on one side of the Galactic plane. 


C. Frequency Dependence of Responsivity with Amplitude 

The observations described in Section IV.A.4 established that flux densities greater than 10 Jy at 60 
and 100 pm are in error of at least 30% at 60 pm and up to 70% at 100 pm because the responsivity of 
the detectors varied with the total flux on the detectors. This effect depends on the background as well as 



on the brightness of the source itself. The photometry of any source in a background greater than 10 
MJy sr -1 at 60 and 100 pm is therefore suspect. 

D. Errors in Cross-Scan Uncertainties Related to Failed Detectors 

Several changes were made in the course of the processing to ameliorate the effects of failed 
detectors. While all of the changes resulted in overall improvements, the last change introduced an error 
as well. The last change handled the situation when an edge detection occurred opposite a failed detector 
and when the fluxes of the detections differed by more than a factor of two, so that the detections failed 
to seconds-confirm. Since only one detection could band-merge with the rest of the source, an improve- 
ment was made to select always the brighter detection and to delete the fainter one. An error in the 
software assigned a wildly incorrect cross-scan uniform uncertainty to the remaining detection one-third 
of the time. 

This change still resulted in a net improvement in the quality of the sources. The erroneous cross- 
scan uncertainties were immediately eliminated if band-merging was successful. A problem was created 
when the faulty detection was merged with another real source that had transited the focal plane at pre- 
cisely the same time. No restriction on cross-scan position existed because of the impossibly large cross- 
scan uniform uncertainty associated with the detection. The resulting, incorrectly band-merged source 
failed to have upper limits from being assigned to unobserved (or unmerged bands). Thus, the Working 
Survey Data Base (WSDB) contains some hours-confirmations with some bands with FSTAT-0, and 
detectors all 0 (see Section V.D.8). A detailed comparison of the processing of the data with and without 
this error showed that no sources with signal-to-noise ratio greater than 9 were lost due to to this error. 

E. Photon-Induced Responsivity Enhancement 

The observations described in Sections IV.A.8 show that a correction of as much as 20% should be 
applied to the photometry in the 100 pm band because of the photon-induced responsivity enhancement 
caused by passage over the Galactic plane. This correction depends on the detailed history of the indivi- 
dual scans which go into making up a source. The correction scheme described in Section VI.B.4 is only 
valid in a statistical sense and is based on guess-work for those areas which were not covered by ascend- 
ing and descending scans. It should be emphasized that these corrections were not incorporated into the 
instrumental fluxes recorded in the WSDB. 

No correction has been attempted for the extended emission maps. 

F. Artifacts in the Digital Image Data Base 

The final calibration corrections for baseline and scale factor in the 16.5° image digital data base 
delivered in November 1984 (third HCON coverage only) were based on the Zodiacal History File. 
There was insufficient time before the delivery date to perform the corrections properly on the time 
ordered data and reassemble the high resolution maps. As a result, the 2' digital data show a patchwork 
structure of 30' squares which varies in intensity from one part of the sky to another. The maximum 
discontinuity across the edge of one of the squares amounts to about 5 digitization levels of the satellite’s 

XI-2 


: 1 I i 



electronics in all bands (Section II.C) regardless of the instantaneous sky brightness. This translates to 
maximum jumps of about 0.9, 2.0, 0.5 and 1.0 MJy sr _1 at 12, 25, 60 and 100 pm, respectively. The 
problem should be of little concern in high brightness regions such as the Galactic plane but may 
compromise the data for some purposes in low brightness regions near the Ecliptic and Galactic poles. In 
most cases the patchwork amplitude is smaller than the residual striping left by the flat fielding pro- 
cedure. 

No attempt was made to smooth out the discontinuities. A number of fields where the problem 
was most severe were deleted and are listed in Table XI.F. 1. The calibration of the third coverage will be 
corrected properly and a complete data set reissued. The photographic data products released in 
November, 1984 do not suffer from the problem because no corrections were applied to these data as dis- 
cussed in Section X.D.3. The first and second HCON coverages as well as the reprocessed third HCON 
will be properly corrected. 


Table XI.F.l Fields Deleted from 3rd HCON 


7 

18 

19 

20 

21 

22 

23 

24 

36 

43 

44 

56 

58 

82 

92 

116 

127 

140 

141 

170 

173 

185 

186 

187 

188 

190 

191 

192 

201 

203 

205 

209 

210 

- 

- 

- 


G. Photometric Processing 

As stated in Section VII.D, the flux densities quoted for the secondary standard NGC 6543 differed 
from those assumed as inputs to the processing by up to 4%. The cause of these discrepancies is unknown 
at this time and the effect, although statistically significant, was considered too small, compared to the 
relative photometric uncertainties, to be remedied. 

H. Insufficient Specification of HCON Coverage 

In order to understand the reliability of a given source, a knowledge of the number of total HCONs 
possible is necessary. This information is not available in either the catalogs or the WSDB and can be 
obtained only by recourse to the raw data. 

I. Position Uncertainties 

As described in Section VU.C.l.b, the IRAS position uncertainties are significantly overestimated 
for brighter objects. An additional source of overestimation of the uncertainties for some objects was a 
processing error that added 3" in quadrature to the in-scan uncertainty for all sources with a cross-scan 


XI-3 



uniform uncertainty of l". It had been intended to do this only for faint sources, about 40% of the cata- 
log, but was inadvertently applied to sources of all flux levels meeting the uncertainty criterion, about 
60% of the catalog. 

J. Overestimated Weak Fluxes 

There is a systematic error associated with the measured flux densities, particularly those of 
moderate quality, near the noise limit of the detectors, which can increase the quoted value by as much 
as a factor of two compared with the true value. Weak sources were often detected when positive noise 
excursions pushed the source signal above the 3a signal-to-noise (SNR) cutoff imposed by the processing. 
Negative excursions dropped the signal below the SNR cutoff and so were not included in the flux 
averaging. While the problem is complicated by the uneven sensitivities of the detectors, a simple 
analysis demonstrates the problem and leads to a simple correction factor. 

Assume that all of the noise sources are Gaussian and neglect spikes and cosmic ray hits as 
second order effects (see Sec. VII. D of the Supplement). Define 

The noise level — a. 

The true source flux density, f lrue = no. 

The flux threshold - ma. 

The measured flux density, f ohs — n 'a. 

A source will be detected whenever the measured flux density exceeds the threshold, i.e. 
n' > m. For a perfectly uniform focal plane with no failed detectors the fraction of times, p, that a 
source was seen is given by: 


p-{2nT m j e~ yl/2 dy (XI.J.l) 

m—n 

The expectation value of the observed flux density is given by the average of all detected values, 

oo 

f ye~^~ n ^ l2 dy 

<n'> - (XI.J.2) 

/ e-(y- n)1/2 dy 

m 

Then the expectation value of the ratio of the observed flux density to the true flux density is given by 

j e -\/2 (m-n) 2 

<n'/n> — < observed/ true > - 1 H (XIJ. 3) 

m “ 

/ e~ y2/2 dy 

m 

The fraction of possible times the source was detected, p, (Eqn. XI.J.l) can be used to estimate 
the fractional amount, <«'/«> (Eqn. XI.J.3) by which the source’s flux density was overestimated. 


XI-4 



Table 1 gives the relationship between p and <n'/n> for two SNR thresholds, 2 a and 3<r, to bracket 
the range of sensitivities of the focal plane detectors. To use the table, one must determine from the 
Working Survey Data Base the number of times the source should have been seen and the number of 
times it actually was seen. The ratio of those two numbers is approximately equal to p which 
can then be used to estimate a correction factor to the quoted flux density using Table XI.J.l. 

The correction factor has been verified by using the moderate quality detections at 25 pm of 
hot stars with accurate 12 pm detections. At low flux levels the 12pm-25pm flux density ratio deviates 
from the Rayleigh Jeans limit. Application of this correction algorithm reduces this error consider- 
ably. 

It should be emphasized that this correction procedure ignores the effects of radiation hits 
and of extra sensitive detectors. For cases where very accurate photometry is required, one should co- 
add the survey data. 


Table XI.J.1 Relation between Fraction of Possible 

Sightings (p) and Flux Overestimate ( <ri/n > ) 


Flux Overestimate 

Frac. of 

SNR THRESHOLD 

Sightings 

(- 2a) 

(- 3a) 

P 

<n r /n> 

<n'/n> 

0.16 

2.51 

1.76 

0.31 

1.76 

1.60 

0.50 

1.40 

1.27 

0.69 

1.20 

1.15 

0.84 

1.10 

1.07 

0.94 

1.04 

1.02 


K. Minor Problems 

1) An unmodelled source of position errors was present in the data, as evidenced by the need to 
increase in-scan as for sources that had a cross-scan position error without a significant uniform 
component. 

2) Radiation hits and noise spikes caused some sources to have significantly larger cross-scan errors 
than quoted. 


XI-5 




3) In order to improve completeness, it was necessary to allow detections without a confirming partner 
due to a failed detector to bypass seconds-confirmation. This resulted in a large number of spurious 
detections being accorded a status equal to that of a truly seconds-confirmed source. These false 
detections could at times replace valid detections in band-merging. 

4) A second effect of failed detections was that an edge detection pair opposite a failed detector was 
often not seconds-confirmed, allowing two separate detections to be present at band-merging. The 
weaker detection, with its larger error basis, was usually chosen in band-merging resulting in a 
depressed flux in that band. 

5) A third effect of failed detectors resulted from detectors that were impaired but not dead. They 
would produce detections that were too weak to seconds-confirm with detections on the partner 
detectors. Again, the weaker detection would often be chosen in band-merging resulting in a 
depressed flux. 

6) Strong radiation hits could also result in a seconds-confirmation failure which, in a small fraction of 
cases caused incorrect fluxes or flux status. 

7) Responsivity changes due to particle or photon radiation caused baseline changes of order 10% in 
some observations. 

8) The optimum thresholds for accepting the seconds-confirmed sources were set by varying the thres- 
hold level and evaluating the numbers of sources passed as a function of threshold value (Section 
VILE. 6). The method worked satisfactorily in all the wavelength bands but the 25 |im band where 
the total number never reached the predicted plateau. 

9) The CIRR2 flag is incorrectly set to 0, implying no all sky data, for about 40 sources at the 0°-360° 
ecliptic longitude boundary. 

10) For many catalogs, the truncation instead of rounding of optical magnitudes leads to a 0.1 mag 
error in the reported magnitudes. 

11) For the ESO/Uppsala Survey of the ESO (B) Atlas, only the first two characters of the catalog 
description were used for the type field. 


Authors: 


T. Chester, G. Neugebauer, C. Beichman and T.N. Gautier. 



XII. ERRATA AND REVISIONS AS OF 1987 


The second release of some portions of the IRAS data reflects a number of changes, 
corrections and improvements to the data. This chapter, written in late 1986-early 1987, 
describes changes to each of the major datasets and records a number of changes to the formats 
of the data which will be of importance to people using the various products. The history of 
the releases of the IRAS data is described in Table XII.A.l. 


A. Version 2.0 of the Point Source Catalog 

The most important difference between the first and second versions of the IRAS Point 
Source Catalog (hereafter denoted PSC- 1 and PSC-2) is the application of a statistical correction 
to reduce the overestimation of the flux densities of sources near the detection threshold (see 
XI.J for a discussion of the effect); changes of as much as a factor of two were made for some 
weak sources. 

A number of other, less significant changes were also made. Flux densities for almost all 
sources changed by a few percent due to refined calculations of certain calibration factors; a few 
sources changed by as much as 10%. Corrections to errors in some calibration and confused 
source processing algorithms resulted in the loss of 6 sources from and the addition of 56 new 
objects for a revised total of 245,889 individual sources in PSC-2. Table XII. A.5 lists 12 
sources which were omitted from both PSC-1 and PSC-2 due to a software error, but which are 
of sufficient quality to be noted for completeness. A number of new catalogs were added to the 
list used for associations. Finally, a number of minor errors in calculations of some of the 
quantities associated with point sources, e.g. the number of neighboring sources, were corrected. 

A.l The Flux Overestimation Correction 

As described in detail in Chapter V, sources were extracted from individual detector 
streams by means of a zero-sum filter followed by a template fit to the data. Only if the detec- 
tion exceeded thresholds in its signal to noise ratio (SNR) and in its correlation with the tem- 
plate of a “perfect” point source was the detection accepted for subsequent processing. A 
source was accepted into the catalog only if it had enough valid detections to confirm its 
existence on time scales of seconds, hours and weeks. However, a failing of this strategy (as 
described in more detail in Section XI.J) is that the flux density of weak sources is overes- 
timated, since the combination of a weak source and positive-going noise (Gaussian excursions 
or spikes due to radiation hits) resulted in a “valid” detection, while the combination of a weak 
source and negative-going noise (Gaussian excursions) forced the detection below the acceptable 
thresholds. Thus, for weak sources, the detections used to form the average flux density 


XII- 1 



Table XII.A.l IRAS Data Products 



Status as of March 1987 



PRODUCT 

DESCRIPTION 

MEDIUM 

VERSION 

DATE 

Explanatory 

Detailed description 

Book 

1.0 

Nov. 1984 

Supplement 

of hardware, data 
processing, and products 
Revisions, New Chap XII, 

Book 

2.0 

Jun. 1987 


Index 




Point Source 

245,839 point sources 

Tape, 

1.0 

Nov. 1984 

Catalog 

(PSC-1) 

Microfiche, 

Tape 

1.1 

Jan. 1986 


245,889 point sources 

Tape 

2.0 

Nov. 1986 


(PSC-2). Updates described 
in this Chapter 

Book 

2.0 

Jun. 1987 

Ancillary File 

More detailed information 

Tape 

2.0 

Nov. 1984 

on point sources 


2.1 

Jan. 1986 



3.0 

Nov. 1986 

Working Survey 

More detailed information 

Tape 

2.0 

Nov. 1984 

Data Base 

on point sources 


3.0 

Feb. 1986 



4.0 (PSC-2) 

Nov. 1986 

High Source 
Density Bins 

Bins processed by high 
source density processor 
for Catalog 

Tape 

1.0 

Nov. 1984 

Point Source 

372,774 rejected point 

Tape 

available at IRAS 


Reject File 

sources 


data centers only 


Reject File 

More detailed information 

Tape 

available at IRAS 


Ancillary File 

on rejected point sources 


data centers only 


Reject File 

More detailed information 

Tape 

available at IRAS 


Working Survey 
Data Base 

on rejected point sources 


data centers only 


LRS Spectra 

Spectra of 5,425 catalog 

Tape, Hard 

1.0 

Nov. 1984 

Catalog 

point sources (8-22 
pm). See XII.C 

Copy 



Zodiacal History 

Time-ordered data at 

Tape 

1.0 

Nov. 1984 

File 

0.5° resolution 


2.0 

May 1986 

All Sky Maps 

All sky images at 

Tape 

1.0 

Nov. 1984 

0.5° resolution 

Negatives 




XII-2 




Table XII.A.l IRAS Data Products 



(Continued) 



PRODUCT 

DESCRIPTION 

MEDIUM 

VERSION 

DATE 

Sky Brightness 

16 x 16° field 

B/W, Color 

1.0 

Nov. 1984 

Images (HCON 3) 

2' Resolution 

Negatives 


(Sept-Nov)" 



Tapes 

1.0 

Nov. 1984 



B/W 

2.0 

May 1986 



Negatives 


(Dec. 1985)" 



Tapes 

2.0 

May 1986 





(Mar. 1986)" 

Sky Brightness 

16 x 16° field 

B/W 

1.0 

Aug. 1985 

Images (HCON 1 ) 

2 ' Resolution 

Negatives 


(Feb- Apr)" 



Tapes 

1.0 

Aug 1985 





(May 1985)" 

Sky Brightness 

16 x 16° field 

B/W 

1.0 

June 1986 

Images (HCON 2) 

2 ' Resolution 

Negatives 


(Apr 1986)" 



Tapes 

1.0 

June 1986 

Sky Brightness 

16 x 16° field 

B/W Negatives 

1.0 

Nov. 1984 

Images (Overlays) 





Galactic Plane 

2° x 15° field 

B/W 

1.0 

Jan. 1985 

Images (HCON 3) 

2 ' resolution 

Negatives, 

2.0 

July 1986 



Tapes 



Galactic Plane 

2° x 15° field 

B/W 



Images(HCON 1) 

2 ' resolution 

Negatives, 

1.0 

Oct. 1985 



Tapes 



Galactic Plane 

2 ° x 15° field 

B/W 



Images (HCON 2) 

2' resolution 

Negatives, 

1.0 

Jul. 1986 



Tapes 



Cataloged 

1 1 ,444 point sources 

Book 

1.0 

Feb. 1985 

Galaxies and 

associated with cataloged 




Ouasars 

galaxies and quasars 




Pointed 

13,853 images with 0.25'- 1.0' 

Tapes 

1.0 

Oct. 1985 

Observations 

resolution, each covering 





~ 1 square degree 




Pointed 

Description of AO’s, 

Book 

1.0 

Nov. 1985 

Observations 

data reduction, and 




User’s Guide^ 

released grids 




Serendipitous 

Description of Catalog 

Book 

1.0 

Dec. 1986 

Survey Explanatory 





Supplement 





Serendipitous 

43,866 point sources 




Survey Catalog*" 

derived from the 

Tape, micro- 

1.0 

Dec. 1986 


Additional Observations 

fiche 




XII-3 




Table XII.A.l IRAS Data Products 
(Continued) 


PRODUCT 

DESCRIPTION 

MEDIUM 

VERSION 

DATE 

Small Scale 
Structure Catalog 
Explanatory 
Supplement 

Description of 
Catalog 

Book 

1.0 

Dec. 1985 
Mar. 1987 

Small Scale 
Structure Catalog 

16,740 sources 
with sizes < 8* 

Tape, 

Microfiche, 

Book 

1.0 

Dec. 1985 
Mar. 1987 

CPC Explanatory 
Supplement 

Detailed description of 
CPC hardware, data, 
processing, and products 

Book 

1.0 

Aug. 1985 

CPC data 

Chopped photometric 
channel image data 

3 tapes 

1.0 

Jan. 1986 

IRAS Asteroid 
and Comet Survey 

IRAS and derived data 
on known asteroids 

Tape 

1.0 

Oct. 1986 

IRAS Asteroid 
and Comet Survey 
Explanatory 
Supplement 

Description of Catalog 
and Summary Information 

Book 

1.0 

Oct. 1986 


“Internal product dates on tape versions show months in this range 
b Young et al. 1985 
c Kleinmann et al. 1987 


reported in PSC-1 were systematically high. This problem predominantly affected moderate 
quality fluxes of weak sources; brighter sources detected with high quality fluxes generally had 
enough detections that flux overestimation was not a large problem. 

The most direct way of measuring the overestimation is to compare the brightness of 
sources in PSC-1 with the values obtained from the more sensitive pointed mode of the satellite 
(Young et al. 1985). Observations made in the pointed mode were three to five times more 
sensitive than the scans making up the survey so that the sources detected in the pointed mode 
would not suffer from threshold effects at the same flux level. The Serendipitous Source Cata- 
log (Kleinmann et al. 1987, hereafter denoted as the SSC) is the catalog of sources extracted 
from selected pointed mode observations and was used as a “truth table” to determine the mag- 
nitude of the flux overestimation effect and to help develop a correction for it. Figures 
XII. A. la-4a show the ratio of SSC to PSC-1 flux densities for the four bands as a function of 
SSC flux density; the presence of an overestimate in the PSC-1 values below about 2 Jy is obvi- 
ous. 


XII-4 


I II: 










LN (F v (SSC)/F „ (PSC 2)) LN (F „ (SSC)/F „ (PSC I)) 



LOG (F„ (SSC)) Jy 



Figure XII. A.3 Ratio of SSC to PSC-1 (a) or PSC-2 (b) 60 nm flux densities v: 
sity before (a) and after (b) the correction of the overestimation 


SSC flux den 
effect. 







A simple model for the overestimation described in Section XI.J leads to an algorithm for 
the correction of the effect. Consider a detector stream with an intrinsic Gaussian noise level a 
and a source with an intrinsic brightness of no. The zeroth moment of the Gaussian distribu- 
tion above the threshold of mo for a source with true flux no gives the probability of detection 
(Eqn. XI.J. 1) and the first moment above mo gives the observed flux (Eqn. XI.J.2). These 
moments were used to relate the flux correction factor to the observed fraction of possible sight- 
ings for a given source. 

The algorithm was implemented in the following manner. The Working Survey Data Base 
(WSDB) contains the number, N, of all accepted detections. If the total number of possible 
detections is denoted by M, then the ratio N/M is an estimate of the probability of detection, 
denoted by p, which can be used in Equation XI.J.l to estimate the quantity n—m. This value 
can be used with Equation XI.J. 3 to derive the ratio n'/n which is an estimate of the amount 
by which the true brightness of the source was overestimated in PSC-1. 

For each source fainter than a certain, band-dependent level (2 Jy at 12, 25 and 60 pm 
and 3 Jy at 100 pm), the number of accepted detections, N, was determined from the WSDB. 
At the same time the total number of possible detections, M, was determined from a detailed 
calculation of the satellite’s pointing history over the entire mission. If the source passed within 
the central portion of a working detector then a possible sighting was recorded. The central 
portion of the detector was defined by a distance in from the edge of the detector and was con- 
sidered a free parameter, denoted by 8Z, in each band; the value of this parameter affected the 
number of possible detections for a given source.The resultant value of N/M was used along 
with an estimate of the cutoff threshold, m, in that band to derive the correction factor as 
described above. 

The SSC data represented in Figure XII. A. la-4a were assumed to give the true brightness 
for faint sources and were used to adjust the parameters 5Z and m on a band by band basis to 
make the curves as flat as possible. The algorithm was applied only to those sources which 
were neither confused nor located in regions of high source density in a given band. Approxi- 
mately 100,000 sources were adjusted in at least one wavelength. Values of 5Z and m and the 
maximum derived correction factors are given in Table XII.A.2. 

Table XII.A.2 Flux Overestimation Parameters 


Wavelength 

(pm) 

m 

(SNR Threshold) 

8Z 

(') 

Max. Correction 

12 

4 

0.2 

1.72 

25 

2 

0.4 

3.00 

60 

3 

0.7 

2.01 

100 

3 

1.0 

2.38 


XII-9 




The photometric uncertainty associated with each flux was also adjusted slightly by exa- 
mining how uncertainties in the pointing of the telescope could affect the number of possible 
sightings. Variations in the N/M ratio were propagated through Eqs. XI.J.l and XI. J. 3 and the 
resulting uncertainty added in quadrature to the uncertainty already given in the catalog. 

The above formalism obviously ignores many complications such as differing sensitivities 
among the detectors within a band and non-Gaussian noise sources like radiation hits and 
cirrus which could increase a source’s flux on some occasions, but not others. Detections could 
have fallen below thresholds for any number of reasons including sightings by a detector of 
poor sensitivity or passage over the edge of a detector rather than over its center. A sighting 
could have been rejected for failing to exceed either the SNR or template thresholds. All of 
these effects were ignored in the above algorithm. As a result, although the corrections are 
quite good in a statistical sense, they may be considerably in error in any individual case. 

Despite these simplifications, the algorithm worked remarkably well. Figure XII.A.lb-4b 
shows the comparison of the PSC-2 and SSC fluxes. While some structure including a slight 
offset in some bands is still evident in the data, the curves are much flatter than those in figure 
XII.A.la-4a; the offset may be due to residual uncertainties in the calibration of the photometry 
in PSC-2 and the SSC. 

Figure XII.A.5 shows the effect of the flux overestimation on the photometry for weak 
sources in PSC-1 and the effectiveness of the correction applied in the generation of PSC-2. 
The figure plots the 12/25 pm color of stars as a function of their 12 pm brightness. In Fig. 
XII.A.5a the flux density ratio comes from PSC-1 and the 12 pm brightness from the SSC. 
Above about 2 Jy a distinct population of stars with colors characteristic of hot (>3000 K) pho- 
tospheres is apparent, but below 2 Jy the population of stars with photospheric colors disap- 
pears and all stars appear to have a 25 pm excess. This astrophysically puzzling result can be 
understood in terms of the 25 pm flux being systematically overestimated for faint stars. The 
corrected colors from PSC-2 (Fig. XII.A.5.b) reveals a distinct population of stars with normal 
photospheres at faint levels. 

As a result of this correction algorithm, the machine readable version of PSC-2 includes 
five new fields: MHCON, the number of possible hours confirmed sightings a source could 
have had and the four band dependent correction factors that were applied to the PSC-1 flux 
densities to produce the PSC-2 fluxes. These quantities are given only for those sources for 
which a flux overestimation correction was made. It should be noted that the value of 
MFICON for any given source could be uncertain by one or two sightings. 

The format of the Catalog tape has been revised to incorporate these new quantities. Table 
XII.A.3 lists the changes to the original format (Table X.B.l, page X-4). 


XII- 10 


UNCORRECTED 


ORIGINAL PAGE IS 
OF POOR QUALITY 



((w rf 91) /O^l) "d) OOl 


Figure XII. A. 5 Ratio of 12 ^m to 25 |^m flux densities from the PSC-1 (a) or PSC-2 (b) vs. 12 
|a.m flux density from the SSC before (a) and after (b) the correction of the 
overestimation effect. 



Table XII.A.3 Revised Portion of Catalog Tape Format 

(see Table X.B.I) 

Start 


Byte 

Name 

Description 

Units 

Format 

139 

MHCON 

Possible number of HCONs 

— 

12 

141 

FCOR 

Flux Correction Factors 





Applied (xl000) 





( 1 value per band) 

— 

414 

149 

SPARE 

3 Spare bytes 

— 

3A1 


A.2 Additional Flux Density Changes 

Flux densities for all sources in the PSC-2 were adjusted by a few percent to make them 
consistent with the final calibration described in Chapter VI. The differences arose because 
PSC-1 used a preliminary version of the all-sky intensity maps to correct for the non-linearities 
of the feedback resistor (Section VI. A. 5). Also, as a result of this change, minimum upper limit 
fluxes are now exactly 0.25, 0.25, 0.40, and 1.0 Jy, instead of varying by a few percent from 
these values. 

In two areas of the sky, from 0° to 29° ecliptic longitude at -90° to -67° latitude and at 60° 
to 65° latitude, flux densities reported in PSC-1 were factors of 1.076, 1.099, 1.013, and 0.967 
too high at 12, 25, 60 and 100 pm, respectively. Approximately 1% of all sources were affected 
by this error which has been corrected in PSC-2. 

A. 3 New and Deleted Sources 

A number of effects changed the total number of point sources present in PSC-2. First, six 
sources present in PSC-1 were deleted from PSC-2 because the small flux changes mentioned 
above pushed the fluxes for these sources below a threshold used by the high source density 
processor. Second, fifty six sources were added to PSC-2 because the very near neighbor win- 
dow was set incorrectly in parts of the production of PSC-1. Thus, the total number of sources 
in PSC-2 is 245,889. The names of the changed sources are given in Table XII.A.4. 

Finally, a software error, discovered after the release of PSC-2, led to the exclusion of a 
small number of sources from both PSC-1 and PSC-2. Examination of the Working Survey 
Data Base (WSDB) led to the discovery of sources with similar fluxes located within ±30" in- 
scan and ±90" cross-scan of one another. In each case only a single hours confirmed sighting 
was reported for each object. Evidently, the weeks confirmation software failed to confirm and 
merge these objects into a single source which would then have appeared in the catalog. 

XII- 12 


I 1 1 i 




Table XII.A.4 New or Deleted Point Sources 




New Sources 


01241-7332 

05265-6840 

08220-4404 

18373-0918 

04531-6708 

05299-6830 

08573-4718 

19514+4306 

04532-6710 

05303-6951 

09040-7402 

21097+4953 

04534-6657 

05305-6952 

09101-5100 

21099+4954 

05044-7012 

05341-6632 

09179-7034 

21314+5805 

05046-7014 

05371-6944 

09215-7028 

21329+5126 

05122-6829 

05374-6946 

09437-6034 

21345+5706 

05150-6629 

05375-6650 

10075-5747 

21445+5653 

05150-6631 

05392-6847 

10075-5824 

21446+5655 

05207-6636 

05393-6930 

10182-5742 

21478+5649 

05236-6702 

05448-6720 

12376-6122 

22126+6905 

05239-6939 

05486-7001 

13191-6245 

22299+6440 

05242-6940 

08003-5012 

18092+4419 

22510+7138 

05264-6730 

08005-5013 

18285-0830 

22512+7140 


Deleted Sources 


04526-695 1 

13046-6222 

20250+4316 

21503+5105 

04538-6952 

15574-0052 




A procedure was developed for finding these objects systematically in the WSDB. First, the 
WSDB was searched for all sources with only one hours confirmed sighting, located above an 
absolute galactic latitude of 10° and above an ecliptic latitude of -80° (to avoid the Large Magel- 
lanic Cloud). Second, sources detected only at 25 pm or only at 100 pm were eliminated to 
avoid contamination by asteroids or cirrus. This sample was examined for objects having com- 
panions within ±30" in-scan and ±90" cross-scan. Those objects with nearby companions and 
meeting the final catalog selection criteria (Sections V.H.2 and V.H.7.) were taking as real 
objects. A final position was derived by simply averaging the position from each sighting; final 
flux densities were computed by averaging the logarithms of the flux densities without using any 
weights. 

Forty-three objects were selected by this procedure. Examination of the raw data led to the 
elimination of seven of these 43 as being due to either cirrus or confused objects. Finally, to 
avoid any problems of source reliability or flux overestimation only sources with a flux density 
greater than 1 Jy in at least one band were selected for inclusion in Table XII.A.5. The format 
of the table is similar to that of the printed version of the catalog, as described in Section X.B.2. 
It should be emphasized that these sources appear only in the table and are not included in 
the printed or machine readable versions of either PSC-1 or PSC-2. 


XII- 13 





Table XII.A.5 

Bright Sources Missing from the Point Source Catalog with \b\ > 10° 

— Name — Gal Flux Density in Janskys N C 

a (1950) 5 a 5 Coord (Not Color-Corrected) C P S I A — Association — 


HHMMT DDMM 

(S) 

(") 

1 b 

NH 

12pm 

25pm 

60pm 

100pm 

1 

W 2 

D 

T 

-Name & Type- 

MAG 

00125-0723 

315 

10 

98-68 

2 

25L 

.44L 

1.46 

4.62 

0 


I 

10 

M-0 1-0 1-060 

999 

1 1434+2042 

26.3 

50 

234+74 

2 

.25L 

.46: 

3.19: 

6.45 

1 


2 

10 

M+04-28-050 

999 

12337+2616 

47,4 

51 

230+86 

2 

.52: 

.53 

4.50 

23.31 

0 

1 7 

6 

9 

U07772 

103 

21492+3716 

13.4 

I 

88-12 

2 

.36: 

.25L 

.48 

2.95 

15 

4 

1 

2 

DO 20951 

1 

22261+8025 

9.1 

59 

117+20 

2 

1.01 

.28L 

.40L 

1.63L 

7 


1 

13 

3746 K0 

67 

22308+4105 

50.3 

26 

97-14 

2 

2.05 

.54 

.40L 

6.60L 

3 


2 

13 

52130 K0 

77 

22324+4024 

28.4 

35 

97-15 

2 

4.43 

18.49 

29.10 

64.39 

12 


2 

27 

MKN 914 

999 

22325+4054 

32.0 

25 

97-15 

2 

.25L 

.26L 

1.08 

7.01 

8 

1 4 

1 

32 

X2232+408 

4 

22326+4031 

37.8 

9 

97-15 

2 

.4 1 L 

,25L 

1.10 

7.33L 

9 


1 

23 

DG 187 

999 

22376+2426 

36.8 

3 

89-29 

2 

.29L 

.25L 

1.05 

2.68 

0 






23019+3405 

57.7 

32 

99-23 

2 

.25L 

.25L 

1.45 

2.50 

0 






23132+2449 

14.4 

12 

97-33 

2 

.25L 

.25L 

.48 

1.35 

4 


3 

9 

U 12460 

155 


A.4 Revised Completeness Estimates for PSC-2 

The completeness of PSC-2 was investigated in two different ways. First, PSC-2 was inter- 
nally checked through analysis of the differential source counts. Second, PSC-2 was externally 
checked against the SSC which is complete to significantly fainter flux densities than PSC-2. 
Results from both methods corroborate the general statements made in Chapter VIII that the 
survey (PSC-1 or PSC-2) is complete, in unconfused regions of the sky, to 0.4, 0.5, 0.6 and 1.0 
Jy at 12, 25, 60, and 100 pm. 

PSC-1 and PSC-2 differ in the flux density level associated with a given completeness level. 
Figures XII.A.6a,b show the differential log Nf\ogS curves for 12 and 60 pm for PSC-1 and 
PSC-2 at high Galactic latitudes covered with two sets of hours confirming scans. Because of 
the correction for flux overestimation, the peak of the curves is broader in PSC-2 than in PSC- 
1. The shapes of the curves implies that PSC-2 has more weak sources than PSC-1, but that 
those sources are very incompletely represented, e.g., —10% completeness at 0.3 Jy at 12pm. At 
the same time, PSC-2 does not achieve a 90% completeness level until a slightly higher flux 
density than PSC- 1 . 

Tables XII. 6a, b summarize the flux densities at which a given completeness is reached 
derived from the source counts in PSC- 1 and PSC-2. Values are given for regions with both two 
or three sets of hours confirming scans (2 or 3 HCONs). To the extent that the flux overestima- 
tion algorithm was successful, the table for PSC-2 reflects the true completeness of the IRAS 
survey as a function of flux density. Note that while it is quite accurate to derive the complete- 
ness from the log V/logS 1 curves below completeness levels of about 50%, it is difficult to derive 
accurate estimates above that level. Thus the flux density values at completeness levels of 90% 
and 95% are indicated in the tables as uncertain. 







Table XII.A.6a 

Flux Densities for a Given Completeness Level 
(PSC-1 ) 



Flux Density (Jy) 


Completeness 


2 HCON sky 


3 HCON sky 

(%). 

12pm 

60pm 

12pm 

60pm 

95 

[.421* 

[.591 

[.38] 

[.52] 

90 

[.40] 

[.56] 

[.36] 

[.50] 

50 

.33 

.47 

.32 

.44 

10 

.29 

.40 

.29 

.39 

5 

.28 

.37 

.28 

.38 


* Numbers in brackets are uncertain 


Table XII.A.6b 

Flux Densities for a Given Completeness Level 
(PSC-2) 


Completeness 

(%) 

Flux Density (Jy) 

2 HCON sky 3 HCON sky 

12pm 60pm 12pm 60um 

95 

[.47]* 

[.60] 

[.46] 

[.60] 

90 

[.44] 

[.59] 

[.44] 

[.57] 

50 

.35 

.48 

.30 

.45 

10 

.24 

.38 

.23 

.33 

5 

.22 

.34 

.22 

.30 


* Numbers in brackets are uncertain 


The second way to evaluate the completeness of the survey is to use the SSC, which, as 
discussed above, reaches three to five times fainter than the survey. A comparison was made of 
all SSC sources at Galactic latitudes b > 30° and b < 50° (to avoid the Magellanic Clouds). 
The SSC sources were divided into stars [/ v (12pm) > / v (25pm) and / v ( 1 2pm) > / v (60pm)] 
and galaxies [/ v (l2pm) < / v (60pm)], in order to isolate the 12 and 60 pm bands as much as 
possible. A search radius of 120" was used to find counterparts in PSC-2. Tables XII. A. 7a, b 
summarize the completeness derived from the percentage of SSC sources that were found in 
PSC-2 for both the 2 and 3 HCON sky at 12 and 60 pm. Given in the table are the number of 
actual objects in PSC-2 (denoted A r ') compared with the possible number of objects (denoted 
M') found in the more complete SSC. 





As a check on the consistency of these two approaches and on the agreement of the results 
in the 2 and 3 HCON sky areas, all the above numbers can be converted to an estimate of the 
completeness of a single HCON and plotted on the same graph. To do so, recall from Chapter 
VIII that in terms of the single HCON completeness C (= 1 — p ), the completeness of the PSC 
in the 2 HCON sky is C 2 and in the 3 HCON sky is C 2 (3 - 2C). 

Figures XII. A. 7a, b show the single HCON completeness at 12 and 60 pm. There is good 
agreement between the logV/logS 1 and the SSC vs PSC-2 results. In addition, the numbers 
from the 2 HCON sky and the 3 HCON sky are also consistent. However, the 12 pm results 
are significantly different from the results shown in Figure VIII. D.l derived from sources in the 
minisurvey. Because those sources were in the 7 HCON sky, they suffered tremendously from 
the effects of flux overestimation. This results in a shift of the minisurvey completeness curves 
to erroneously higher flux densities. Table XII.A.8 summarizes the Catalog completeness at 12 
and 60 pm as derived from the curves shown in Figure XII.A.6a,b. 


Table XII.A.7a 


12 pm Completeness from SSC 


Flux density 
(Jy) 

N' 

2 HCON sky 
M 1 Completeness 

N' 

3 HCON sky 
M' Completeness 

0.50-0.55 

6 

6 

1.00 

16 

16 

1.00 

0.45-0.50 

4 

5 

.80 

18 

18 

1.00 

0.40-0.45 

13 

15 

.87 

17 

17 

1.00 

0.35-0.40 

7 

11 

.64 

15 

15 

1.00 

0.30-0.35 

8 

17 

.47 

15 

23 

.65 

0.25-0.30 

6 

24 

.25 

24 

39 

.62 

0.20-0.25 

3 

26 

.12 

4 

50 

.08 




Table XII.A.7b 




60 pm Completeness from SSC 


Flux density 

N' 

2 HCON sky 

3 HCON sky 

(Jy) 

M' Completeness N' 

M' Completeness 


0.60-0.65 

14 

17 

.82 

11 

11 

1.00 

0.55-0.60 

10 

13 

.77 

17 

20 

.85 

0.50-0.55 

14 

25 

.56 

20 

24 

.83 

0.45-0.50 

8 

22 

.36 

18 

29 

.62 

0.40-0.45 

5 

29 

.17 

14 

35 

.40 

0.35-0.40 

4 

36 

.11 

13 

66 

.20 

0.30-0.35 

5 

49 

.10 

10 

70 

.14 


XII- 17 









SINGLE HCON COMPLETENESS SINGLE HCON 



F y ( 60 pim) (Jy ) 


Figure XII. A. 7 a) The single HCON completeness versus 12 pm flux density. A smooth solid 
line has been drawn through the data. The other solid line is from the 2 
HCON SSC results, the dashed line from the 3 HCON SSC results, the solid 
points from the 2 HCON logAVlogS results, and the open points from the 3 
HCON logA/logS' results. Uncertain points are in parentheses. 

b) The single HCON completeness versus 60 pm flux density. The smooth 
solid line has been drawn through the data. The other solid line is from the 2 
HCON SSC results, the dashed line from the 3 HCON SSC results, the solid 
points from the 2 HCON logA/logS 1 results, and the open points from the 3 
HCON logA7log.S results. Uncertain points are in parentheses. 


TIT 


XII- 18 




Table XII.A.8 


Best Estimate of Completeness Level 
(PSC-2) 


Completeness 

(%) 

Flux Density (Jy) 

2 HCON sky 3 HCON sky 

12um 60pm 12pm 60pm 

95 

.46 

.65 

.41 

.58 

90 

.45 

.64 

.38 

.56 

75 

.41 

.59 

.34 

.51 

50 

.35 

.52 

.29 

.45 

25 

.29 

.45 

.25 

.38 

10 

.25 

.38 

.23 

.32 

5 

.23 

.34 

.22 

.30 


A. 5 Associations 

The PSC-2 contains associations with sources in four additional catalogs not used in the 
preparation of PSC-1: the IRAS Small Scale Structure Catalog, the IRAS Serendipitous Survey 
Catalog, the OSU catalog of radio sources and the Michigan Spectral Catalog. Parameters for 
these new catalogs are given in revised versions of Tables V.H. 1 and X.B.4. These catalogs were 
added at the end of the queue for the printing priority in the printed version of the PSC. How- 
ever, in the printed version of PSC-2, an asterisk in the column giving the number of 
associations (NID) denotes an association of a PSC-2 source with an object in the Serendipitous 
Survey Catalog. 

There was an error in associating PSC-1 sources with stars in the Gliese catalog since the 
proper motion in right ascension was erroneously taken as seconds of arc per year instead of 
seconds of time per year. Remedying this error resulted in the addition of 37 associations in 
PSC-2 and the deletion of five associations originally presented in PSC-1. The changed associa- 
tions are given in Table XII.A.9. 

The ESO/Uppsala Catalog, catalog 14, was previously given a "multiple" classification 
instead of "other”. This has been corrected in PSC-2. 


XII- 19 





Table XII.A.9 Changed Gliese Associations 



Newly 

Associated Sources 


00027-3737 

01365-1812 

06562-4413 

16451-4737 

00156+4344 

01416-1611 

08053+6952 

16564+4726 

00176-6509 

01504-2240 

11027+4347 

17023-0459 

00235-7731 

02085-5103 

13275+1038 

20017+2312 

00295+6657 

02334+0639 

13432+1508 

20542-4419 

00348-2502 

03168-6245 

13469-2151 

21141-3904 

00461+5732 

03172-6241 

14006-4634 

21598-5700 

00589+7124 

03180-4315 

14260-6227 

22070-0452 

01051+5439 

05100-4502 

14359-6031 

23029-3607 




23110+5653 


Sources No Longer Associated 


07483+8023 

08355-3958 

16267+1831 

19145+0505 

21362-2732 





The following minor changes exist in the associations with all catalogs for sources in the 
area bounded by ecliptic longitude 260°-280° and ecliptic latitude -60° to 60°. The truncation 
error mentioned in Supplement XI. K. 10 has been fixed. One new association with the catalog 
of Suspected Variables, catalog 16, was made exactly at the maximum 90” radius allowed 
(source 18237-2417). The latest version (1984) of the SAO catalog was used in the generation 
of PSC-2; six new associations were made and eight old associations were lost. These sources 
are given in Table XII. A. 10. All 14 of these associations were at the maximum radius allowed. 



Table XII.A.10 

Changed SAO Associations 



Newly Associated Sources 


17167-3229 

17370-3843 

18311-1734 

18402-7755 

17326-3324 

17510-3726 




Sources No Longer Associated 


16516-6705 

17337-0220 

17458-0937 

17555+3324 

17236-2125 

17394+2611 

17519-3035 

17557+3351 


A. 6 Source Names 

Source names were derived from the equatorial position by taking the hours, minutes and 
tenths of minutes (truncated, not rounded) of right ascension as well as the sign, degrees and 
minutes of declination. For example a source at a = 12 /! 22" 1 15.5* and 8 = — 15°2015" has the 
name 12222-1520. Separate sources which would have been given the same name based on the 
above scheme were distinguished from one another by appending letters of the alphabet to their 





name. It should be noted that since sources were named before the decision to retain them in 
the catalog was made, it is possible that not all of the sources with names distinguished only by 
a letter will be present in the Catalog, e.g. 12222-1520B might be in the catalog, but not 1222- 
1520A. 

A. 7 Revised Positional Uncertainties for Bright Sources 

As described in the Section Vll.C.l.b, an additional 3" was supposed to be added in qua- 
drature to the in-scan uncertainty of sources that were both faint and had the minimum possi- 
ble (l") in-scan uncertainty. That correction was inadvertently applied in PSC-1 to all sources, 
faint and bright, with the minimum uncertainty. PSC-2 corrects this problem and has 84,287 
sources with this increased uncertainty, as opposed to 144,403 sources in PSC-1. The 
minimum uncertainty ellipse is now 3 " x 5 ", as opposed to 3 x 7 . 

A.8 Correction of Point Source Neighbor Counts 

In PSC-1 the counts of hours and weeks confirmed neighbors, denoted PNEARH and 
PNEARW, of a given source are in error for objects near which a forced weeks confirmation 
took place (Section V.H.3). This problem has been fixed for PSC-2. 

A.9 Spurious 25 pm Only Sources 

The following three sources have been found to be spurious on the basis of examination of 
co-added detector data, but were left in the catalog since no general rule could be used to delete 
them. They are: 05570-6722, 08291-6146, and 18021+6556. The first and third sources were 
produced by noise and radiation hits; the second resulted just from noise. 

A. 10 Working Survey Data Base, Ancillary File and Reject Files 

The version of the WSDB corresponding to PSC-2 is 4.0 (Nov. 1986) and differs from 3.0 
only in the revised number of sources in the catalog described above (56 additions, 6 deletions). 
Version 3.0 (Feb. 1986) corrected a calibration error in the flux densities of individual HCON 
sightings. Version 2.0 was the first publicly released version of the WSDB. 

Two errors were remedied in Version 3.0 of the Ancillary file. First, the "fault" byte in 
HSDDROC (see Table X.B.7c) was incorrect for a small number of sources due to the incorrect 
very near neighbor window. This is now corrected. Second, the first publicly released version 
of the Ancillary File (version 2.0) contained associations with the weeks-confirmed file for the 
Small Scale Structure Catalog: only about half of these weeks-confirmed sources survived to 
the final SSS catalog. The hex coding of the bands in which an extended source was detected 


XII-21 



was calculated incorrectly. These problems have been fixed in version 3.0 of the Ancillary File. 

Listings of point sources that failed the catalog selection criteria, e.g., confusion or 
insufficient number of HCONs, were not publicly released, but are available from the IRAS 
data center in Pasadena. 


B. Total Intensity Data 

B. 1 Total Intensity Maps 

A number of fields were omitted from the first release of the third HCON sky images due 
to the misapplication of a calibration factor (Section XI.F.l). This error has been corrected and 
the second release of the third HCON includes all the possible fields. 

B.2 Version 2.0 of the Zodiacal Observation History File 

The Zodiacal Observation History File (ZOHF) was recreated to correct a 0.5° in-scan 
position error found in Version 1.0. Several additional improvements were incorporated at the 
same time. A calibration correction that partially adjusts for errors in the electronic baseline 
was added. The sampling interval was decreased to 4 seconds of time (—1/4°) from 8 seconds 
of time (—1/2°). The 8 second (—1/2°) boxcar filter size remains the same. All records in Ver- 
sion 2.0 represent 8 seconds of contiguous data, whereas in Version 1.0, non-contiguous data 
were incorporated in the file when possible. Data from the edge detectors were not used in the 
new file. This means the averaged values represent, in cross-scan, 27. 3' for the 12 pm 
wavelength band and 28.5' for the 60 pm wavelength band. Though data from edge detectors 
in the 25 and 100 pm bands were also removed, the remaining full-sized detectors in these 
bands cover the full 30' width of the focal plane. 


C. Low Resolution Spectrometer 

No significant errors in the processing of the data from the Low Resolution Spectrometer 
(LRS) have been reported. A complete atlas of LRS spectra appears in Astronomy and Astro- 
physics (Supplement), 1986,65, 1. 



D. Other Anomalies Fixed In this Release 


The preceding chapter (XI) lists a variety of anomalies in the data processing that either 
were discovered too late in the processing to be fixed in PSC-1 or for which no fix was deemed 
possible. The status of each is detailed in Table XII.D.l. 



Table XII.D.l Status of Anomalies From First Release 

A. 

Processing of Cirrus 

No change 

B. 

Noise Estimator 

No change 

C. 

Frequency Dependent Responsivity 

No change 

D. 

Cross-scan Uncertainties 

No change 

E. 

Photon-induced Responsivity Enhancement 

Analysis in Progress 

F. 

Artifacts in Maps 

Fixed 

G. 

Photometric Processing of NGC 6543 

No change 

H. 

Unknown Number of Possible HCONs 

Partially fixed 

I. 

Position Uncertainties 

Fixed 

J. 

Flux Overestimation 

Fixed 

K. 

Minor Problems 9,10 

Fixed 

K. 

Minor Problem 7 

Analysis in Progress 

K. 

Other Minor Problems 

No change 


Authors: 

T. Chester, C. Beichman and T. Conrow 
References: 

KIeinmann,S.G., Cutri, R.M., Young, E.T., Low, F.J. and Gillett, F.C., 1986, Explanatory Sup- 
plement to the IRAS Serendipitous Survey Catalog. 

Young, E.T., Neugebauer, G., Kopan, E.L., Benson, R.D., Conrow, T.P., Rice, W.L., and Gre- 
gorich, D.T., 1985, JPL/IPAC Preprint 008. 


XII-23 




1 . 1,1 III Hi 


i 


i in 



XIII. CONTRIBUTORS TO IRAS 

Many people, agencies, and companies have contributed to the success of IRAS. Those listed below 
were among the many who have worked on the project since its inception in 1975. 

The US National Aeronautics and Space Administration (NASA) was responsible for the design and 
development of the telescope. NASA also provided the launch of the satellite by Goddard Space Flight 
Center (GSFQ with a Delta 3910 Rocket at the Western Test Range and ground tracking and communi- 
cation during the early weeks after launch. NASA was also responsible for the final data processing of 
the survey data. 

The Netherlands Agency for Aerospace Programs (NIVR) was responsible for the design and 
development of the spacecraft and the intergration of the satellite. NIVR also was responsible for the 
ground facility and services for the operation of the satellite while in orbit. 

The UK Science and Engineering Research Council (SERC) carried out the design, development 
and operation of the real time tracking and data acquisition facility which controlled the satellite, and 
supported the design, development and implementation of the non real time control software and the 
preliminary analysis facility at the Rutherford Appleton Laboratory (RAL). 

The overall IRAS project management was co-chaired by Peter Linssen and Gene Giberson. A 
series of managers have played important roles during the lifetime of IRAS. These include: W. Bloemen- 
dal, E. K. Casani, D. Compton, R. Dalziel, J. de Koomen, W. de Leeuw, E. Dunford, T. Harmount, B. 
Martin, G. M. Smith, and G. Squibb. 

Participating at the headquarters of the three agencies were: M. Bensimon, N. Boggess, J. Clapp, L. 
Dondey, B. Edelson, R. Halpem, C. Hartman, N. Hinners, J. Holtz, L. Jones, L. Kline, W. Logan, F. 
Martin, B. Norris, C. Pellerin, A. Stofan, D. Stoughton, and D. Wrublick (NASA); N. de Boer, E.F.F.M. 
Braun, D. de Hoop, A.P. Hoeke, and M. van der Matten (NIVR); and H.H. Atkinson (SERC). 

A joint science team defined the overall mission and saw that the scientific requirements were 
fullfilled. Initially the joint team consisted of H.H. Aumann, D.A. Beintema, N. Boggess, J. Borgman, 
P.E. Clegg, T. de Jong, F.C. Gillett, H. J. Habing, M.G. Hauser, J. R. Houck, R.E. Jennings, F.J. Low, 
P.L. Marsden, G. Neugebauer, S.R. Pottasch, B.T. Soifer, R. van Duinen, and R.G. Walker. G. Neu- 
gebauer and R. van Duinen served as American and European co-chairmen of the joint team. During 
the course of the project, B. Baud, C.A. Beichman, J.P. Emerson, T.N. Gautier, S. Harris, G. Miley, 
F.M. Olnon, E. Raimond, M. Rowan-Robinson, P.R. Wesselius and E. Young joined the science team; 
in 1982 H. J. Habing became the European co-chairman of the joint science team. 


Telescope Development 

The Jet Propulsion Laboratory (JPL), operated by Caltech, and the NASA Ames Research Center 
managed the telescope development. 

The participants in the project management activities at JPL were: M.J. Alazard, D. Bane, A.P. 
Bowman, F.E. Bristow, J.W. Clough, A.G. Conrad, S.L. Conrad, K.R. Cooke, M.Y. Cook, G.G. Coyle, 
W. Crosson, R.R. Dagelen, D.A. Elliott, R. Hernandez, G.J. Hodges, R.F. Klotz, K.G. Korkus, L.L. 
Lievense, D. Low, S.L. Molina, M.B. Murrill, H. Otake, G.K. Robinson, C.J. Silvio, C.W. Snyder, J.W. 


XIII- 1 



Stockeraer, R.F. Stott, J.J. Van der Woude, and R.H. White. 

The participants in the telescope system management at Ames Research Center were: D.L. Ciffone, 
D. Compton, P. Dyal, R.R. Nunamaker, G.W. Thorley, and L.S. Young. 

JPL personnel who participated in the telescope development and testing were: K. Ahlberg, B. 
Anderson, D.M. Anderson, B. Beidebach, T.J. Borden, L. Broms, B.D. Brown, G.G. Coyle, D.M. 
Engler, E.L. Floyd, J.A. Garba, J.R. Gatewood, A. Giandomenico, W. Gin, C.S. Guernsey, E.J. Heising, 
K.G. Holmes, L.E. Hovland, H.B. Hotz, W.P. Hubbard, H.W. Jackson, J.C. Koenig, J.J. Landeros, D.Er 
Langford, H. Lin, H.R. Long, R.A. Mallgren, D.P. Martin, P.V. Mason, S.D. Mayall, L.F. McGlinchey, 
F.A. Morelli, K.I. Moyd, F.L. Murphy, W.C. Neiderheiser, C.D. Newby, W.G. Orchard, H. Otake, T. 
Ozawa, J.D. Patzold, D. Petrac, S.W. Petrick, W.I. Purdy, B. Rax, J. Real, R. Richter, D.K. Rubin, 
W.M. Ruff, D.R. Rupnik, R.P. Salazar, C.N. Sasaki, K.B. Sigurdson, J. Slonski, J.F. Smith, Jr., J.B. 
Stephens, G.E. Tennant, E.F. Tubbs, E. Tward, H.D. Von Delden, L. Wen, J. Winther, J.L. Wolfe, C. 
Wong, E.C. Wong, W. Wood, and F.H. Wright. 

In addition, JPL personnel were added that specifically worked during the period of the focal plane 
redesign. These personnel were: R.H. White, A. Bailey, M.N. Carney, C.D. Carter, S. Chavez, C.T. Cru- 
zan, B.C. Debusk, C.G. Derkson, S. Dodge, D. Eastwood, H.L. Fitzhugh, R. Frazer, D.E. George, J.K. 
Hofman, R. Irigoyen, C.C. LaBaw, J.O. Lonborg, C.G. Lowell, M. Marquess, M.J. McKelvey, W. 
Powell, J.J. Simmonds, S. Thompson, A. Toppits, L.S. Vamell, and J. Vasbinder. 

At Cornell University, D. Briotta, P. Graf, and G. Gull helped develop the retrofit 100 pm detec- 
tors. 

At the University of Arizona, K. Armstrong, and R. Kurtz helped develop the 60 and 100 pm 
filters and the JFET modules. The spectral response of the Ge detectors and the long-term radiation 
effects on the Si and Ge detectors were also measured. E. Young devised the "bias boost” strategy. The 
internal reference sources, components of the spectral filters and the JFET moules were manufactured at 
Infrared Laboratories, Inc., Tucson, Arizona with the assistance of L. Richardson. 

Ames Research Center personnel who participated in the telescope system development were: G. 
Anderson, S. Baker, C. Ball, W. Barrows, T. Bridges, W. Brooks, D. Cusano, G. De Young, M. Dix, L. 
Edsinger, A. Femquist, W. Gilbreath, J. Goebel, T. Harmount, R. Hedlund, G. Hillen, E. lufer, P. Kit- 
tel, M. Kiya, R. Lavond, C. Leidich, K. Lorell, H. Lum, Jr., C. McCreight, E. Melugin, G. Nothwang, 
C. Neel, R. Pittman, L. Polaski, F. Prevost, J. Prucha, S. Rathjen, C. Ritchie, C. Robbins, E. Somer, W. 
Vanark, J. Vorreiter, R. Walker, T. Weber, C. Yetka, and E. Zimmerman. 

The definition study management team at GSFC consisted of E. Hymowitz, S. Mosier, and M. 
Mumma. 

The major companies participating in the design and development of the telescope were: 

Ball Aerospace Systems Division - Telescope System 
Rockwell International - Focal Plane 
Perkin-Elmer Corporation - Optics 

Personnel participating in the design and development of the telescope system at Ball Aerospace 
Systems Division were: T. Abbott, D. Adam, T. Alarez, L. Anderson, R. Anderson, L. Andreozzi, R. 


Arentz, H. Argue, J. Austin, J. Bamberg, C. Barkley, A. Bamer, R. Barnes, W. Beck, R. Bemis, W. 
Bengston, M. Berger, V. Berry, R. Bingham, K. Booth, R. Bradford, J. Briggs, J. Brown, J. Byrnes, W. 
Cash H Chameroy, C. Chambellam, J. Conlan, L. Cotsamire, J. Cowder, J. Cowley, J. Cox, F. Cutter, 
N. Daly, R. Darnell, D. Davis, T. Davis, D. Dehogg, L. Derouin, W. Deshler, W. Devereux, C. Downey, 

J. Droge, J. Duncan, G. Emerson, J. Emming, P. Encinias, D. Erickson, M. Erickson, B. Etheridge, D. 
Evans, B. Evans, L. Femer-Sinn, L. Fisher, R. Fisher, T. Fleener, B. Fox, H. Freeman, L. Frobom, A. 
Gabriel, J. Gallegos, J. Godden, D. Grabosky, E. Gray, R. Greenwall, R. Grunz, T. Hadsell, R. Haight, 
W. Hammel, R.D. Hamlin, R.J. Hamlin, S. Hanley, M. Harrell, K. Hegy, R. Herring, R. Hershey, 
Hicks D. Hillis, L. Hovda, R. Hopkins, L. Housewright, and L. Hughes. Others include P. Iverson, R. 
Jackman, G. James, S. Johnson, J. Johnson, J. Jones, M. Kerr, J. Kinsey, W. Lamb, K. Laughlin, T. U 
Blanc, D. Lennon, J. Lester, D. Livingston, D. Lloyd, E. Long, R. Loomis, P. Lostroh, M. Maine R. 
Manning, J. Marcantonio, J. Mateyka, D. Mathews, K. McDaniel, M. McKeever, V. McNeil, R. e 
Ion R Misch, A. Mord, W. Morley, G. Morris, D. Mount, H. Mynleiff, L. Nall, P. Nelson, M. Noble, 

P Olbert A. Olsen, C. Olson, R. Ortega, A. Pankaskie, D. Payne, J. Penner, K. Peterson, C. Pherson, H. 
Poehlmann, S. Porrine, M. Poyer, R. Price, P. Puzo, C. Rafferty, D. Regenbrecht, J. Reidy, R. Reinker, 
D Roalstad, J. Rodgers, J. Rodriquez, K. Roller, L. Rouse, C. Rowe, H. Royer, S. Scott, G. Seig , 
Smeins, G. Smiley, R. Snook, J. Sparks, N. Stoffer, D. Strecker, A. Stroeve, R. Sullivan, M. Swoboda, J. 
Tarpley J. Taylor, D. Tennant, J. Thomgren, B. Tolhurst, J. Tracy, A. Urbach, D. Van Gundy, L. Van- 
dello, S. Varlese, R. Viano, H. Vogt, J. Wanger, D. Warlick, R. Weary, H. Wells, R. Wendler, R. 
Werholz, W. Whitehead, W. Wickstrom, K. Willis, B. Wise, R. Wolfkiel, D. Wood, R. Woolley, R. 

Venne and N. Zaun. 

Participants at Rockwell International were: G. Audick, B.R. Bailey, E.C. Banks, D.J. Chiavenni, 
C.M. Cornelius, M.E. Dews, R.F. Flanagan, W.H. Flaugh, A.L. Gable, C.L. Hall, D.O. Hopwood, E.R. 
Hutchinson, C. Ju Wu, M.S. Keith, L.W. Kelsey, R.M. Lack, E. Lax, W.C. Milo J.C. Monson, D.G 
Moss, A.J. Nicoli, M.D. Petroff, V.S. Pottas, C.A. Randolph, D. Randolph, D.L. Rawlins, J.C. Roth, J. 
Savela, N. Sclar, L. Silverstein, W.E. Southworth, A.S Squillance, P.G. Tally, J.C. Torres, P.S. Vigneault, 
J.M. Walz, and J.V. Westling. 

Participants at Perkin-Elmer were: B. Alte, R. Altenhof, W. Arndt, J. Ashinasi, J. Bacich, T. Bar- 
letto, B. Boyce, W. Craighead, J. Dunn, G. Erdtmann, G. Fabich, G. Ferrera, D. Gabnel, L. Gardella, 
M. Gillen, F. Gillespie, S. Gowrinathan, R. Grosso, W. Gunther, H. Hall, N. Hamed, R. Hamed, G. 
Hawkins, A. Hellerung, G. Huse, R. Jones, R. Jung, D. Kittel, T. Konoski, G. Lester, E. Lotocki, C. 
McGlynn, M. McGuirk, J. Malloy, J. Mandle, D. Marshall, M. Matrullo, P. Mehta, H. Moeller, P. 
Naiden, B. Nardella, V. Nichols, R. Noll, J. Oberheuser, T. O’Neil, A. Osanitch, W. Papas, R. Paqum, 
W. Petrie, D. Phillips, C. Radcliff, R. Rowley, J. Russo, M. Schreibman, D. Shafer, F. Sileo, J. Stites, R. 
StoU, D. Stramiello, J. Tolan, D. Trost, N. Vancho, B. Voytek, A. Westfall, N. Woodbury, P. Young, 


and P.S. Young. 

The additional focal plane instrument (DAX) was designed, developed, manufactured, and tested by 
the Laboratory for Space Research, University of Groningen, with assistance from the Technical Physics 
Department of the Netherlands Organisation for Applied Research (TPD-TNO). Participants were J. 
Evers, A. de Jonge, W. Luinge, K. Wildeman, W. Werner, and J. Achthoven. 


XIII-3 



Spacecraft 


The spacecraft was designed and built by ICIRAS, an industrial consortium of Fokker and Signaal, 
the Dutch National Aerospace Laboratory (NLR) as a subcontractor. Participating at Fokker, Signaal 
and NLR were: W. van Es, R. Grijseels, P. Kant, H. Koot, J. Keijzer, W. Ligtenberg, C. de Leng, T. 
Olivierse, P. de Pijper, W. Pasteuning, G. Rietdijk, R. Schuren, J. Seeboldt, M. Vreeman, H. Vreeling, R. 
Waayer, D. Weber. 

Participants in the design and development at Fokker, Signaal, and NLR were: H. Bakker, W. 
Berkepeis, R. van Bezooyen, H. Buiter, J. Bisschop, P. Buurs, J. Field, A. Fok, R. Gibson, W. de Graaf, 
J. van de Heuvel, D. Jongeling, J. Kanis, A. Koekenberg, F. Koorevaar, G. van der Kruys, W. van der 
Laan, P. Lindhout, S. Maltaric, A. Mooij, J. Nolte, D. Noordhom, N. Pennings, J. Piebinga, F. Rekers, 
H. Rouws, W. Schenck, D. Schermer, B. Scholte van Marst, H. Sprengers, N. Smilde, M. Thijssen, P. 
Verveen, J. Wantsing, M. Wattel, J. van der Weerdt, C. van Wesel, H. Witte, J. Wijker, J. Wijnen. 

Apart from involvement in design and development, the following persons also prepared and car- 
ried out the spacecraft check-out and monitoring during operations: E. Boom, J. van Casteren, R. van 
Doom, M. van Dijk, G. Hameetman, L. Karsten, K. de KJuiver, M. Lamers, W. van Leeuwen, J. Prins, 
A. Pouw, C. P. R. C. Slippens, A. van Swieten, F. Teule, and C. van der Voort Maarschalk. 


Integration and Test 

The ICIRAS industrial consortium (Fokker and Signaal), assisted by NLR, had the prime responsi- 
bility for the satellite integration and test effort. Their Integration and Test team cooperated with that of 
JPL and with specialists from Ball Aerospace. The success in making this a unified effort was a major 
accomplishment of the IRAS project. 

Personnel supporting the integration and test of the satellite from JPL were: M.J. Alazard, G. Alba, 
D.M. Anderson, M.J. Argoud, J.C. Beckert, C. Beichman, D.J. Boatman, K. Breski, L. Broderick, B.D. 
Brown, R. Brown, R.W. Burt, L.V. Butler, N. Carte, W.J. Castellana, A.G. Conrad, G.G. Coyle, R. 
Daniel, F. Geno, D.C. Hammond, J. Harrel, E.J. Heising, J. Holbrook, W.R. Irace, J. Johnstone, J. 
June, F.L. Lane, T. Laney, P.V. Mason, J. MacConnell, J. Meehan, F.L. Murphy, E. Nave, J.D. Patzold, 
K. Port, D.L. Potts, D. Rosing, D.S. Ross, C.N. Sasaki, W. Schaefle, T. Shain, R.L. Sicol, A.G. Silliman, 
J.J. Simmonds, J.P. Slonski, S.T. Smith, L. Steimle, M. Temple, P. Van Velzer, J. Vasbinder, W. Walker, 
R. Weaks, C. Weidmann, R.H. White, V.A. Wirth, Jr., and W.R. Woods. 

Those individuals who supported the satellite test and integration from BASD were: L.C. 
Andreozzi, D.K. Chaffey, L.D. Davis, D.A. Durbin, R.O. Einertson, N. Erickson, E.C. Long, N.E. 
Magette, V.B. McNeill, R.M. Paris, W.E. Pinon, D.P. Runyon, C.A. Springer, J.L. Tucker, J.L. Tracy, 
Jr., A.R. Urbach, R.G. Voorhees, and J.A. Wells. 

Support provided by individuals at Fokker, Signaal, and NLR were: A. Bleekrode, C. Blom, R. van 
de Brink, J. Buesink, H. Carrington, H. Van Daalen, A. Van Dorsten, H. Hanekamp, D. van ’t Hof, L. 
Huesken, J. Klaaskate, H. Keppel, J. Kollen, W. Kollen, J. Lageman, T. Lamberts, M. Van Leeuwen, P. 
Moes, R. Nicolai, W. van Nifterick, R. Overgauw, T. Pasteuning, C. Schmeitink, A. Van Soolingen, B. 
Stenneke, J. van der Straaten, C. Taylor, S. Tensen, H. Vaassen, D. Van de Vegt, J. Verhey, A. Vialle, J. 



Webbers, P. Wielsma, J. Zuidam, and J. Zwarts. 


Launch 

The NASA Delta Project Office at the Goddard Space Flight Center and the Delta Launch Opera- 
tions of the Kennedy Space Center were responsible for the launch of the satellite. 

Ground Operations Software 

The software for the IRAS Ground Operations (IGO) was developed by a joint team from NLR and 
RAL under the management of R. van Holtz, R. Holdaway and A. Buck, and consisted of L. Baldwin, 

L. Barendse, I. Beharrel, A. Chipperfield, T. Dimbylow, D. Drummond, R. Ely, J. McDougall, B. Miller, 
K. Mount, E. Oord, P. de Pagter, W. Pol, M. Reid and H. Young. The onboard software was the 
responsibility of the NLR. 

The Preliminary Analysis Software, managed by G. Thomas, was written by J. Abolins, P. Barber, 
J. Fairclough, S. Green, S. Martin, M. Oliver, J. Renes, P. Richards, A. Stevens and B. Stewart. 

Operations 

The overall running of the mission was coordinated by the Mission Operations Manager, R. van 
Holtz, and his deputy A.J. Rogers. Operations at the Rutherford Appleton Laboratories, Chilton, 
Oxfordshire, were managed by H. Bevan, A. Buck, R. Holdaway, A. Lowe, J. McDougall, P. McPherson, 

M. Reid and G. Thomas. 

The generation of the IRAS Satellite Observation Plans was the responsibility of the IRAS Ground 
Operations Team consisting of J. MacDougall, L. Baldwin, J. Gourlay, S. Martin, E. Oord, W. Pol, G. 
Spalding and H. Young. 

Commanding of the satellite and receipt of the data were performed by the Chilton Operations 
Team led by D. Ewart, L. Harris, B. Rathbone, J. Stenning and J. Wright. 

Monitoring of the housekeeping data and a preliminary analysis of the scientific data were the 
responsibility of the Post Pass Analysis and Science Support Team consisting of J. Abolins, T. Dim- 
bylow, J. Fairclough, K. Mount, M. Oliver and P. Richards. H. Walker monitored the short-term pro- 
gress of the survey and modified the Observations Plans accordingly. 

J. MacDougall and B. Stewart had overall responsibility for maintaining the uplink and downlink 
software respectively. 

The Chilton Operations Support Team of H. Bevan, J. Cathrew, B. Champion, A. Chipperfield, A. 
Smith and G. Walkers maintained the ground hardware and provided support facilities. 

W. McLaughlin, S. Lundy, and D. Wolff designed and implemented the details of the survey stra- 
tegy with the assistance of C. Lau, H. Ling, R. Schlaifer and V. Wang. 

JPL personnel supporting operations at RAL were: A.I. Beers, B.D. Brown, J.B. Carraway, P.E. 
Doms, D.M. Engler, W.R. Irace, D.E. Langford, P.V. Mason, J.J. Rakiewicz, R.P. Salazar, C.W. Snyder 
and T. Sweetser. 


XIII-5 



Personnel supporting operations at JPL were: D. Bender, D.C. Bluhm, D.F. Finnerty, J.A. Holla- 
day, E.D. Jahelka, and RJ. Springer. 

Data Reduction 

The survey observations were reduced to their final form at the Science Data Analysis System 
(SDAS) facility at JPL. The final data processing system was developed under the management of J. 
Duxbury, T. Chester, P. Poulson, W. Scholey and W. Underwood. System design was by P. Poulson, D. 
Wittman and D. McCreary. Raw data input and conversion was by B. Sorensen and S. Pang. Calibra- 
tion was by D. Elliot and S. Wheelock. Data detection was by T. Hibbard. Pointing reconstruction was 
by H. McCallon. Confirmation was by J. Fowler, E. Rolfe and H. Pham. Small extended sources was by 
E. Rolfe and R. Pomphrey. Sky flux was by R. Stagner and R. Narron. Deep sky was by E. Kopan. 
Final product processing for the point source catalog was by T. Chester, C. Beichman, R. Beck, J. Ben- 
nett, J. Chillemi, N. Chin, T. Conrow, C. Oken, T. Sesplaukis and D. Wittman. Data base design was by 
T. Sesplaukis, C. Biller, J. Chillemi and F. Akers. Asteroid extraction was by S. Peters, T. Kia and J. 
Fowler. Simulation was by S. Kitzis and K. Moyd. User support software was by J. Bennett. Science 
and operations analysis by T. Chester and R. Benson. Image processing was by B. Hartley. 

Post processing and product analysis was by T. Conrow, J. Good, D. Gregorich, P. Hacking, G. 
Helou, S. Pohjala, C. Oken, W. Rice, D. Walker, P. Ramsey and K. Sedwick. 

Operations and technical support at JPL was accomplished by a team managed by J. Duxbury, T. 
Chester, R. Von Allmen, G. Smith and G. Lairmore. The data management team consisted of G. Smith, 
G. Lairmore, R. Beck, M. Alexander, H. Ashby, R. Bailey, S. Banks, S. Bedrossian, G. Benn, A. Bollin, 
M.R. Boykins, B.R. Brewer, I. Chan, J. Cisneros, T.J. Crawford, W. Currie, D.M. Engler, I.M. Esquivel, 
L. Fischer, D.E. Fritsche, L. Fullmer, G.W. Gameau, M. Garza, D. Hermsen, D. Hines, O. Hodges, 
L.M. Hughes, C.N. Hull, D. Hurwitz, D. Jackson, T. Jay, G. Johnson, B. Kaneshiro, B. Kahr, L. Lamp- 
ley, L.A. Lloyd, C. Lonsdale, P.J. Lynn, Jr., T.G. MacDufF, F. Macias, I. Miller, N. Parson, G.S. Pate, 
L.E. Perrine, W.H. Peters, B. Pondo, J. Radbill, P. Ramsey, L. Ray, D. Richardson, E. Robles, A.L. 
Sacks, M.S. Saul, K. Sedwick, B. Smith, A.E. Stansel, W. Starr, S. Taylor, N.W. Thomas, W.J. Trimble, 
R. Urban, B.M. Vance, J.E. Walsh, L.L. White, C. Wiley, J.L. Wilson, Jr., and C.A. Wolfe. 

The compilation of astronomical catalogs for the final processing was the responsibility of a group 
at GSFC consisting of J. Mead, T. Nagy and R. Hill. 

Pre-Publication Catalogs 

The following people and organizations outside the IRAS project provided catalogs on short notice 
and before publication: S. Kleinman and R. Joyce-catalog of accurate positions of sources from the Two 
Micron Sky Survey; M.P. Veron-Cetty and P. Veron-the Catalog of Quasars and Active Nuclei; the 
European Southern Observatory-the ESO/Uppsala Survey of the ESO(B) Atlas; J. Huchra-catalog of 
Seyfert galaxies; D. Gezari-the Catalog of Infrared Observations; and R. Hill-the most up-to-date version 
of the CLAS catalog. We also acknowledge the efforts of W. Warren and the Astronomical Data Center 
at the NASA Goddard Space Flight Center in providing many of the astronomical catalogs used during 
the IRAS mission and data processing. 


XIII-6 


l 1 1 



The Explanatory Supplement 

The Explanatory Supplement was assembled with the enthusiastic assistance of J. Boyd, S. 
Conrad, R. Dumas (Hernandez), S. Foster, R. Hernandez, S. Livingston, P. Neill, C. Race, J. 
Serpa. J. Boyd and R. Dumas (Hernandez) helped prepare the second edition of the Supple- 
ment. 


The Printed Version of the IRAS Catalogs 

The computer typeset versions of the IRAS catalogs were prepared by C. Oken, R. Beck 
and D. Turney. K. Simon oversaw the publication of the IRAS catalogs and the Supplement. 


XIII-7 




XIV. SURVEY SKY COVERAGE 


Plots of the sky surveyed by IRAS are given in the following pages so that the user of the catalog 
who fails to find a source at some position can verify that the position was, in fact, scanned enough times 
to result in a confirmable object. At least two hours-confirmed sightings were required for a source to be 
included in the catalog. A region of the sky was considered as a hole in one of the hours-confirming cov- 
erages if: 1) the telescope simply did not observe that region; 2) only failed detectors covered the region; 
or 3) the particle radiation level was sufficient to increase the detector noise by more than a factor of 
two. 

The broad overview of the sky coverage was given as Fig. I.C.l The more detailed maps presented 
here in equatorial coordinates show the departures from a "perfect coverage" which is considered to con- 
sist of three or more sets of hours-confirming scans. A region receiving three or more coverages shows no 
deviation from the baseline level, while a region that received only two coverages, i.e., a 'level 1 hole", is 
marked by a single-height box spanning the appropriate range of coordinates; a region that received only 
one coverage, a "level 2 hole", is marked by a double-height box; and a region that was not scanned at all 
is shown as a triple-height box. Regions with double or triple height boxes, i.e., having only single or no 
HCON coverage, do not or cannot contribute sources to the catalog. 

Coverage holes due to radiation effects are shown indistinguishably from geometrical holes in these 
plots. However, data for these regions were processed normally, so that sources may be found in some 
nominally forbidden regions (Section III.D.2). This applies particularly in the regions of the polar horns 
where radiation effects were occasionally severe enough to increase the noise level to qualify as a hole, 
yet not so bad as to prevent the detection of sources. 


XIV- 1 




67 . 
















ORIGINAL PAGE IS 
OE POOR QUALITY 


n i 























HOLES IN COVERAGE Of EQ. R.A. 60 TO 120 DEG 




m 





















HOLES IN COVERAGE OF EQ . R . A . 120 TO 160 DEG 



XIV- 17 


HOLES IN COVERAGE OF EQ . R.fi. 120 TO 160 DEG 







HOLES IN C 



XIV- 19 








XIV-20 


















HOLES IN COVERAGE OF EQ. R.FL 180 TO 240 DEG 



-23 . 
- 24 . 
- 25 . 
- 26 . 






















































HOLES IN COVERAGE OF EQ . R.A. 300 TO 360 DEG 


ORIGINAL PAGE IS 

oe poor quality: 



XIV-34 


10 . 





HOLES IN COVERAGE OF EQ . R . A . 300 TO 360 DEG 




HOLES IN COVERAGE OF EQ . R . A . 300 TO 360 DEG 















XV. INDEX 

Note: Some topics discussed in Chapter XII (Errata) may not appear in this Index. 

A 


absolute calibration 
ADC 

Aitoff projection 
all sky maps 
format 
alpha Lyr 
alpha Tau 
amplifiers 
analog data flow 
analog electronics 
analog signal path 
analog to digital conversion 
saturation 
extended source 
point source 
analysis 

extended emission products 
point source catalog 
ancillary file 
angular resolution 
annealing detector 
anomalies in processing 
aperture cover 
array 

artifacts in images 
associations 
catalogs used 
low resolution spectra 
asteroids 
calibration 
calibration model 
color temperatures 
source contamination by 
source density 
atlases and catalogs 
attitude calibration strategy 
attitude control 
constraints 

attitude reconstruction 
accuracy 
biased FACs 
errors 
absolute 

cross scan and in scan 
process noise 
thermal misalignment 
avoidance angles 


see under ’calibration’ 

see ’analog to digital conversion’ 

X-32 

V- 51 
X-37 

IV- 9, VI- 19, VI-21, VI-28 

VI- 19, VI-20, IX-6 

see ’electronics, analog’ 

11-22, 11-23 

II- 19ff, see also ’electronics, analog’ 
11-21 

11.25, VI- 1, 

III- 20 

V- 48 

X-7, X-16, X-26 

VII- 38ff 
Vll-lff 

X- 12, X-18 

1-1 

11-20 

XI- Iff 

II- 3, II-4 

see ’survey array’ 

XI-2 

VII-35, X-20ff 

V-64ff 

V-64 

V- 5, V-47 

VI- 22 

VI-19 

VI- 24 

VII- 34 

VII- 3 3 

see ’catalogs and atlases’ and ’catalogs’ 

III- 17 

II- 2 

III- 2 

II-2, II-2, V-3, V-6ff 

VII- 10 

V-8 

II- 2 
V-3 
V-7 
V-8 

III- 13, V-48 


XV- 1 


B 


Bamberga 
banana effect 
band filling, point source 
band merging 
point source 
small extended source 
threshold, small extended source 
band pass 

calculation of effective band pass 

effective frequency 

effective wavelength 

nominal 

spectral 

band, rejection of 
bandwidth 
baseline, electronic 
baseline stability 
bessel filter 
bias 

bias boost 
effects of 
biased FACs 
bin number program 
binning, extended emission 
bins in ecliptic coordinates 
black, optical 
blanking time, radiation 
block control word 

bright source neighbors removed as cross talk 
bright source problems 


C 


calibration 

absolute 

estimated accuracy 
extended mission 
point source 
uncertainty 

accuracy of relative calibration 
cross scan correction factor 
extended source consistency 
low resolution spectrometer 
problems 
stellar 
strategy 
attitude 


VT-22ff 

III- 8, III-9 

see ’confirmation, point source’ 

see under ’confirmation, point source’ 

see under ’confirmation, small extended source 

see under ’confirmation, small extended source 

11-17,18 

VI-20 

VI-27 

VI- 27 
X-13 
1-2 

VII- 26 

see ’band pass’ 

see ’detector, baseline’ 

see ’detector, baseline stability’ 

11-21 

11-20 

11-20, III- 13, VI-2, VI-6, VI- 12 

IV- 12, VI-2 

see ’attitude reconstruction, biased FACs’ 

X-46 

V- 50, V-73 
X-48 

n-8 

V-48 

X-12 

VTI-31 

VII-28ff 


Vl-lff 

VI-19ff 

VI-24 

VI-28 

VI- 20ff 

VII- 11 
VII-13ff 

VI- 6 

VII- 39 
IX -4 
VI-12ff 
VI-20 

III- 17 


XV-2 


photometric 
calibration equation 
catalogs 

association catalogs 
extragalactic 
low resolution spectra 
selection of spectra 
point source 

completeness and reliability 

description 

maps 

signal to noise cutoff 
source distribution 
source selection 
statistics 

sky brightness images 
small extended source 
analysis 

catalogs and atlases, description 

cautionary notes 

caveats 

extended emission 
extended source 
general 

high density regions 
low resolution spectra 
point source catalog 
small extended source 
working survey data base 
cavities 

charged particle hits 

Chilton, England 

chronology of mission 

CIRRI 

CIRR2 

CIRR3 

cirrus 

anomalous processing 
cirrus flagging 
cirrus flags 
class On 
class In 
class 2n 
class 3n 
class 4n 
class 5n 
class 6n 
class 7n 
class 8n 
class 9n 

classification of LRS spectra 
clean-up processing 
close neighbors 
cluster analysis processing 
clustering threshold 


III- 17 

VI- 4 

V-64ff 

1-3 

1-3, V-l 
IX- 12 
V-l 

VIII-4ff, XII-5ff 
1-3 

I-6ff, vn-40ff 
V-3 

1-4, VII-2 
V-64 

vn-i 

see ’extended source, map images’ 
1-3, V-l,V-46 

VII- 38 

1-3, see also under ’catalogs’ 
see ’caveats’ 


V-6 

V-51 

1-2 

V-62 

IX- 20 
V-54 
V-37 

X- 2 

see ’detector, integrating cavities’ 
see ’radiation hits’ 

III-7, V-3 
III- 18 

VII-37, X-9, X-15, see also ’cirrus flags’ 

VII-37, X-9, X-15, see also ’cirrus flags’ 

VII-37, X-9, X-15, see also ’cirrus flags’ 

1-2, V-4, V-54, VII-34, Vin-2, VIII-5 

XI- 1 
V-54 

V-54, VII-37 
IX- 19 
IX- 13 
IX- 13 
IX- 13 
IX- 13 
IX- 14 
IX- 14 
IX- 14 
IX- 19 
IX- 19 

see ’low resolution spectra, classification’ 

V-53, vn-l,x -16 
see under ’neighbors’ 

see under ’confirmation, small extended source’ 
see under ’confirmation, small extended source’ 


XV-3 



coarse window 
color correction 
color correction factors 
color map images 
comets 

commandable gain 
commandable offset levels 
communication 
completeness 
formalism for 
estimation of parameters 
in the galactic plane 
low resolution spectra 
outside of galactic plane 
point source 
small extended source 
verification of 
compressed detector data 
compression, dynamic range of sky prints 
compression scheme for digital data 
computer, onboard 
confirmation 

extended emission 
hours confirmation 
point source 
band filling 
band merging 

band merging bright sources 
double detection mode 
false alarms 
hours confirmation 
confusion processing 
decision 

photometric agreement 
photometric refinement 
position agreement 
position refinement 
statistical processing 
optical cross talk removal 
photometric refinement 
position probability density 
position reconstruction 
position refinement 
position uncertainty 
seconds confirmation 
confusion processing 
decision 

position refinement 
statistical processing 
threshold 

triple detection mode 
weeks confirmation 
decision 

position refinement 
statistical processing 


V- 18 

I- 2, VI-20, VI-27 

VI- 26, inside back cover 
see under ’map images’ 

V-5, V-47, VII-33 

II - 21 
11-21 
II-2 

m-i, vn-i, viii-iff 
vm-5 

VIII-7 

VIII- 10 

IX- 20 
VIII-8 

see under ’catalogs, point source’ 

VIII- 11 

VTII-9 

V-49 

X- 33 

n- 23 , n -25 

II-2 


1-2, V-6 

m-i 

V-3ff, V-13ff 

V-25 

V-24 

Vn-28 

V-19 

V-16 

V-4, V-28ff 

V-29 

V-28 

V-29 

V-29 

V-29 

V-15, V-22 

V-30 

V-18 

V-23 

V-15 

V-17 

V-15, V-22 

see ’position uncertainty’ 

V-4, V-16ff 

V-20fF 

V-18 

V-22 

V-23 

VII- 3 2 

V-19 

V-4, V-31ff 
V-32 
V-32 
V-32 


XV-4 


seconds confirmation 
small extended source 
band merging 
band merging threshold 
close neighbors 
cluster analysis processing 
clustering threshold 
decision function 
hours confirmation 
seconds confirmation 
source construction 
weeks confirmation 
weeks confirmation threshold 
weeks confirmation 
confuse flag 
confused neighbors 
confusion 

low resolution spectrometer 
point sources 
confusion flags 
confusion limit 
confusion noise 

confusion processing, point source 
confusion status word 
constraints 
attitude 

Earth IR radiation 
Earth horizon 
eclipse operation 
half orbit 
lune 
planet 

south Atlantic anomaly 
station passage 
Sun angle 
contributors 
control axes 
coordinate overlays 
coordinate transformation matrices 
cosmic ray hits 
cover, aperture 
coverage 
depth 

first HCON 
gaps 

second HCON 
third HCON 

cross scan correction factor 
cross scan direction 
cross scan response 

low resolution spectrometer 
cross scan uncertainty anomalies 
cross talk 

infrared detector 
optical 


V-39 

V-43 

V-44 

V-37 

V-40 

V-39 

V-5, V-36 
V-5, V-36 
V-36 

V-5, V-38 
V-42, V-44 
m-i 

VII- 37 
V-63 

I- 2, V-53, VUI-lff 

see under ’low resolution spectrometer’ 

VTII-2ff 

see under ’flags’ 

V-57 

see ’noise , confusion’ 

see under ’confirmation, point source’ 

V-20fF, V-33 

III-2 

m-2, III-4fF 
III-2, IV- 19 
III- 8 

III-2, m-17 
m-i7 

m -2 

m-3 

III-5 

m- 2 , iii- 5 

XIII- Iff 

II- 2 

X- 37 
V-17 

see ’radiation hits’ 

II- 3 

VIII- Iff, XIV- Iff 
I-4ff, III-22 
m-14 

m-19 

m -15 

III- 16 

see under ’calibration’ 

V-l 

IV- 3ff, VI-6 

see under ’low resolution spectrometer’ 

XI- 2 


IV-3 

VII-28, X-16 


XV-5 



in-flight tests 
removal 
cryogenics 
CSTAT 

cylindrical projection 


IV- 19 

see under ’confirmation, point source’ 
II-3, II-3ff 

see ’confusion status word’ 

X-31 


D 


data compression 
data flow, analog 
data processing 
analysis of 
extended emission 
final steps 

low resolution spectrometer 
point source 
point source catalog 
small extended source 
cluster analysis processing 
small extended sources 
data processing summary 
data reconstruction 
data, housekeeping 
dead detectors 
debris 

deleted detectors, extended emission 
deleted fields in first release of 3rd HCON 
density of sources 
depth of coverage 
despiking 

low resolution spectrometer 
survey array 
destriping 
detection 
detector 
amplifier 
annealing 
anomalies 
baseline 

particle radiation effects 
baseline stability 
bias 

construction 
cross scan response 
cross talk 
failed 
effects of 
infrared 

integrating cavity 
linearity 

low resolution spectrometer 
noise 
coherent 


11-23, 11-25 
II-22ff 

VIM If 

V-5ff, V-48ff 

V-6 

IX-9ff 

V-lff 

V-52ff 

V-5 

V-37 

V-34ff 

V-I 

V- 3 

II- 23 

see ’detector, failed’ 

III- l 

see under ’extended emission’ 

XI-3 

VIII- 2 

see under ’coverage’ 

IX- 10 

see ’nuclear pulse circumvention circuitry’ 
see ’extended emission, destriping’ 
see ’source detection’ 

see ’electronics, analog’ 

11-20 

IV- 3 

VI- 2 

rv-i 2 

IV- 11 

11-20 

11-14 

rv-3ff 

IV-3 

IV-3 

VTI-32, XI-2 
II- 11 

11-11, 11-14 
IV-3 

IX-3, IX -4 
11-20, IV-3 
IV-3 


XV- 6 



particle radiation effects 
noise spectrum 

photon induced responsivity enhancement 
anomalous processing 
extended emission 
low resolution spectrometer 
maps 

point source 
reliability 
responsivity 

non-linearity processing 
particle radiation effects 
saturation 
solid angle 
stability 
visible 
detector 5 

coherent noise 
detector 9 
detector 13 
detector 17 

detector 19, coherent noise 

detector 20 

detector 25 

detector 26 

detector 28 

detector 36 

detector 40 

detector 41 

detector 42 

detector 43, coherent noise 
detector 44 
detector cavities 
detector masks 
detector noise 
detector numbers 
detectors 

diffraction pattern 

diffraction spikes 

digital data compression scheme 

digital electronics 

dimensions of satellite 

distribution of catalog sources 

dose, photon 

double detection mode 

drop dead detection 

dust particles 

dynamic range compression 


IV- 12 
IV- 3 

IV-12ff, V-52 

XI-2 

VI-19 

IX- 8 
VI-16ff 
VI-13ff 

rv-3 

11-20, IV- 1, IV-3, IV-9 

XI- 1 

IV-llff 

m -20 

IV-3, IV-8 
IV-3 

II-11, 11-23 

IV- 3 

VI- 3 

VII- 32 
VII-32 

V- 31, VII-32 

VI- 3 

V- 31, VII-32 
IV-3, vn-32 
IV-3, vn-32 
IV-3, VII-32 

VII- 32 
VII-32 
VII-32 
IV-3, vn-32 

VI- 3 

VII- 32 
II- 11 
II- 11 

see ’noise, detector’ 

X- 25 

see ’detector’ 

IV-19 
IV- 19 
11-23, 11-25 

11-23, see also ’electronics, digital’ 

II- 1 

see under ’catalogs, point source’ 

IV- 13ff 

see under ’confirmation, point source’ 

V- 18, V-24 

III- l 
X-33 


E 

early eclipse HI-22 

Earth V-48 


XV-7 



Earth horizon constraint 
Earth radiation constraint 
Earth shield 
Earth’s magnetic field 
eclipse 

eclipse operation constraint 
ecliptic pole differences 
ecliptic poles 
edge detections 
edge overlap 
effective frequency 
effective wavelength 
electronics 
analog 

commandable gain and offset 
pole-zero amplifier 
digital 

frequency response 
low resolution spectrometer 
telescope 
electronics boxes 
equivalent cylindrical projection 
Errata 
Europa 

explanatory supplement 
extended emission 

consistency checking and data removal 
deleted detectors 
deleted fields in 
first release of 3rd HCON 
destriping 
detector weighting 
effective resolution 
filter 

final map generation 
gaps 

map images 

missing fields in 3rd sky coverage 
phasing 

projection into maps 
quality checking 

extended emission data processing 
extended emission formats 
extended emission photometry 
extended emission processing anomalies 


III-2, IV- 19 
m-2, m-4, m-5 

II- 7 

n-2, V-8 
n-2, III-22 

III- 8 
VI-9 
VI-6 

V-19, VTI-28 

V- 20 

see under ’band pass’ 
see under ’band pass’ 

11-14, II- 1 9ff 
11-21 
11-21 

11-23, VI- 1 
rv-9 

IX- 3 
n-i9ff 
II-7 

X- 31 
XII- Iff 

VI- 22ff 
M 


V-51 

V-49 

XI-3 

V-50 

V-48 

Vn-39 

see ’filter, Lanczos single smoothed’ 

V-52 

V-49 

1-3 

X-33 

V-49 

V-50 

V-48 

see under ’data processing’ and ’extended emission’ 
see under ’format’ 
see under ’photometry’ 

V-51 


fabry lens 
FAC 

failed detectors 
false alarms 
feedback resistor 
nonlinearity 


11-15 

see ’fine attitude calibration’ 

see ’detector, failed’ 

see under ’confirmation, point source’ 

II-19ff, VI-4 

II-20ff, IV- 15 


:i I r 


XV-8 



field of view 
detectors 
telescope 
filter 

electronic 

infrared 

transmission 

Kalman 

Lanczos single smoothed 
filter leaks 

filter transmission verification 
final processing steps 
fine attitude calibration 
fine attitude sensors 
first HCON coverage 
FITS format 
FITS headers 
five degree gap 
flags 
cirrus 
confuse 
confusion 
hex encoded 
meaning of point source 
small extended source 

flux 

average calculation 
conversion to flux density 
determination of relative flux 
discrepant 
high quality 
moderate quality 
uncertainties 
upper limit 
flux density 
conversion to 
flux overestimation 
flux reconstruction 
flux status word 
focal plane 
assembly 
diagram 
filter and lens 
geometry analysis 
format 

catalogs and atlases 
extended emission 
galactic plane maps 
low resolution all sky maps 
low resolution spectra 
map images 
prints 
tapes 

point source printed 
point source tape 


see ’detector, solid angle’ 

II-8, II- 1 1 

11-21 

n-is 

II-17ff 
V-7, V-9 
V-49 

II- 17 
IY-17 

see under ’data processing’ 

V-7 

V-7 

III- 14 
X-l, X-36 
X-51ff 
III-21 

X-9, X-l 5 

VII-37 

VII-36 

X-6 

VII-36 

X- 8, X-l 5 

V- 55, V-63 

VI- 27 
VI-5 

VIM 3, Vn-24 
V-55 

V-55, V-63 
V-55, V-63, VII- 18 

V- 55, V-63 

VI- 27, X-l 3 

XI- 4, XII- Iff 

Vl-lff 

V-33, V-55, V-63 

II-9ff 
II- 11 

11-15, see also ’survey array’ 
V-26 

X-lff 

X-30ff 

X-36 

X-37 

X-38ff 

X-32ff 

X-33ff 

X-36 

X-5, X-lOff, inside back cover 
X-2ff 


XV-9 



small extended sources 

X-29 

working survey data base 

X-12ff 

zodiacal observation history file 

X-37, X-61ff 

frequency response 

see ’electronics, frequency response 

variation with total flux 

VI-19 

FSTAT 

see ’flux status word’ 

G 



gain 


commandable 

11-21 

high 

11-23 

low 

11-21 

nominal 

11-23 

galactic anti-center 

V-6 

galactic center 

III-20, V-6, V-57 

galactic plane 

IV- 12, V-6, VI- 13 

galactic plane maps 

V-52 

galactic plane shadow 

see’shadow, galactic plane’ 

Gamma Dra 

VI-28 

gap, 5 degree 

III-2 1 

gaps 

m-i9 

gaps in extended emission maps 

see ’extended emission, gaps’ 

Gauss-Markov process 

V-7 

gaussian position errors 

see under ’position uncertainty' 

glints 

see ’Moon glints’ 

gnomonic projection 

X-30 

ground based observations 

VI-22, VI-28ff, VII- 11 

gyro 

II-2, V-3, V-7, V-8 

effect of Earth’s magnetic field 

V-8 

z-axis 

II-2 

ZA 

V-8, VII- 10 

ZB 

V-8, VII- 10 

gyro package 

II- 1 


H 


half orbit constraint 
half survey rate response 
HCON 
header files 
headers, FITS 
helium tank 
supercritical 
superfluid 
hex encoded flags 
high density bins 
high density processing 
high density regions 
location 

selection criteria 


m-17 

IV-9 

see ’hours confirmation’ I 

X-l 

X-51ff 


II- 3 
II- 3 


see ’flags, hex encoded’ 
X-47 

Vn-2, VII-25ff 


V-6, V-56 

V-57ff 

V-62 


XV-10 


1 li 



high gain 
high quality flux 
hole recovery strategy 
horizon sensor 
horns, polar 
hours confirmation 
decision 

photometric agreement 
point source 
position agreement 
statistical processing 
hours confirmed source 

relative photometric accuracy 
housekeeping data 
Hygiea 
hysteresis 


II- 23 

see under ’flux’ 

III- 17 

II- 2 

see ’polar horns’ 

III- 1, see also under ’confirmation’ 
see under ’confirmation, point source’ 
see under ’confirmation, point source’ 
see under ’confirmation, point source’ 
see under ’confirmation, point source’ 
see under ’confirmation, point source’ 
v-4, vn-i 

VII- 15 

11-23 

VI-22ff 

see ’detector, photon induced responsivity enhancement’ 


I 


in-flight modifications 
in-flight tests 
in-orbit checkout 
in-scan direction 
incompleteness, sources of 
index 

infrared cirrus 
infrared detectors 

insufficient spec, of HCON coverage 

integrating cavities 

intensity image 

internal reference sources 

interplanetary dust cloud 

IOC 

IRAS mission 
IRC +10216 


J 


JFETS 

Johnson noise 

junction field effect transistors 
Jupiter 

Jupiter avoidance strategy 
Jupiter constraint 


III-19ff 

iv-iff 

III- 17 
V-l 
VII-31 
XIV- Iff 
see ’cirrus’ 

see ’detector, infrared’ 

XI-3 

II- 11 

V- 50 

see ’reference sources, internal’ 

VI- 11 

see ’in orbit checkout’ 
see ’mission’ 

III- 20, IV- 15, VII-28ff 


see ’junction field effect transistors’ 
11-23 

II- 14, 11-19, VI- 1, VI-2 

IV- 19, V-48 

III- 13 
III-2 


K 


Kalman filter 
KAO 


XV- 11 


see ’filter, Kalman’ 

see ’Kuiper Airborne Observatory’ 



key value 11-25 

known processing anomalies see ’anomalies’ 

known source ID’s X-25 

known source file V-4 

known sources 

correlation with point sources V-26 

discrepancy analysis V-27 

prediction V-26 

Kuiper Airborne Observatory VI-28fF 


L 


Large Magellanic Cloud 

leaks 

life time 

limiting flux density 

limiting noise at low background 

line numbers 

linearity, infrared detectors 
low gain 
low quality flux 
low resolution all sky maps 
format 

Low Resolution Spectra 
blue energy distribution 
classification 
examples 
line spectra 
miscellaneous 

performance of classification scheme 
red energy distribution 
selection for catalog 
Low Resolution Spectrometer 
confusion 
cross scan response 
data base 
data processing 
despiking 
detectors 
electronics 
linearity checks 
memory effects 
multiplexer glitches 
optical properties 
performance and calibration 
photon induced responsivity enhancement 

radiation effects 
spectral resolution 
wavelength dependent responsivity 
wavelength scale 

low resolution spectrometer extractions 
low signal to noise detections 


V-36, V-57 
see ’filter leaks’ 
n-4 
1-1 
11-23 
X-35 

see ’detector, linearity’ 

11-21 

see ’flux, upper limits’ 

V-51, V-52 

X-37 

IX-lff 

IX- 13 

IX-13ff 

IX-16ff 

IX- 19 

IX- 19 

IX- 19 

IX- 14 

IX- 12 

IX-lff 

IX-8 

IX-6 

IX-9 

see under ’data processing’ 

IX- 10 

see under ’detector’ 
see under ’electronics’ 

IX-9 
IX-9 
IX-6 
IX- 1 

see under ’calibration’ 
see under ’detector, photon induced 
responsivity enhancement’ 

IX-6 

IX- 1, IX-4 
IX-6 
IX-6 
V-33 

see under ’source detection’ 


XV- 12 


LRS extractions 
lune 

lune constraint 


see ’low resolution spectrometer extractions’ 

m- 9 , m-io 

III- 17 


M 


M17 

magnetic field 
magnetic torque coils 
magnitudes, photometric 
main cryogen tank 
main shell 
map image format 

map image surface brightness formula 
map images 
artifacts 
color 

coordinate systems of prints 
deleted fields in first 
release of 3rd HCON 
final calibration factors 
label of prints 
location 
location on sky 

missing fields in 3rd sky coverage 
map projections 
maps, projection into 
Mars 

Martin optical black 
masks, detector 
memory, random access 
minisurvey 

minor processing anomalies 
missing maps from 3rd sky coverage 
mission 

chronology 
constraints 
life time 
requirements 
moderate quality flux 
Mon R2 
Moon 

Moon avoidance strategy 
Moon constraint 
Moon glints 
multilayer insulation 
multiplexer glitches 


III-20 

see ’Earth’s magnetic field’ 
II- Iff 

vi-20ff, vn-iiff 

II- 3 
II-3ff 

see under ’format’ 

X-33 


XI-2 

X-35 

X-35 


XI-3 

X-33 

X-35 

X-49 

X-34 

X- 33 

see ’projections’ 

see under ’extended source’ 

V-48 

II- 8 

n-n 

n-2 

m- 19 , V-53, VIII-2, VIII-8 

XI- 4 
X-33 
1-1 

III- 18 

m-2, ra-2ff 

n-4 

m-i 

see under ’flux’ 

VII- 30 

IV- 19, V-48 
III- 13 

III- 2 

IV- 19 
II-4 

see under ’low resolution spectrometer’ 


N 


NCOVR 


V-39, V-47 


XV- 13 



near neighbors 
NEFD 

neighbor flags 
neighbor tagging 
neighbors 

SES close neighbors 
SES flags 

bright sources removed as cross talk 

confused 

near 

point source flags 
very near 
weak 
Neptune 
NGC 2024 
NGC 2071 
NGC 6334 
NGC 6357 
NGC 6543 

noise 

coherent 

confusion 

detector 

Johnson 

limiting, low background 
process 

noise equivalent flux density (NEFD) 

noise estimator 

noise shadow 

noise spectra 

nominal gain 

non-boosted SAA crossings 
non-gaussian photometric errors 
non-gaussian position errors 
non-inertial sources 
non-seconds confirmed detection 
due to failed detector 
NSC 
NSCF 

nuclear particle hits 

nuclear pulse circumvention circuitry 


V-13 

see ’noise equivalent flux density’ 

VII-36 

V-53 

V-44 

VII-36 

VII-31 

V-63 

V-13 

VII-36 

V-63 

V- 62 

VI- 28 

VII- 30 
VII-30 

m-20 

III- 20 

IV- 3, IV- 17, IV-22, VI-5, 

VI-6, VII-11, IX-5, XI-3 


VI- 3 

VIII-4 
1-2 
11-23 
11-23 

see ’attitude reconstruction, process noise’ 

IV- 1 

see under ’source detection’ 

vm-n 

IV- 3 
11-23 

V- 51 
VTI-15ff 

see under ’position uncertainty’ 

VII- 35 
V-19 

V-19, V-25ff, Vn-32 

see ’non-seconds confirmed detection’ 

see ’non-seconds confirmed detection, due to failed detector’ 

see ’radiation hits’ 

11-21 


O 


offset levels, commandable 11-12, 11-21 

offset voltage VI- 1, VI-2, VI-5 

measurement VI-6ff 

time dependence of bias boost effect VT-2ff 

OMC 1 VII-30 

onboard computer II-2 

operation problems III- 19 

Ophiuchus V-57 


XV- 14 


TT1T 



optical black 
optics 

low resolution spectrometer 
survey array 
telescope 
orbit 
Orion 

Orion nebula 

out of band spectral leak verification 
out of band spectral leaks 
out of field radiation 
overlays for map products 

P 


particle radiation 
particles, dust 
phasing 

photometric accuracy 
photometric refinement 
hours confirmation 
seconds confirmation 
photometry 

extended emission 
point source 
processing anomalies 
small extended source 
photometry calibration strategy 
photon dose 

photon induced responsivity enhancement 
planet constraint 
point source catalog 
point source distribution 
point source format 
point source neighbor flags 
point source photometry 
pointing accuracy, absolute 
pointing errors 
cross scan 
in-scan 

pointing reconstruction 
polar horns 
pole to pole scans 
pole-zero amplifier 
porous plug 
position accuracy 
absolute 

of catalog sources 

position probability density function 
position reconstruction, point source 
position refinement 
hours confirmation 
seconds confirmation 


II-8 

IX- Iff 
II-15ff 

II- 7ff 

III- 2 

V-57, V1I-30 

III- 20 

IV- 17ff 

n-n 

see ’stray light’ 

X- 37 


see ’radiation’ 

III-l 

see under ’extended emission, phasing’ 
Vll-llff 

see under ’confirmation, point source’ 
see under ’confirmation, point source’ 

VI-6ff 

VI-6, VI-20, VII- 1 Iff, XI-4, XII- Iff 

XI-3 

VI-6 

III- 17 

IV- 1 3ff 

see under ’detector’ 

III-2 

see under ’catalogs’ 

see under ’catalogs, point source’ 

see under ’format’ 

VTI-36 

see under ’photometry’ 

II-2 

V- 3 

V- 3 

see ’attitude reconstruction’ 

11-21, III- 19, III- 19, VI- 12, IX-6 

VI- 9, VI- 11 
11-21 

n-4 

1-1 

VTI-5 

VTI-2ff 

see under ’confirmation, point source’ 
see under ’confirmation, point source’ 

see under ’confirmation, point source’ 
see under ’confirmation, point source’ 


XV- 15 



weeks confirmation 
position refinement 
position uncertainty 
anomalous processing 
gaussian position errors 
non-gaussian position errors 
pre-survey observations 
probability density function 
probability of variation 
problems, operations 
process noise 
projections, map 
pulse circumvention circuitry 


R 


RUMi 

radiation 

radiation blanking time 
radiation effects 

low resolution spectrometer 
radiation hit circumvention circuitry 
radiation hits 
radiators 

random access memory 
reaction wheels 
reconstructed data values 
recorder, tape 
recovery strategy 
reference source 
internal infrared 
internal optical 
rejected stim flashes 
stability 

rejected sources 
rejected stim flashes 
relative flux 
reliability 

formalism for 
estimation of parameters 
in the galactic plane 
moderate flux quality sources 
outside of galactic plane 
point source 
small extended sources 
resistor, feedback 
resolution 
angular 

extended emission, angular 
spectral 

response, spectral 


see under ’confirmation, point source’ 
see under ’confirmation, point source’ 
V-15, VTI-8 
XI- 3 

V-14, V-22ff, V-32, V-34, VII-8 
V-14, V-27, V-29, V-34 
III- 17 

see under ’confirmation, point source’ 
see ’variability, probability of 
III- 19 

see under ’attitude reconstruction’ 
X-30ff 
11-21 


IX-6 

VI-12 

V- 48 
IV- 11 
IX-6 

11-21, V-48 

11-21, III- 19, V-53, VII-32 

II-3 

II-2 

n-1, II-2 
n-25 

II- 1 

III- 17 

n-8, VI-5 

II- 8 

VI- 12 

IV- 22 

I- 3, VII-25ff 

see under ’reference source’ 

see ’flux, determination of relative flux’ 

III- 1 , VII- 1, Vlll-lff 
VIII-6 

VIII-7 

vm-io 

VIII- 10 
VIII- 8 

see under ’catalogs, point source’ 

VIII- 1 1 

II- 1 9ff 

1-1 

see under ’extended emission’ 
see under ’low resolution spectrometer’ 
11-18 


XV- 16 


: I I ! 


see under ’low resolution spectrometer’ 


LRS 

responsivity 

DC 

infrared detectors 
point source 

frequency dependence 


S 


SAA 

sample numbers 
sampling rates 
satellite 

description 

dimensions 

satellite operations plan 

saturation of detectors and electronics 

Saturn 

SC 

scan rate 

science data analysis system 
SDAS 

second HCON coverage 

second survey 

secondary mirror 

seconds confirmation 

seconds confirmation decision 

seconds confirmation processor 

seconds confirmed sources 

seconds confirmed, band merged source 

segment control word 

selection of point sources for catalog 

SES1 flag 

SES2 flag 

shadow 

galactic plane 
noise 

shadowing 
shift code 

signal to noise cutoff 
significand 

silicon diode detectors 
single HCONs 

sky brightness variation model 
sky coverage 
sky plate formats 

sky plate surface brightness formula 
sky plates 

small extended source 
clusters 
size 

small extended source catalog 
small extended source 


VI-5 
IV- 1 
VI-5 
IV-9 


see ’south Atlantic anomaly’ 

X-35 

11-23 

n-iff 

ii- 1 

II- 2, III-9 

III- 20 

IV- 12, V-48, VI- 13 

see ’seconds confirmed detection’ 

III- 13 

V- l 

see ’science data analysis system’ 
in- 15 
III- 13 

II- 7, II-9 

III- l 

see under ’confirmation, point source’ 
see under ’confirmation’ 

V-4 

V- 4 
X-12 

see under ’catalogs, point source’ 
VTI-36, VII-38, X-8, X-15 
VII-36, X-8, X-15 

vni-n 

see ’noise shadow’ 

see under ’source detection’ 

11-25 

see under ’catalogs, point source’ 

11-25 

see ’detector, visible’ 

Vn-25 

VI- 11 

see ’coverage’ 

see ’format, map image’ 

X-33 

see ’map images’ 

V-5, see also under ’data processing’ 
V-5 

see ’catalogs, small extended source’ 


XV- 17 



decision function 
small extended source format 
small extended source neighbor flag 
small extended source photometry 
solar panel 

solar radiation constraint 
solid angle of detectors 
SOP 

source detection 
detection rate 
filter 

point source 
small extended source 
noise estimator 
anomalies 
point source 
point source 
correlation coefficient 
low signal to noise detections 
noise estimator 
shadowing 
square wave filter 
template 
timing estimate 
small extended source 
square wave filter 
threshold 

correlation coefficient 
signal to noise 
source follower 
source shadowing 
south Atlantic anomaly 

map 

strategy 

south Atlantic anomaly constraint 
spacecraft 

control axes 

special calibration observation 
spectra 

spectral bandwidth 
spectral coverage 

low resolution spectrometer 
survey array 
spectral hole 
spectral leaks 

spectral passband verification 
spectral response 
square wave filter 
stability 

DC electronic 
detectors 

standard asteroid model 
star crossings 
star sensor 


see under’confirmation, small extended source 
see under ’format’ 

VII-36 

see under ’photometry’ 

II- 1 

III- 2, III-4 

see ’detector, solid angle’ 
see ’satellite operations plan’ 

V-9ff 
VII- 1 


V-3, V-9 
V-35 

XI- 1 

V-3 

V-3 

V-llff 

V-13 

V-10 

V-13 

V-9 

V-3, V-ll, V-13, V-14 

V-ll 

V-3 

V-35 


V-3 

V-3 


11-19 

see under ’source detection, point source’ 

II- 20, IV- 12, V-51, VI-2, 

VIII-2, IX-4, IX-6 

m-7, m-21 

III- 13 


III-3 

II-l 

n -2 

VT-8 

see ’low resolution spectra’ 
see ’band pass, spectral’ 


IX-1 

II-17ff 

V-57 

11-17 

rv-i6ff 

11-18, VI-20 

see under ’source detection’ 


VI-5 

IV-3 

VI-23 

11-23 

II-2 



station pass constraint 
statistics, point source 
stellar calibration 
stim flash 
stimulators 
stray light 

performance 
rejection monitoring 
stripes 
Sun 

Sun angle constraint 
Sun sensor 
spiking 

y-axis variability 
Sun shade 

surface brightness density, conversion to 
surface brightness formula 
surface brightness, conversion to 
survey array 
characteristics 
electrical 
optical 
construction 
diagram 
survey design 

recovery strategy 
second six months 
strategy 
survey rate 

survey strategy realization 
symmetry plane of zodiacal dust 


III- 5 

see under ’catalogs, point source’ 
see under ’calibration’ 

VI- 5 

see ’reference sources, internal’ 
n-8, V-48 

II- 8, II- 10 

IV- 20ff 

V- 50 

III- 2, IV- 19, VI- 19 
III-2, m-5 

II- 1, II-2, V-3, V-7 

V-7 

V-8 

n-3, ii-6 
V-52 
X-33 
V-49 

11-12 

11-20 

11-16 

11-14 

II- 11 

III- 17 
III- 13 
III-9ff 
III- 13 
III- 17 

VII- 3 8 


T 


tape recorder 
telescope 
electronics 
optical performance 
in-flight tests 
optics 
orientation 
overview 

system characteristics 
thermal control 
template 
tests, in-flight 
TFCAL 
TFPR 

thermal calibration sources 
thermal control 
thermal misalignment 
third HCON coverage 
TIA 


II- Iff 
II- 3 

II-7, II-8 

IV-19ff 

II-7 

II- 3 

n-3ff 

II-5 

II-7 

see under ’source detection, point source’ 

IV-lff 

see ’special calibration observation’ 
see ’total flux photometric reference’ 
see ’reference sources, internal’ 

II-6 

see under ’attitude reconstruction’ 

in- 16 

see ’trans-impedance amplifier’ 


XV- 19 



tic marks 
tie points 

time constant, offset voltage 
time ordered files 
total flux photometric reference 
annual variation 
flux determination 
model parameters 
total flux photometric region 
trans-impedance amplifier 
frequency response 

transformation equations, map projections 

transformation matrices 

TRIAD 

triple detection mode 


X-35 

V- 50 

VI- 3 

V- 49 

IV- 1 1, VI-6ff, VTI-39 

VI- 6, VI- 10, VI-11 
VI-7 

VI-10 

see ’total flux photometric reference’ 
n-19, VI-3 

n-20 

X-30ff 

see ’coordinate transformation matrices’ 


VI-23 

see under ’confirmation, point source’ 


U 


uncertainty, flux 

universal time seconds, zero point 

upper limit flux 

Uranus 

UTCS 


V 


van Allen belts 

vapor cooled shields 

variability, probability of 

variability analysis 

variable sources 

vent nozzels 

very near neighbors 

visible detectors 

visible star crossings 

visible star sensors 

visible wavelength channels 

visual stars 

Vtialoffl 


W 


warm up 
warnings 
wavelength scale 
weak neighbors 
weeks confirmation 
point source 
position refinement 


see ’flux, uncertainties’ 

V-3, X-62 
see under ’flux’ 

IV- 18, VI-28, VI-30 
see ’universal time seconds’ 


III-3, III- 19 
II-4 

V-56, Vn-22 

V-55 

VII-22ff 

II-6 

V-63 

II- 1 1 , 11-23, see also ’visible star sensors’ 

n-23 

V-7 

II- 11 

V-3 

see ’offset voltage’ 


III-22 

see ’caveats’ 

see under ’low resolution spectrometer’ 

V-62 

III-l 

see under ’confirmation, point source’ 
see under ’confirmation, point source’ 



small extended source 
threshold 

statistical processing 

weeks confirmation decision, point source 
weeks confirmation 
weight image 
working survey data base 
format 
WSDB 


Z 


zero clamp 

zero sum square wave filter 
zodiacal dust emission 
effects 
modeling 
variation 
zodiacal emission 
zodiacal observation history file 
format 
ZOHF 


see under ’confirmation, small extended source’ 
see under ’confirmation, small extended source’ 
see under ’confirmation, point source’ 
see under ’confirmation, point source’ 
see under ’confirmation’ 

V-50 

1-3, V-4, V-6, V-52ff, VII- 1 

see under ’format’ 

see ’working survey data base’ 


IX-3 

see under ’source detection’ 

VI- 11 
VTI-38 

vn-39 

VII- 38 

see ’zodiacal dust emission’ 

V-6, V-49, V-52, VI- 19, VII-39 
see under ’format’ 

see ’zodiacal observation history file’ 


XV-21 







NASA 

National Aeronautics and 
Space Administration 


1. Report No. 

NASA RP-1190 volume 1 


Report Documentation Page 


2. Government Accession No. 


4. Title and Subtitle 

Infrared Astronomical Satellite (IRAS) Catalogs and 
Atlases 

Volume 1 Explanatory Supplement 


3. Recipient's Catalog No. 

5. Report Date 

1988 

6. Performing Organization Code 


7. Author(s) 


C.A. Beichman, G. Neugebauer,* H.J. Habing,' 
P.E. Clegg,**** and T.J. Chester, editors 


8. Performing Organization Report No. 


10. Work Unit No. 


9. Performing Organization Name and Address 

Jet Propulsion Laboratory 
Pasadena, CA 91109 


12. Sponsoring Agency Name and Address 

Office of Space Science and Applications, Astrophysics 
Division 

National Aeronautics and Space Administration 
Washington, DC 20546 

15. Supplementary Notes 

-California Institute of Technology, Pasadena, Calif. 
**Sterrewacht , Leiden, The Netherlands 
***Queen Mary College, London, United Kingdom 


11. Contract or Grant No. 

13. Type of Report and Period Covered 
Reference Publication 

14. Sponsoring Agency Code 


16. Abstract 

The Infrared Astronomical Satellite (IRAS) was launched on January 26, 1983. 
During its 300— day mission, IRAS surveyed over 96% of the celestial sphere at 
four infrared wavelengths, centered approximately at 12, 25, 60, and 100 microns. 

Volume 1 describes the instrument, the mission, and data reduction. 

Volumes 2-6 present the observations of the approximately 245,000 individual 
point sources detected by IRAS; each volume gives sources within a specified 
range of declination. Volume 7 gives the observations of the approximately 
16,000 sources spatially resolved by IRAS and smaller than 8*. 


17. Key Words (Suggested by Author(s)) 


astronomy interstel 

infrared IRAS 

all-sky survey cosmology 
galaxies catalogs 

star formation stars 
19. Security Classif. (of this report) 
Unclassified 


interstellar matter 
IRAS 


18. Distribution Statement 

Unclassified - Unlimited 

Subject Category 90 


20. Security Classif. (of this page) 

21 . No. of pages 

Unclassified 

462 


A01 A18 


NASA FORM 1626 OCT 86 

For sale by the National Technical Information Service, Springfield, VA 22161 


11 



TABLE Suppl.VLC6 Color Correction Factors, K 1 


INTRINSIC 
POWER LAW 2 

RATIO OF FLUX DENSITIES 
BEFORE COLOR-CORRECTION 


CORRECTION FACTOR 


a 

/v( 12 pm) 

/v( 25 pm) 

/ v ( 25 pm) 

/v(6 0 pm) 

/»( 60 pm) 
/v<l00pm) 

K( 12 pm) 

K(25 pm) 

K(60 pm) 

K( 100 pnr 

-3.0 

0.113 

0.063 

0.21 

0.91 

0.89 

1.02 

1.02 

-2.5 

0.162 

0.102 

0.275 

0.92 

0.91 

1.00 

1.01 

-2.0 

0.232 

0.164 

0.355 

0.94 

0.93 

0.99 

LOO 

-1.5 

0.333 

0.262 

0.460 

0.97 

0.96 

0.99 

1.00 

-1.0 

0.480 

0.417 

0.600 

1.00 

1.00 

1.00 

LOO 

-0.5 

0.694 

0.662 

0.786 

1.04 

1.04 

1.02 

LOO 

0.0 

1.005 

1.045 

1.037 

1.10 

1.10 

1.05 

1.01 

0.5 

1.459 

1.642 

1.378 

1.17 

1.16 

1.09 

1.02 

1.0 

2.123 

2.567 

1.843 

1.25 

1.23 

1.15 

1.04 

1.5 

3.094 

3.992 

2.484 

1.35 

1.32 

1.23 

1.06 

2.0 

4.519 

6.170 

3.373 

1.47 

1.41 

1.32 

1.09 

2.5 

6.610 

9.480 

4.617 

1.61 

1.53 

1.44 

1.12 

3.0 

9.681 

14.475 

6.370 

1.78 

1.67 

1.59 

1.16 

INTRINSIC 

RATIO OF FLUX DENSITIES 
BEFORE COLOR-CORRECTION 


CORRECTION FACTOR 


TEMP (K) 

M 1 2 pm) 
U25 pm) 

/ v (25 pm) 
/ v ( 60 pm) 

/ v (60 pm) 
/ v (100pm) 

K(12 pm) 

K(25 pm) 

K(60 pm) 

K(100prr 

10000 

4.345 

6.050 

3.350 

1.45 

1.41 

1.32 

1.09 

5000 

4.172 

5.931 

3.327 

1.43 

1.40 

1.32 

1.09 

4000 

4.086 

5.872 

3.316 

1.42 

1.40 

1.31 

1.09 

3000 

3.944 

5.773 

3.297 

1.41 

1.39 

1.31 

1.09 

2000 

3.666 

5.578 

3.259 

1.38 

1.38 

1.31 

1.09 

1000 

2.891 

5.005 

3.145 

1.27 

1.34 

1.29 

1.08 

800 

2.545 

4.730 

3.088 

1.22 

1.32 

1.28 

1.08 

600 

2.036 

4.287 

2.995 

1.15 

1.29 

3.27 

1.08 

500 

1.692 

3.950 

2.920 

1.09 

1.26 

1.26 

1.08 

400 

1.272 

3.478 

2.810 

1.0! 

1.22 

3.24 

1.08 

300 

0.785 

2.780 

2.630 

0.92 

1.15 

1.21 

1.07 

290 

0.734 

2.693 

2.606 

0.91 

1.15 

1.21 

1.07 

280 

0.684 

2.602 

2.580 

0.90 

1.14 

1.20 

1.07 

270 

0.633 

2.506 

2.553 

0.89 

1.13 

1.20 

1 .07 

260 

0.583 

2.407 

2.523 

0.88 

1.12 

1.19 

1 .07 

250 

0.534 

2.304 

2.491 

0.87 

1.11 

1.19 

1.07 

240 

0.486 

2.196 

2.457 

0.86 

1,09 

1.18 

1.07 

230 

0.438 

2.084 

2.420 

0.85 

1.08 

1.18 

1.07 

220 

0.392 

1.967 

2.381 

0.85 

1.07 

1.17 

1.07 

210 

0.347 

1.845 

2.338 

0.84 

1.06 

1.16 

1.06 

200 

0.304 

1.719 

2.291 

0.83 

1.04 

1, 16 

1.06 

190 

0.263 

1.589 

2.240 

0.83 

1.02 

1.15 

1.06 

180 

0.224 

1.455 

2.184 

0.83 

1.01 

1.14 

1.06 

170 

0.188 

1.317 

2.124 

0.83 

0.99 

1.13 

1.06 

160 

0.154 

1.176 

2.057 

0.84 

0.97 

L 12 

1.06 

150 

0.124 

1.034 

1.983 

0.85 

0.95 

1.11 

1.05 

140 

0.097 

0.892 

1.901 

0.87 

0.93 

1.09 

1.05 

130 

0.073 

0.751 

1.810 

0.90 

0.91 

1.08 

1.05 

120 

0.053 

0.614 

1.709 

0.94 

0.89 

1.06 

1.04 

110 

0.036 

0.484 

1.595 

1.01 

0.86 

1.04 

1.04 

100 

0.023 

0.363 

1.468 

1.12 

0.84 

1.02 

1.04 

95 

0.018 

0.307 

1.400 

1.19 

0.83 

1.01 

1.03 

90 

0.014 

0.256 

1.326 

1.28 

0.83 

1.00 

1.03 

85 

0.010 

0.208 

1.249 

1.39 

0.82 

0.99 

1.03 

80 

0.007 

0.165 

1.168 

1.54 

0.81 

0.97 

1.02 

75 

0.005 

0.127 

1.082 

1.74 

0.81 

0.96 

1.02 

70 

0.003 

0.095 

0.993 

2.01 

0.81 

0.95 

1.01 

65 

0.002 

0.067 

0.898 

2.40 

0.82 

0.94 

1.01 

60 

0.001 

0.045 

0.801 

2.97 

0.83 

0.93 

LOO 

55 

— 

0.028 

0.700 

3.86 

0.86 

0.92 

LOO 

50 

— 

0.016 

0.597 

5.35 

0.90 

0.91 

0.99 

45 

— 

0.008 

0.493 

8.09 

0.97 

0,92 

0.98 

40 

— 

0.003 

0.391 

13.79 

1.08 

0.93 

0.98 


1 / v [actual] - / v Iquotedl/A' (See Suppl.VI.C.3); spectral response given in Table Suppl.II.G5. 

O <J 

2 /v- V" 


TITTI '