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OPTICAL BRIGHTNESS AND POLARIZATION 
OF QUASARS AND RELATED OBJECTS 



By 



BEN Q. McGIMSEY, Jr, 



A DISSERTATION PRESENTED TO THE GRADUATE COUNCIL OF 

THE UNIVERSITY OF FLORIDA 
IN PARTIAL FULFILLMENT OF THE • REQUIREMENTS FOR THE 
DEGREE OF DOCTOR OF PHILOSOPHY 



UNIVERSITY OF FLORIDA 
1974 



ACKNOWLEDGMENTS 

The author wishes to express his appreciation and 
thanks to Dr. A. G. Smith, the chairman of the author's 
Graduate Committee, for his guidance and support during the 
course of the research reported in this dissertation. The 
author also wishes to acknowledge the contributions made by 
Drs. A. E. S. Green, R. C. Isler, F. B. Wood, and S. T. 
Gottesman as members cf his Graduate Committee. 

Thanks are extended to A. G. Smith, R. J. Leacock, G. H. 
Folsom, R. L. Hackney, K. R. Hackney, R. L. Scott, and P. L. 
Edwards for their help in obtaining and reducing much of the 
data presented in this study. Special appreciation is 
expressed to R. L. Scott and P. L. Edwards for their assis- 
tance in the preparation of the photographs. The assistance 
of W. W. Richardson, H. W. Schrader, and E. E. Graves in the 
construction and maintenance of equipment used in the QSO 
program is gratefully acknowledged. Without their help it 
would have been difficult to conduct the research reported 
here. The efforts of Mrs. Elizabeth Godey in preparing the 
typed manuscript are appreciated. The data presented here 
were reduced with computer time donated by the Northeast 
Regional Data Center of the State University System of Florida 



11 



The author has been partially supported by an NDEA Title 
IV fellowship and by Graduate School and Arts and Sciences 
fellowships from the University of Florida. This support is 
gratefully acknowledged. 

The support of the author's family never failed during 
the years of his graduate career and is deeply appreciated. 
The author's greatest debt of gratitude is owed to his wife, 
Karen, for her understanding and encouragement and it is to 
her that this dissertation is dedicated. 



111 



TABLE OF CONTENTS 

Page 

ACKNOWLEDGMENTS ii 

LIST OF TABLES v 

LIST OF FIGURES vi 

ABSTRACT viii 

CHAPTER 

I INTRODUCTION . 1 

Quasars 1 

Optical Properties 3 

Models 10 

Polarized Component 20 

II OBSERVATIONAL TECHNIQUES AND EQUIPMENT 2 4 

The General Photographic Program 24 

The Polarimetry Program 30 

III RESULTS OF THE QSO MONITORING PROGRAM 39 

Discussion of Individual Objects 48 

IV POLARIMETRIC OBSERVATIONS OF BL LACERTAE , 

3C 120, AND 3C 273 99 

Error 99 

BL Lacertae 102 

3C 120 108 

3C 273 112 

V SUMMARY AND CONCLUSIONS 115 

BIBLIOGRAPHY 12 3 

BIOGRAPHICAL SKETCH 128 



IV 



LIST OF TABLES 

Table Page 

1 Calibrated Sources 42 

2 Sources with Tertiary Calibrations 46 

3 Sources with No Calibrations 50 

4 Sources with Known Polarization 103 

5 Polarimetric Observations 106 



LIST OF FIGURES 

Figure Page 

1 A spinar model of a QSO 14 

2 RHO 30-inch reflector showing the Cassegrain 
camera and Polacoat 105 UV filter 34 

3 Light curves of OB 338, 01 318, and OK 290 ... 56 

4 Comparison sequence of OQ 208 60 

5 Light curves of OQ 208, OX 074, and 

PKS 0202-17 62 

6 Comparison sequence of PKS 0202-17 66 

7 Light curves of PKS 0222-23, PKS 0336-01, 

and PKS 0458-02 68 

8 Comparison sequence of PKS 0222-23 70 

9 Comparison sequence of PKS 0458-02 74 

10 Light curves of PKS 0518+16, PKS 1252+11, 

and PKS 1347+21 77 

11 Comparison sequence of PKS 1252 + 11 79 

12 Comparison sequence of PKS 1347+21 82 

13 Light curves of PKS 1607+26, PKS 2209+08, 

and PKS 2345-16 85 

14 Comparison sequence of NRAO 140 89 

15 Light curves of NRAO 140 and NRAO 512 91 

16 Light curves of 3C 371 and 3C 446 95 

17 Comparison sequence of 3C 446 98 

18 Percent transmission versus wavelength 
for Kodak GG-13, Polaroid HN-32, and 

Polacoat 105 UV filters 101 



VI 



LIST OF FIGURES - Continued 



Figure Page 

19 The B magnitude, percent linear polariza- 
tion P, and position angle of polarization 

9 for BL Lacertae 105 

20 The B magnitude, percent linear polariza- 
tion P, and position angle of polarization 

6 for 3C 120 Ill 

21 The B magnitude, percent linear polariza- 
tion P, and position angle of polarization 

6 for 3C 273 114 



VII 



Abstract of Dissertation Presented to the Graduate Council 
of the University of Florida in Partial Fulfillment of the 
Requirements for the Degree of Doctor of Philosophy 

OPTICAL BRIGHTNESS AND POLARIZATION 
OF QUASARS AND RELATED OBJECTS 



By 

Ben Q. McGimsey, Jr. 
December, 19 74 

Chairman: Dr. Alex G. Smith 
Major Department: Physics 

Photographic monitoring of over 130 quasars and related 
objects has shown definite optical variability to be present 
in nineteen of these objects and suspected variability in 
forty-four others. Regular observations of the optically 
violent variables OB 338, OX 074, PKS 0518+16, PKS 2345-16, 
NRAO 512, and 3C 371 have revealed that these objects continue 
to undergo violent optical outbursts of a magnitude or more 
with the flare and subsequent return to minimum sometimes 
occurring within a few weeks. The earlier observations of 
PKS 2345-16 and NRAO 512 have been reported in a paper by 
G. H. Folscm, A. G. Smith, and R. L. Hackney, "Optical Flares 
in the Quasi-Stellar Radio Sources PKS 2345-16 and NRAO 512," 
in A strophysical Letters , 1_, 15 (1970); and the observations 
of OB 338 and OX 074 in a paper by G. H. Folsom, A. G. Smith, 
R. L. Hackney, K. R. Hackney, and R. J. Leacock, "Optical 
Changes in Eleven Ohio Radio Sources with Unusual Spectra," 
in The Astrophysical Journal (Letters) , 169 , L131 (1971). 

viii 



In addition, optical variability was observed for the first 
time in the objects PKS 0202-17, PKS 0222-23, PKS 1252+11, 
and NRAO 140. Both OQ 208 and PKS 0222-23 showed a total 
range of variability of a magnitude or more. The optically 
violent quasar 3C 446 underwent a 2.5 flare in 1974 with an 
extremely rapid decay time. 

Two types of characteristic variability were observed. 
The first was a violent optical outburst with rapid rise and 
decay times, observed in the objects described above. The 
second type of characteristic variability was a slowly varying 
component which appeared as a gradual rise and fall in optical 
brightness. Both types of variability were observed in 
several sources, while other sources exhibited only one type. 

There is observational indication that major flares in 
optically violent objects tend to be followed by a succession 
of minor flares of decreasing amplitude superimposed upon the 
decline in brightness following the major flare. This effect 
was observed in NRAO 512, OX 074, PKS 2345-16, PKS 2209+08, 
and possibly OQ 208 and PKS 0222-23. 

Optical linear polarimetric observations were conducted 
photographically on BL Lacertae, 3C 120, and 3C 273. These 
observations indicate that the peculiar object BL Lac has a 
very large, rapidly varying linearly polarized component. 
The polarized component of the Seyfert galaxy 3C 120 is 
smaller and more slowly varying than that of BL Lac. The 
quasar 3C 273 seems to lack a linearly polarized component. 
There seems to be a positive correlation between the optical 

ix 



activity of an object and the magnitude and variability of 
its linearly polarized component. These results were par- 
tially communicated in a paper by B. Q. McGimsey, A. G. 
Smith, R. L. Scott, R. J. Leacock, and P. L. Edwards, 
"Optical Linear Polarization in the Extragalactic Sources 
BL Lacertae and 3C 120," presented at the 38th Annual Meeting 
of the Florida Academy of Sciences , 21-2 3 March 19 74 in 
Orlando, Florida. 



CHAPTER I 
INTRODUCTION 

Quasars 

The attempt to determine the true nature of quasars and 
related objects has been one of the most important efforts of 
the past decade in astronomy. In 1960 Sandage and Matthews 
photographed the radio source 3C 48 v;ith the 200-inch Palomar 
telescope. The optical counterpart to this small angular 
diameter radio source was found to be a 16th magnitude 
"stellar" obiect with a wisp of nebulosity. Spectra taken by 
a number of astronomers indicated an abnormally blue object 
with broad emission lines which could not be identified. This 
problem was solved by Schmidt (1963) when emission lines in 
5C 273, a similar object, were identified as the Balmer lines 
of hydrogen and a line of Mg II shifted toward longer wave- 
lengths of the spectrum with a redshift z = AX/X = 0.15 8. 
This led to the determination by Greenstein and Matthews 
(1963) that the emission lines of 3C 48 exhibited a redshift 
z = 0.367. Later discoveries have shown that these values of 
z are low for quasars. The quasi-stellar object with the 
largest redshift known at present is OQ 172 with z = 3.53 
reported by Wampler e_t al_. (1973a). 



Quasi-stellar sources can be briefly described as being 
star-like, often variable in both radio and optical emission, 
with large ultraviolet and infrared fluxes , and exhibiting 
broad emission lines with large red shifts. In addition, 
several objects have a large linearly polarized component. 
The field has grown to include "quasi-stellar objects," which 
have the same optical properties but are very weak radio 
emitters; and Lacertids, which have radio and optical spectra 
similar to quasi-stellar sources, are very bright and highly 
variable, but lack emission lines, thus not allowing a red- 
shift to be determined. 

The assumption of a cosmological interpretation of the 
observed redshifts, i.e., that the redshifts are due to the 
relativistic expansion of the universe, leads to difficulties 
in explaining the physical mechanism of radiation. These 
objects appear to be at least as luminous as galaxies but 
the time scale of the optical variability of at least some 
objects indicates optical sources of light day to light 
months in size. Explanations of quasars assuming the red- 
shift to be due to a ncn- cosmological mechanism are not satis- 
factory at the time of this writing. 

Assuming the cosmological interpretation, quasi-stellar 
objects appear to be the most distant objects in the known 
universe, and thus the youngest observed. The study of these 
objects will hopefully lead to information about their origin 
and radiation mechanisms and perhaps to knowledge of the uni- 
verse at earlier epochs, and its evolution. To contribute to 



this study is the purpose of the University of Florida pro- 
gram. 

The present program was established in 1968 by A. G. 
Smith and G. H. Folsom, v:ith the Rosemary Hill 30-inch reflec- 
tor and has grown into a photographic monitoring program of 
about 150 quasars and related objects. This program has been 
continued by Smith and others and has been expanded to in- 
clude multicolor photometry. Recently linear polarization 
observations of several objects have been initiated by the 
author. 

The terminology used in describing quasars has remained 
somewhat unsettled. In general both radio-strong and radio- 
weak quasars will be referred to as quasi-stellar objects or 
QSO's by the author. Similar objects, such as Lacertids , 
Seyfert galaxies and N galaxies, will be referred to individ- 
ually by their proper names, and collectively as "related 
objects . " 

Optical Properties 

Stellar Imag e 

Quasi-stellar objects, as the name implies, are star-like 
optical objects identified with compact radio sources. Some 
have associated wisps or jets of nebulosity, as does the QSO 
3C 273B, which is associated with an optical jet 3C 273A. 
However, this is not the general case. 



Quasi-stellar objects are usually identified by checking 
a radio source error box for optical objects. If no other 
obvious source, such as a galaxy, appears within the error 
rectangle, the identification of the radio source is tenta- 
tively made with a blue stellar object within the error box, 
pending further confirmation. This confirmation is usually a 
spectrogram showing a QSO-like appearance. A highly red- 
shifted line spectrum is considered definite proof. 

Variability 

Many quasi-stellar objects show optical variability. 
About thirty-four of two hundred and two objects in a catalog 
by De Veny et al_. (1971) are classified as optical variables. 
Several types of variability have been recorded. Some show 
small-amplitude (less than 0.5 magnitude) irregular variations 
over a period of weeks or months (type V variations) . One 
such source is the QSO 4C 05.34, which has been reported by 
Hackney (1973) to show a statistically significant one-tenth 
magnitude variation on the time scale of a year. Others, 
classified as optically violent variables (OW) , show varia- 
tions of a magnitude or more in days or weeks. One example 
of such an object is PKS 2345-16 (Folsom et al_. 1970;1971). 
This source has shown a total change of about 1.5 in a period 
of three months in late 1969 and early 1970. More recently, 
the Rosemary Hill program shows that it has undergone type V 
variations in the 1972 observing season, suggesting that a 
QSO can show more than one type of variability. 



As this is written, there has been only one QSO that 
shows possible periodicity in its optical radiation. Kinman 
et al_. (1968) have seen an eighty-day periodicity in optical 
bursts from 3C 345 with one chance in several thousand that 
this periodicity is fortuitous. They also see a pattern in 
the phases of the eighty-day outbursts indicating a 321.5-day 
period. There is a chance of one in eighty that this second 
period is fortuitous. Pour out of five expected outbursts in 
the 321.5 day period have been seen from 1965 to 1969 (Morri- 
son 1969). However, data taken from 1970 to 1972 at Rosemary 
Hill tend to disprove this periodicity (Hackney 1973). 

It is interesting to note that there is a correlation 
between violent optical variability and the radio spectral 
index (Folsom et_ al_. 1971; Folsom 1970; Andrew et_ al_. 1972). 
The violent optical sources usually exhibit radio spectra 
with a small spectral index (flat) or spectra which show an 
increase in emitted flux toward short (centimeter or milli- 
meter) wavelengths. 

The diameter of a variable optical source at cosmologi- 
cal distances is given by Terrell (1967) for non-relativis- 
tically expanding sources. The observable change in luminosity 
of a sinusoidally varying source is approximately 

at 4 cT 



where T is the intrinsic period, D the diameter, L the 
o r 

observed luminosity and c the speed of light. This merely 



gives the familiar result that a change is observable only 
if the travel time of light across the object, D/c, is less 
than or on the order of the period of variation. The observed 
period T is equal to T times 1+z where z is the redshift. 
Solving for D and expressing AL/L as a change in magnitude 
implies 

Using observed quantities from some typical QSO's leads to 
diameters in the order of a light year for slowly varying 
objects, light weeks for PKS 2 345-16, and a light day for the 
optically violent variable 3C 446 (Burbidge 1967). For a 

relativistically expanding source, the right side of equation 

2 2 1/2 
(1.1) must be multiplied by the factor y = (1 - v /c ) in 

order to give correct results. 

Ultraviolet Excess 

A striking property of most quasi-stellar objects is a 
large excess in the ultraviolet component of the electro- 
magnetic radiation as compared to galactic stars. The U-B 
color index ranges from -1.2 to -0.4. In a plot of U-B versus 
B-V color indices, QSO's are above the region of main sequence 
stars, in the region of white dwarfs. This ultraviolet excess 
is so striking that it is considered a strong factor in favor 
of the identification of a stellar-like source as the optical 
counterpart of a compact radio source. 

This ultraviolet excess has let to discovery of another 
subclass of QSO's. During an optical search for QSO's, 



objects were found with ultraviolet excess but with no radio 
counterpart (Burbidge 1967). When spectra were taken, some 
of these objects proved to be galactic stars, while many were 
proved to be QSO's that are either radio quiet or weak radio 
emitters . 

Line Spectra and Redshifts 



Possibly the most striking feature of quasi-stellar 
objects is the presence of large redshifts in their line spec- 
tra. The spectra may consist of both emission and absorption 

o 

lines. The emission lines are broad O100A), of the type 
which may be expected to arise in hot gaseous nebulae, in 
radio galaxies, and in the nuclei of Seyfert galaxies. The 
stronger lines seen are the Balmer lines for objects with 
small redshifts, Lyman «, C IV A 1549, C III X 1909 and Mg II 
X 2798. Absorption lines are seen in many, but not all, 
QSO's. They are narrow and are usually resonance lines. The 
only line feature observed in the radio region has been the 
21 cm hydrogen line in absorption recently found by Brown 
and Roberts (1973) in 3C 286, believed to be due to a fore- 
ground object. The redshifts for 3C 286 are z = 0.849 and 
b J em 

z , „ = 0.692. 
abs 

The size of the emission line redshifts z = AX/X range 
from 0.06 for B 234 to the recently established 3.53 for OQ 
172. Some objects, e.g., PKS 0237-23, exhibit several differ- 
ent values of z , , all of which may be different from z^ . 

abs * em 

Usually the absorption redshifts are smaller than the 



emission redshift. There is some evidence that absorption 
lines may be more prevalent in high redshift objects (Bur- 
bidge 1967). 

Continuous Energy Distribution 

Knowledge of the shape of the continuous energy distri- 
bution of quasi-stellar objects is important to construction 
of models and mechanisms for radiation emission. Stein (1967) 
has plotted a spectrum for 3C 273 covering a large part of the 
spectrum. This is the only object for which there are enough 
data to make such a plot. The steep rise into the infrared 
suggests a non-thermal process, either the synchrotron or 
inverse Compton mechanism. 

There is no "average" QSO spectrum, because of individual 
differences among the sources (Oke 1966). Even so, one can 
fit a simple equation to the form of the optical spectrum: 

F(v) - v" n (1.3) 

F(v) - e ~ v/v ° t (1.4) 

where F(v) is the observed energy distribution and v is fre- 
quency. Since there is no typical QSO spectrum, there is 
some disagreement about the form of the distribution. Schmidt 
(1968) found a value of n = 0.7 using equation (1.3) after 
correcting UBV colors of several sources for interstellar 
reddening. Lari and Setti (1967) prefer equation (1.4) with 

v = 1.2xl0 15 Hz. 
o 



X-rays have been observed from the direction of 3C 2 73 
(Friedman and Byram 1967). If 3C 2 73 is the source, the 
authors claim that the intensity in the X-ray region is about 
the same as the optical intensity. 

In the radio region QSO's have spectra of the form of 
equation (1.3) with a median value of n near 0.7. Some QSO's 
have spectra that show curvature, getting flatter towards 
longer wavelengths. In some cases a maximum appears and the 
flux decreases toward longer wavelengths (Burbidge 1967). 
Williams (1963) proposed that the shape of this part of the 
spectrum indicates a synchrotron origin (Slish 1963). 

Related Obje cts 

There are several classes of objects showing many simi- 
larities to QSO's. These objects may help to determine the 
true nature of QSO's. The three main types of these related 
objects are Seyfert galaxies, N galaxies, and Lacertids. 

Seyfert galaxies are spiral galaxies with bright, 
stellar-appearing nuclei. They exhibit strong emission lines 
which are greatly broadened. Seyferts are strong emitters in 
the infrared but show enough ultraviolet excess to appear 
bluish. The observed redshifts in the emission lines are 
assumed to be due to cosmological expansion of the universe. 

The nuclei of N galaxies are even bluer and more star- 
like than those of Seyferts. The envelope is smaller and 
more amorphous. N galaxies also exhibit both the strong 
emission lines and large infrared flux seen in Seyfert galax- 
ies and strong low frequency radio emission. 



10 

"Lacertid" is the name given to a class of bright objects 
of which BL Lacertae is the prototype. The spectrum of these 
objects is non-thermal with a peak in the infrared, an ultra- 
violet excess, and no emission or absorption features. Lacer- 
tids also exhibit rapid, large amplitude variations in the 
optical, infrared, and radio flux. There is some indication 
that there may be appreciable optical variability on a time 
scale of less than one hour (Scott et_ a_l. 1973a). 

Recently reported observations by Oke and Gunn (19 74) 
indicate that Lacertids may be at cosmological distances. 
The authors, using the 200-inch telescope at Mt. Palomar, 
observed a redshift of z = 0.0 7 in the nebulosity surrounding 
BL Lac when the center was occulted by using an annular aper- 
ture. The spectrum shows lines seen in ordinary giant ellip- 
tical galaxies. This leads to the assumption that the very 
strong spectral continuum masks redshifted emission features 
in Lacertids. 

Models 

There is considerable disagreement among astronomers 
concerning the nature of quasi-stellar objects. Many, for 
example Morrison (1969) , believe these objects to be extra- 
galactic at the cosmological distance "D" indicated by the 
observed redshift z. 

D = J- (1.5) 

H o 



11 

2 2 1 / ° 
where v is determined by (1+v/c)/ (1-v /c ) ' = 1+z , where 

H is Hubble's constant, v is the velocity of recession and 

c the speed of light. Others, such as Kellerman (1972), 

suggest that QSO's are extragalactic but closer than distance 

D. Thus the redshift would be of Doppler or gravitational 

origin. Receiving more attention recently has been the idea 

that the redshifts are partly of cosmological origin and 

partly of some other origin. All of these hypotheses and - 

several models to account for the observed phenomena will be 

discussed below. 

C osmological Hypothesis 

The cosmological hypothesis assumes that the redshifts 
of QSO's are due to the expansion of the universe. The main 
objection to this assumption is that the observed variations 
in optical flux indicate a small diameter (see equation 1.1) 
for QSO's while the assumption of cosmological distances 

A 1 

indicates large amounts of radiation being emitted, 10 
erg/sec in 3C 273, for example (Burbidge 1967). On the 
other hand, QSO's are similar to Seyfert and N galaxy nuclei, 
which are assumed to be at cosmological distances. The 
results of Oke and Gunn (19 74) showing the similarity of the 
spectrum of the nebulosity surrounding BL Lac to spectra of 
giant elliptical galaxies support this hypothesis. In addi- 
tion, Bahcall and Hills (1973) have found that for the optic- 
ally most luminous quasars with redshifts from 0.2 to more 
than 2, the slope of the magnitude-redshift relation is 



12 

consistent with the value of five expected from the expansion 
of the universe. The fact that there are no observed blue- 
shifted QSO's adds credence to this theory. 

Spinar Model 

Morrison (1969; 1973) has proposed that QSO's are super- 
ficially similar to pulsars, though on a much larger scale. 
Quasi-stellar objects are assumed to be compact, spinning 
masses with a corotating magnetic field surrounded by a pool 
of relativistic material (see Figure 1). The ultimate energy 
source is gravitational. As the object contracts gravita- 
tional energy is converted to rotational energy and, through 
the magnetic stirring and accelerating processes in the pool, 
into particle and electromagnetic energy. Mass is emitted 
from the surface in a relativistic beam which corotates with 
the spinar. Optically, one observes the varying emissions 
from the pool except for the QSO 3C 345. In this case, 
Morrison theorizes , the observer is in the plane of the 
emitted beam and directly observes the 321.5-day rotation 
period. The primary emission is infrared, by the synchrotron 
process. Optical and x-radiation arise from the inverse 
Compton process involving the infrared photons and the rela- 
tivistic particles. The radio bursts are due to expanding 
lumped mass ejections beyond the critical surface. Optical 
emission and absorption lines arise in hot and cool gas 
regions beyond the critical surface. 



Figure 1. A spinar model of a QSO (Morrison 1973). 

Reproduced by permission of Physics Today . 
The numbered regions are (0) spin axis, 
(1) surface of spinar, (2) synchrotron 
emission of infrared, (3) Compton-recoil 
emission of optical and x-ray continua, 
(4) critical surface, (5) radio burst 
clouds, (6) emission-line optical source, 

(7) absorption line optical source, and 

(8) radio emission from weak-field synchro- 
tron plasma. 



14 



* 



^i-',i/,^' l , ivT/"*5« , .':*-V«'5 , .*.V..V.'. 
jMj V * # ; :,-:,»: .-■-;.■:::.;.;■;••■:■.: 

fit* . , \ - V ■ ' . ...»"* . 



v 




\' '■.'- V- \v 



•*% $&$ 



•*c 



Sir 1 -'.-: 

i — 

6 



Component A 



_L 



10 155 <10 16 >10 16 =10 165 10"-10 ,s 10 20 10 23 Radius R (cm) io :t 

;rl0 5 10 4 >10 3 1-10- 3 Magnetic field B (gauss) 10-' 



15 

Magnetic Rotator 

Piddington (1970) assumes that a weak intergalactic 
magnetic field perpendicular to the rotation axis of a proto- 
galaxy is wound several times around the protogalaxy. The B 
field is strengthened and finally erupts in tongues along the 
rotation axis due to the Rayleigh-Taylor instability. The 

Q 

final stage is a star-depleted gas cloud of mass 10 Mo, 

16 9 

radius 10 cm and surface speed of 10 cm/sec with a frozen- 

in field of 10 gauss. Particles are accelerated by pulsar 

mechanisms and account for the optical and infrared continuum 

by the synchrotron process. The shrinking gas cloud heats up 

to provide emission lines. The connection between the tongues 

and the internal field system provides for escape of rela- 

tivistic electrons into the tongues , where they are trapped. 

These electrons radiate by the synchrotron mechanism, leading 

to the familiar double-lobed radio source. Radio-quiet QSO's 

never developed the external field system. The energy source 

is again electrodynamic conversion of gravitational energy. 

Supernova Theory 

Colgate (1969) has proposed that the energy source of 
quasi-stellar objects is a chain reaction of supernovae. 
Stars in a dense cluster collide and coalesce when the rela- 
tive kinetic energy is insufficient for disruption. A few 
stars will rapidly grow to sizes of about 50 Mo, evolve to 
the supernova stage in about 10 years and erupt at a rate of 
about five supernovae per year to account for the observed 



16 



energy radiated. Permeating the cluster is a gas of density 
6x10 particles/cm , the result of the stellar collisions and 
material from previous supernovae. The optical continuum 
arises from the collisional heating of the expanding supernova 
envelope by the ambient gas to give a peak output of 10 
ergs/sec. The optical emission lines arise from the excita- 
tion of the entire gas cloud. Thus, the continuum will fluc- 
tuate depending on the time of the supernova eruptions while 
the line emission will be fairly constant. The emission lines 
are broadened by self- absorption in the centers of the lines, 
causing them to be five to ten times as wide as the absorption 
lines, as is observed. The infrared and millimeter radiation 
is due to the scattering of photons within a plasma (the 
ionized gas cloud) excited into oscillation by the two-stream 
instability. 

Local Hypotheses 

Another suggested model of QSO's is that they are expelled 
from galaxies and that the redshift is due to the velocity of 
recession. For local objects moving at relativistic veloci- 
ties, the number of blueshifted objects should be at least 
4 times the number of redshifted objects for z = 1.5 (Burbidge 
1967; Chiu et al_. 1973). Since no blueshifted QSO's have been 
identified, the expectation is that QSO's must have been 
ejected from some nearby center. Terrell (1967) derived a 
mass of 10 Mo for the average QSO from the upper limit of 
the proper motion of 3C 273. Assuming that the approximately 



17 

10 QSO's (Schmidt 1969) were expelled from the center of the 
Milky Way and radiate with 10% efficiency of conversion of 
mass to energy, all the mass within 100 parsecs of the center 
of the galaxy would be consumed, an improbably high number. 
The same energy problem holds for any nearby source. 

A second local model assumes that the spectral lines of 
QSO's are formed next to a compact massive object, thus caus- 
ing a redshift due to the gravitational potential of the 
object. The gradient of the potential over the emitting 
region must be small since the line widths are much smaller 
than the wavelength shift. The main objection to this model 
is that, according to theory (Burbidge 1967), z must be less 
than or equal to two, which conflicts with observed redshifts, 
A model consisting of a cluster of collapsed objects can show 
a redshift greater than two (Hoyle and Fowler 1967) , but it 
shows much too short a lifetime to be acceptable (Zapolsky 
1968). 

Other Hypotheses 

There have been several observational indications that 
redshifts cannot be entirely cosmological. Although no com- 
plete model has been advanced, these objections to the cosmo- 
logical hypothesis should be noted since they probably will 
have a direct effect on future theories of QSO's. 

Burbidge and Burbidge (196 7) have pointed out that there 
is a frequent occurrence of the redshift z = 1.95 in both 
emission and absorption lines of QSO's. This suggests that 



18 

a mechanism other than cosmological expansion is causing at 
least portions of the redshifts. While not to be dismissed, 
this evidence for non-cosmological z has not been completely 
convincing (Schmidt 1969). 

Stein e_t al. (1971) assumed that the spectrum of the 
Lacertid BL Lac can be attributed to synchrotron radiation 
and that synchrotron self-absorption occurs. From these 
assumptions they have derived an upper limit to the distance 
of BL Lac as 300 kpc. This conflicts with the distance of 
350 Mpc. measured by Oke and Gunn (1974). 

Arp (19 71) has claimed to see a luminous bridge between 
a peculiar galaxy and a QSO but this has been disputed by 
others (Lynds and Millikan 1972; Ford and Epps 1972). Further 
evidence for galaxy-QSO association has been advanced by 
other observers (Burbidge ejt al. 1971; 1972). Using the 3C 
and 3CR catalogs of radio sources, the Burbidges and their 
coworkers found that five of forty-seven QSO's lay within a 
few minutes of arc of bright galaxies of markedly different 
redshift. They calculated the probability of this effect 
being a chance occurrence as 10 . Further, the galaxy-QSO 
separations are inversely proportional to the galaxy red- 
shifts. No statistically significant close association 
between galaxies and QSO's in the Parkes 2.7 GHz survey in 
the ±4° declination zone was found by the authors. However, 
they claim this may be due to selection effects. The authors 
conclude that galaxies give rise to QSO's. 



19 

Hazard et al_. (19 73) have found four QSO's within one 
arc minute of one or more galaxies. The search involved look- 
ing for blue stellar objects near galaxies two or three mag- 
nitudes fainter. The four objects noted were later identi- 
fied as QSO's. The authors note that the chances are one in 
one hundred of finding one such QSO-galaxy grouping, assuming 
cosmological distances. 

There have been two discoveries of close pairs of QSO's 
with very different redshifts. Stockton (19 72) found that 
Ton 155 and Ton 156 are approximately 35 arc seconds apart. 
Wampler et_ al_. (1973b) have discovered that two QSO's, PKS 
1548+115a and 1548+115b, are separated by five arc seconds 
with redshifts of 0.44 and 1.90, respectively. (PKS 
1548+115b had not previously been identified as a QSO.) The 
authors also note the wavelength ratio (1+z) /(1 + z), = 2.02, 

3- D 

a curious coincidence. The probability that this is a chance 
occurrence of two QSO's at cosmological distances is esti- 
mated as 4 x 10" by the authors. However, Bahcall and 
Woltjer (1974) found a much greater probability (0.50) for 
random associations and concluded that these two pairs could 
be merely apparent associations of widely separated objects. 

Chiu et_ al_. (1973) have pointed out in defense of the 
cosmological theory that astronomers may be observing two 
types of objects: (1) the majority of QSO's at cosmological 
distances and (2) a "dwarf" branch originating from explosions 
in nearby galaxies. The authors propose that the Lacertids 
are blueshifted members of the dwarf branch, with the shifted 



20 

infrared continuum swamping the lines. In conclusion, there 
is still much disagreement on the meaning of the redshift. 

Polarized Component 

Recently, there has been great interest in the polarized 
component, of the radiation of QSO ■ s and related objects. As 
this is written, there have been no observations of signifi- 
cant optical circular polarization (Landstreet and Angel 1972; 
Schmidt 1969). However, Conway et al . (1971) have observed 
small (less than 0.1%) circular polarization at 21 cm and 49 
cm in several QSO's and BL Lac (Biraud 1969). 

Significant variable linear polarization has been 
observed in the continuum emission in many QSO's and related 
objects (Burbidge 1967; Schmidt 1969; Visvanathan 1973a). 
There has been no linear polarization observed in emission 
lines of QSO's. The polarized component amounts to 10% or 
more of the optical continuum flux in some objects (Visvana- 
than 1969; 1973a). There is some evidence that optical vari- 
ability in QSO's is accompanied by a higher percentage of 
linear polarization than in non-variable sources (Kinman 
et al . 1969). Lacertids, which tend to be violently vari- 
able, show large amounts of polarization, thus supporting 
this evidence (Visvanathan 1973a; 1973b; Kinman et_ al_. 1967). 
In general, increases in the total flux from an object tend 
to be accompanied by an increase in polarized flux (Visvana- 
than 1973a; Williams et al. 1972). Kinman et al. (1974) 



21 

report interesting behavior of linear polarization in both 
the radio and optical wavelengths of 3C 345 and OJ 287. The 
position angle of polarization in both objects tends to show 
sudden departures from a smooth trend, followed by a return 
to the original trend. 

Possible Mechanisms 

The two mechanisms discussed as possible sources for the 
observed polarization from QSO's and related objects are the 
synchrotron mechanism and the inverse Comptcn effect. In a 
simplified view of the synchrotron mechanism, relativistic 
charged particles move with a component of velocity perpendicu- 
lar to a magnetic field. The particles describe a circle or 
helix about the magnetic lines of force. Radiation is emitted 
in a narrow cone in the direction of the particle velocity. 
This radiation is almost completely polarized with the elec- 
tric vector parallel to the plane of the instantaneous circu- 
lar motion. The assumed reason that only a small percentage 
of linear polarization is measured is that the observed radi- 
ation is the resultant emission from particles moving in an 
inhomogeneous field, or emissions from several discrete 
sources with different alignments of magnetic field. 

The inverse Compton process involves the collision of a 
high energy particle, usually an electron, and a low energy 
photon, which results in a high energy recoil photon and a 
decrease in electron energy. If the low energy photons are 
polarized, the inverse Compton effect would result in high 



22 

energy photons with the polarization preserved. However, 
this preservation of polarization during inverse Compton 
scattering has been disputed by Bonometto and Saggion (1973). 
The inverse Compton process becomes more important as the 
electromagnetic energy density increases and it will dominate 
the synchrotron mechanism as the mechanism for particle 
energy losses if the electromagnetic energy density is greater 
than twice the magnetic field energy density (B /8ir). 

There are as yet no suggestions for any other possible 
mechanisms . 

The facts that a power law representation of the optical 
continuum is valid, and that percentage polarization is inde- 
pendent of wavelength, are taken as indications that the 
optical continua of OJ 287 and BL Lac are of synchrotron 
origin (Visvanathan 1973a; 1973b). The shape of the continuum 
in the infrared has been taken to indicate that this portion 
of the spectrum is also due to synchrotron emission (Ozernoy 
and Sazonov 1971) . 

It is difficult to compare the observed polarization 
properties mentioned in the previous section with QSO models, 
as the models do not give a detailed picture of emission 
mechanisms. The polarization observations are probably con- 
sistent with the spinar or the magnetic rotator models, 
although the spinar model does predict that the optical com- 
ponent is of inverse Compton origin in contradiction to 
Visvanathan' s results. It is difficult for the writer to 



23 

believe that the linear polarization reported could originate 
in the supernova model. Kinman e_t al. (19 74) have reported 
that their observations of OJ 287 are inconsistent with an 
expanding source model. 



CHAPTER II 
OBSERVATIONAL TECHNIQUES AND EQUIPMENT 

The General Photographic Program 

The Rosemary Hill observing program consists of regular 
observations of QSO's and related objects to search for opti- 
cal variability and to monitor known variables. In addition, 
possible changes in spectral slope are monitored by observa- 
tions in the Johnson UBV magnitude system (Hackney 1973). 

The method of photographic photometry was chosen for the 
Rosemary Hill monitoring program for several reasons. The 
first and most obvious is that with a 30-inch reflector photo- 
electric observations of many faint (on the order of 15th 
magnitude or fainter) QSO's would be difficult. Objects 
fainter than this magnitude make up a large percentage of the 
present observing program. While observation of the few 14th 
magnitude or brighter QSO's and related objects would be 
possible photoelectrically with this telescope, these obser- 
vations would require integration times of several minutes 
each for the object, comparison star, and sky background. 
For an object of about 14th magnitude, a typical unfiltered 
photographic exposure time for object and comparison sequence 
is about four minutes. Thus, time resolution is improved by 
the photographic method, enabling the observer either to 

24 



25 

measure a larger number of sources or, by making many expo- 
sures, to check for short term variability in a single source, 
both important segments of the Rosemary Hill program. 

Photographic photometry is also advantageous in that 
changes in transparency affect both the object and comparison 
sequence (typically a one-half degree or smaller diameter 
field) simultaneously and thus do not seriously affect results 
For photoelectric observations, however, good weather with 
constant, excellent sky transparency is required. Weather 
studies have been conducted at Rosemary Hill with results 
that indicate 66 percent of the nights available are usable 
photographically while only 35 percent are useful photoelec- 
trically (Hackney 1973). 

The photographic program also lends itself to observa- 
tions of faint objects or nebulosity through the use of fine- 
grain emulsions. Finally, many known radio sources which are 
unidentified or of doubtful optical identification are moni- 
tored infrequently to search for possible variations and 
identification of new QSO's. These plates are of similar 
scale and easily comparable to the Palomar Sky Survey to aid 
in identification of objects and to check for variability 
over a twenty-year time span. 

Telescope time at Rosemary Hill is awarded on the basis 
of half nights. The QSO monitoring program is typically 
given observing time equivalent to eight whole nights per 
month divided between twelve observing sessions. 



26 

Telescope and Equipment 

Most observations of the optical brightness of objects 
were made at the f/4 Newtonian focus of the 30-inch Tinsley 
reflector. A few observations of the optical brightness of 
BL Lac (shown in Figure 19) were made with the 18-inch 
Ritchey-Chretien telescope. The Newtonian focus of the 30- 
inch was chosen for several reasons. The first is that the 
field is one degree in diameter; while only one-half is 
usable due to comatic aberration, this is a large enough 
field to include a good comparison sequence for almost all 
objects. In comparison, the Cassegrain field is only one- 
fourth degree in diameter. The exposures at the Newtonian 
focus are also much shorter than would be required for the 
f/16 Cassegrain focus since atmospheric turbulence causes 
each star to exhibit a finite disc, thus increasing exposure 
times for the larger focal length Cassegrain system. A final 
reason is that the physical process of removing the photo- 
electric photometer usually mounted at the Cassegrain focus 
is difficult and time-consuming; this process is avoided by 
using the Newtonian focus. The Newtonian program does have a 
disadvantage in that the moonlit sky is too bright for success- 
ful exposures. Consequently, a Cassegrain camera was developed 
by R. L. Scott to monitor bright objects during moon time. 

As this program has been in operation for several years , 
focus, guiding, exposure, and development procedures have 
already been described thoroughly elsewhere (Folsom 1970; 
Hackney 19 73). These procedures are not repeated here. One 



27 

recent change in the program is that most exposures are now 
made on Kodak 103a-0 plates hypersensitized by baking (Scott 
and Smith 1974) , whereas untreated 103a-0 plates were used 
in the past. 

Iris Photometry 

The QSO photographic plates are reduced on a Cuffey Iris 
Astrophotometer. Since Folsom (19 70) has given a complete 
description of the instrument and its use, a simplified dis- 
cussion will be given here to give the reader a working knowl- 
edge of its operation. The density and diameter of an image 
of a point source on a photographic plate are related to the 
magnitude of the source. The iris photometer transmits a 
light beam through a variable iris diaphragm, the photographic 
plate and the QSO or star image. This beam and a reference 
beam adjusted for each plate alternately enter a photocell. 
By changing the diaphragm opening, which is encoded to a 
digital reading, the currents produced by the two beams can 
be equalized, thus nulling a galvanometer. The iris reading 
at the galvanometer null is the recorded reading. 

The magnitude of an object is related to the iris read- 
ing in the form: 

2 
m = al + bl + c (2.1) 

where m is the magnitude, I the iris reading and a, b, and c 
are constants (Hackney 1972). For each plate the iris read- 
ings of the source and several stars of known magnitude are 



28 

determined. A least squares parabolic fit is made to the 
magnitudes versus iris readings of the known stars. The 
magnitude of the unknown is then determined from this curve 
by substituting the iris reading. The distances on the mag- 
nitude axis of the plotted known stars from the fitted curve 
are known as the "residuals." The error for a magnitude 
determination is taken to be the rms of the residuals. 

A complete listing of the factors causing error in these 
measurements is given by Hackney (1972). The most important 
is fluctuation in the background density of the emulsion. 
The best error obtainable in the University of Florida pro- 
gram is about 0.05 magnitude for bright objects on nights of 
excellent transparency. An average value is 0.10 magnitude. 

A possible effect that should be noted is the apparent 
change in the magnitude of an object due to increasing air- 
mass. Much of the optical energy of QSO's is radiated in the 
blue and ultraviolet regions of the spectrum, the regions in 
which optical radiation suffers the greatest scattering by 
the atmosphere. Hackney (1972) found that for the Seyfert 
galaxy 3C 120 this effect is 0.05 magnitude per airmass, thus 
showing that this effect is not of great importance. 

An important part of the program is detecting variability 
in objects not previously reported as variable. In order to 
detect variability, a comparison sequence of nearby stars of 
known magnitude is needed. The method used in the Rosemary 
Hill program to calibrate a comparison sequence is to expose 
an unknown QSO and its surrounding field on the same plate 



29 

with a field containing stars of known magnitude. These two 
exposures are of identical length and are made only during 
excellent sky conditions. When read on the iris photometer, 
the reference beam is the same for both fields. Each poten- 
tial member of the new comparison sequence is treated as an 
unknown and has its magnitude determined from a calibration 
curve of known stars on the second half of the plate. While 
this process is not as accurate as a photoelectric determina- 
tion of the amplitude of each comparison star would be, it is 
much superior to attempts to determine variations in QSO's 
by "eyeball" methods. 

Data Reduction 

Two programs have been developed by Hackney (19 72) for 
the reduction of iris photometer data. The first is for the 
University of Florida IBM 370/165. This program is most suit- 
able for large amounts of data gathered over one or more 
observing seasons. The comparison star magnitudes and the 
iris readings are entered on punched cards. The program fits 
a linear and a parabolic curve to the data and prints the 
solution which minimizes the error. A linear fit is given to 
those objects with high scatter in the comparison sequence or 
with too few stars for a parabolic fit. The output gives the 
magnitude of the QSO with the residuals for each comparison 
star for every plate, the error for each plate (rms of the 
residuals), and the airmass for each observation, plots a 
light curve, and smooths the comparison sequence by printing 



30 

a new magnitude for each star. The new magnitude is derived 
by averaging the residuals from all plates and adding this 
increment to the original magnitude. The use of this feature, 
combined with accuracy in recentering each field, reduces the 
error caused by field effects in the observing system. In 
addition, the program tests the data for significance of vari- 
ation by use of a chi-squared test of the hypothesis of con- 
stant mean devised by Penston and Cannon (1970) and modified 
by Hackney (19 72). 

The second program is for the Hewlett-Packard 9810A cal- 
culator. This machine provides high accuracy with a fast 
turn- around time, but it can process only one plate at a time, 
thus providing no test for variability. This machine is best 
used when accurate magnitudes are needed quickly, for example 
when an object is suspected of undergoing an outburst. The 
program makes a linear or parabolic fit to the data at the 
discretion of the programmer. The calculator print-out gives 
the magnitude of the object, the residual for each comparison 
star and the error. The program and comparison sequence mag- 
nitudes are entered from magnetic cards. The iris readings 
may be entered manually or electronically through a photometer- 
calculator interface built at the University (Hackney 1972). 

The Polarimetry Program 

As noted in the previous chapter, there is presently 
great interest in polarized emission from QSO's and related 



31 

objects. The nature and distance of these objects are the 
cause of much controversy. Further observational knowledge 
can help reduce the controversy and lead to an understanding 
of the physical nature of these objects. One of the most 
valuable of the observational tools is linear polarization. 
From a knowledge of the spectral shape and variability of the 
linearly polarized continuum one can construct source models 
and even estimate an upper limit to the distance of the object 
(Visvanathan 19 73b) . 

In order to substantiate and extend the work of others 
in this field, linear polarization observations of three 
selected sources were initiated by the author at Rosemary Hill. 
As was mentioned in the remarks concerning the regular moni- 
toring program, most of these QSO's and related objects are 
too faint to observe photoelectrical^ with the 30-inch 
reflector. Unpolarized light incident upon a polarizing ele- 
ment is reduced in intensity by at least 50 percent, and some- 
times as much as 65 percent, upon transmission. This in 
effect reduces the brightness of a source by about one magni- 
tude. For this reason, photoelectric polarimetric observa- 
tions of even the brightest QSO's are not possible with the 
present equipment. As an alternative, a photographic deter- 
mination of linear polarization was decided upon by the author. 
This method has the advantage of being able to measure fainter 
sources. In addition, the equipment for exposing, developing, 
and reducing photographic plates was at hand. The method does 
have the disadvantages of long exposure times and increased 



32 

error compared to photoelectric methods. However, it was 
felt that the importance of adding to the data on such objects 
was enough to outweigh these disadvantages. 

Method of Exposure 

The f/16 Cassegrain focus of the 30-inch reflector was 
used for the linear polarization measurements to avoid spuri- 
ous polarization induced by the flat Newtonian secondary 
mirror. The Cassegrain camera designed by R. L. Scott and 
shown in Figure 2 was employed in the measurements. Most 
exposures were made in the Johnson B magnitude system through 
a sheet Polaroid HN-32 filter and a Schott GG-13 filter on 
Kodak 103a-0 emulsions hypersensitized by baking. After June 
1, 1974, a Polacoat 105UVWRMR polarizing coating on a quartz 
substrate was used as the polarizing element. The spectral 
response of the HN-32 has been determined by use of a Beckman 
spectrophotometer. This curve and the GG-13 spectral response 
from the Kodak "Schott Color Filter Catalog #365e" are shown 
in Figure 18. As can be seen, the polaroid does cause some 
change in the spectral response of the system in the region 
of interest (about 390-480 nm) . It also causes a 60 percent 
reduction in the amount of light reaching the plate, result- 
ing in longer exposure times to reach the plate limit. The 
Polacoat 105UVWRMR response from Polacoat Bulletin P-113- 
3/1/71 is also shown. It is nearly flat in the region of 
interest and also transmits only about 40 percent of the 
incident light. 



T3 

C 
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CD 
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34 




35 

Five exposures are made for one measurement. The first 
is a standard magnitude determination taken without the polar- 
izing filter. The subsequent exposures are made with the 
filter in place. The axis of polarization of the filter is 
known precisely and is initially aligned north-south. The 
dark slides are opened and the first polarization exposure 
is made. The camera, together with the plate and filter, 
which are locked in place, is then rotated counterclockwise 
so that the polarization axis is east-west. The second expo- 
sure is made. The camera is rotated clockwise 45° into a 
northeast-southwest position for the third exposure and then 
counterclockwise 90° to a northwest-southeast position for the 
fourth exposure. These four exposures are made whenever pos- 
sible in one-inch diameter fields on one plate to avoid pos- 
sible errors arising from variations in background conditions 
in the emulsion from plate to plate. However, for some objects 
there are not enough stars to form a comparison sequence in a 
field of this size. In this case, exposures are made in two- 
inch diameter fields on two plates. Possible errors are 
minimized by choosing two plates from the same emulsion batch 
and hypersensitizing and developing them together. 

Focusing is accomplished by the knife-edge method. The 
Cassegrain secondary mirror is motorized, and is moved in and 
out to achieve the best focus. Guiding is accomplished with 
the Cassegrain guide box, which contains a large diagonal 
mirror to deflect a portion of the beam at a right angle to 
the optical axis into a microscope. The procedure is 



36 

thoroughly described by Hackney (1972). The development of 
plates follows exactly the same procedure as that described 
by Folsom (19 70) for the plates taken at the Newtonian focus. 

Data Reductio n 

The magnitude of an object is determined for each of the 
four polarization exposures by exactly the same method as 
that described above for normal photographic exposures, using 
the iris astrophotometer and the HP 9810A calculator. This 
procedure gives four magnitude values, one for each exposure. 
The values are reduced in pairs according to the procedure 
given by Hall and Serkowski (1963). A brief description of 
this procedure follows. 

The reduction is accomplished with the use of Stokes 
parameters expressed in stellar magnitudes. The Stokes param- 
eters for linearly polarized light are: 



1=1 + I . (2.2) 

max min v 

Q = (I - I . ) cos29 = PI cos26/100 (2.3) 

x v max mm' J 

U = (I - I . ) sin26 = PI sin26/100 (2.4) 

v max mm' v 

V = (2.5) 



"where I is the maximum intensity observed using a perfect 

max ' or 

polarizer and I • is the intensity observed with the polar- 

r mm J c 

izer rotated by ninety degrees. The position angle of I ax 
in equatorial coordinates is 9. The parameter V is a measure 
of elliptical polarization and is zero for light showing only 



37 

linear polarization. Observational results indicate that the 
presence of optical elliptical polarization in QSO's is doubt- 
ful (Landstreet and Angel 1972). 

The ratio of I to I - can be expressed in terms of 
max mm r 

stellar magnitudes: 



p = 2.5 log j 



max 



(2.6) 



min 



The linearly polarized percentage of the total flux is 



P = 100 



fl - I . I 

max mm 

I + I . 
1 max mm 



I 



(2.7) 



If one expands I /I . loP' ' in a Taylor series and 

r max min ' 

assumes that p<<2.5 then it can be shown 



P z 20 (In 10)p = 46.05 p 



(2.8) 



Equation (2.8) becomes invalid for P of about 40 percent or 
more . 

The Stokes parameters in magnitudes are given by 



Px E P cos2e = 0.4605 1 



U 



Py - P Sin29 = 0.4605 I 



(2.9) 
(2.10) 



and are derived by dividing equations (2.3) and (2.4) by 
equation (2.2) and substituting from equation (2.8). These 
parameters are given in terms of the measured magnitudes by 



p x = m(90) - m(0) 
p = m(135) - m(45) 



(2.11) 
(2.12) 



38 

where the arguments give the alignment of the polaroid axis 
relative to the celestial sphere during the exposure. After 
obtaining p and p , p and are determined by solving equa- 
tions (2.9) and (2.10) simultaneously. The degree of polariza- 
tion P is then determined using equation (2.8). 

Error is determined by using the rms of the residuals as 
explained above. These are taken to be an estimate of the 

error e in p and p . The error is independent of the true 
x y 

values of p and p as explained by Hall and Serkowski (1963) 

The true value of the amount of polarization p is re- 

r r o 

lated to the most probable observed value p by 

r 2 ! 2.1/2 . « 

(P + j ire ) for p Q z 

P = (2.13) 

2 2 

p + e for p o >> £ 

The rms deviation 6 of the observed polarization from the 
true value p is 

e/T for p ~ 

« p - ° (2.14) 

K e for p >>e 

The mean error e of the amount of polarization, which is the 
rms deviation of p from p , is 



o 



1 ^1/2 



e(2 - i- tt)-^ for p = 

e D = ° (2.15) 

^ e for p >>e 

r o 

while the mean error e fi of the position angle 9 is 

rad for p ~ 

/IT ° 

e Q = (2.16) 

\ | rad for Po >>£ 



CHAPTER III 
RESULTS OF THE QSO MONITORING PROGRAM 

The present program of photographic observations of 
extragalactic radio sources at Rosemary Hill includes over 
one hundred fifty sources. In such a large program many 
objects are of necessity observed infrequently, with perhaps 
as few as one or two plates per year. The more interesting 
and active objects may be observed weekly or more often during 
periods of activity. The results of observations of sixteen 
of these objects have been reported recently (Scott e_t al_. 
1973b), while the results of observations of several more 
objects will be reported elsewhere. All of the remaining 
objects currently being monitored are reported in this chapter. 

It should be noted that the QSO monitoring group consists 
presently of four members under the direction of Dr. A. G. 
Smith. Exposure, development, and reduction of photographic, 
plates is a cooperative effort of all members of the group. 
Therefore, many of the data presented in the present chapter 
were taken by present and former members of the group and 
their work is referenced whenever appropriate. However, in 
addition to participation in the joint group efforts, many 
calibrations, the refinement of comparison sequence magnitudes, 
the compilation and plotting of data, the research into the 



39 



40 

literature, and the conclusions drawn in this chapter and 
Chapter V are solely the work of the author. 

In the following tables all radio sources identified by 
Parkes catalog numbers have had the prefix PKS dropped and 
are identified by the right ascension and declination suffix. 
All sources from other catalogs are identified by their full 
designations . 

In Table 1 are listed all objects which have had a photo- 
electric comparison sequence published, or which have had a 
comparison sequence established by photographic transfer from 

selected areas as described in Chapter II. All objects for 

2 
which the x statistic indicates eighty percent or greater 

confidence level for variability are discussed separately 
with light curves given. These constitute 50 percent of the 
calibrated sample. This large number is probably caused by 
the practice by the author and coworkers of giving priority 
to establishing comparison sequences for those objects which 
appear to exhibit interesting activity. There is some evi- 
dence that all QSO's may be variable on a very long time 
scale (years) (Lu 1972; Penston and Cannon 1970). 

In Table 2 are listed those objects for which "tertiary" 
comparison sequences have been established by photographic 
transfer from other QSO fields, rather than by direct cali- 
bration. These are easily accomplished since two QSO fields 
are exposed on one plate as standard practice. Many such 
tertiary calibrations have proven to lead to inaccurate 
results when compared to results from later direct 















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48 

calibrations. Direct calibrations from selected areas have 
not been made for these objects because of the large number 
of program objects and the infrequency of nights with the 
excellent transparency conditions required. However, it is 
the opinion of the author that the use of these tertiary 
calibrations gives a better indication of the behavior of an 
object than a visual inspection gives. About 39 percent of 
these objects appear to be variable as determined by comparing 
their rms error to the range of observed magnitudes. 

In Table 3 are listed those objects with no calibrations 
and usually with very few plates taken. Some may be deleted 
from the program due to lack of optical identification or 
lack of interesting activity. Others will probably be cali- 
brated and observed more regularly in the future. A few 
objects in this table, such as PKS 2135-14, are well-known 
objects which are recent additions to the Rosemary Hill pro- 
gram. 

Discussion of Individual Objects 

OB 33 8 

Kraus e_t al. (1968) found that OB 338 has a slightly 
peaked radio spectrum with a maximum at 500 MHz. Thompson 
et al . (1968) found two stellar objects near the radio posi- 
tion. Folsom (1970) proposed the bluer of the objects as the 
optical counterpart and reported the Rosemary Hill observa- 
tions. Hackney (1973) reported a calibration of the 









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54 



surrounding field by photographic transfer from SA 20 (Brun 
1957) and updated the observations. 

The object is variable with greater than 99.5 percent 
confidence with an rms error of 0^2 . The light curve in 
Figure 3 shows a well-defined peak in 1970 with no further 
flares of that size since. However, due to the few plates 
taken since 1970, this type of burst may have been repeated 
since that time. There seems to be no indication of a slowly 
varying component, the minimum brightness being nearly con- 
stant at about 19.5. A sequence finder is given by Hackney 
(1973). 

01 318 

Kraus e_t al_. (1968) found 01 318 to have a radio spectrum 
peaked at 1600 MHz. Thompson et_ al_. (1968) reported a pair of 
red and blue objects within the radio error box. Blake (1970) 
proposed the blue object as the optical counterpart and this 
was confirmed by Folsom (1970). Burbidge and Strittmatter 
(1972) found a redshift of z = 1.620, confirming the object 
as a QSO. Hackney (19 73) reported a comparison sequence 
established by photographic transfer from SA 50 (Brun 1957) 
and indicated that the object was variable with 99.2 percent 
confidence. The magnitudes of the comparison sequence were 
recalculated more accurately by the author, using an elec- 
tronic calculator. 

The data indicate that 01 318 is variable at the 80 per- 
cent confidence level with an rms error of 0.15. The light 



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57 



curve is shown in Figure 3. An 0™8 burst was observed in 
late 1970. More recently the object shows a 0.25 decline in 
1973-1974 with a 0™5 burst indicated by one plate. There 
appears to be no long-term change in minimum brightness. A 
sequence finder is given by Hackney (1973). 



OK 290 



This object is dissimilar to the two previous ones in 
that the radio spectrum rises smoothly to 10 GHz with no 
maximum observed (Kraus et_ al_. 1968). The optical counterpart 
was identified with a 17th magnitude stellar object (Thompson 
et al_. 1968; Blake 1970), with a finder published by Blake. 
A redshift of 1.620 was observed by Burbidge and Strittmatter 
(1972). Medd et_ al . (1972) found a brightening of 1 flux unit 
at 2.8 and 4.5 cm between 1967 and 1970. Stull (1972) found 
a decline in flux density of 20 percent between January and 
August of 1971 at 3.75 cm. Previous Rosemary Hill observa- 
tions have been reported by Folsom (1970) and Hackney (1973). 
Recently a comparison sequence has been calibrated by a photo- 
graphic transfer from SA 54 (Brun 1957). 

The data indicate that OK 290 is variable with greater 
than 99.5 percent confidence. The average rms error of a 
measurement is 0™09. The light curve in Figure 3 shows a 
"wave-like" shape with maxima in late 1971 and early 1973. 
There seems to be good indication of short-term activity with 
variations of as much as 0.3 observed within a two-week 
period. A sequence finder is given by Hackney (1973). 



58 

OQ 20 8 

This object has a very sharply peaked radio spectrum with 
a maximum at 5 300 MHz (Kraus e_t al_. 196 8; Thompson et al. 
1968). Blake (1970) identified the source with a 14th magni- 
tude Seyfert galaxy. Medd et_ al_. (1972) found that the 
object remained at almost constant brightness at 2 . 8 and 4.5 
cm from 1967 to 1971. Lu (1972) found that the object is an 
optical variable showing a 0?5 peak in 1970. Craine and 
Warner (1973), using the Harvard plate collection, reported 
that OQ 208 was variable with a range of nearly one magnitude 
between 1938 and 1953. They found no variation over a period 
of four months in early 1972. 

The Rosemary Hill data are shown in Figure 5. A compari- 
son sequence was calibrated with a photographic transfer from 
M 3 (Sandage 1953) . The object is variable with greater than 
99.5 percent confidence with an average rms error of .'1 . 
The Florida data agree in general behavior with Lu's (1972) 
results. The 1970 Florida light curve shows two peaks 0:6 
above minimum with a 0.4 drop between them. In 19 74 there was 
a 0.5 flare followed immediately (within two weeks) by an 0.8 
decrease. Inspection of the plates and images leads the 
author to the conclusion that this change is real. The 
object remained at a minimum for four weeks with one 0.4 
flare occurring, rising rapidly and declining within one day. 
Unfortunately, there are no 1972 observations to compare with 
Craine and Warner (1973). A comparison sequence finder is 
shown in Figure 4. 



Figure 4. Comparison sequence of OQ 208. This photo- 
graph is reproduced from a plate taken on the 
night of 13 April 1974, at Rosemary Hill. 
North is at the top and east is to the left. 
Comparison star photographic magnitudes are 
(1) 16.40, (2) 16.52, (3) 14.77, (4) 16.56, 
(5) 15.65, (6) 16.08, (7) 15.56, (8) 15.13, 
(9) 15.74, (10) 15.95. 



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63 

OX 074 

Kraus et_ al_. (1968) reported OX 074 to have a peaked 
spectrum with a maximum at 2500 MHz. Thompson et al. (1968) 
identified the source with a blue stellar object. Stull 
(1972) found a flux density increase of about 35 percent at 
3.75 cm in two months of 1971. Previous Rosemary Hill data 
have been reported by Folsom (1970) and Hackney (1972), who 
identified the object as an optically violent variable (OW) . 
The light curve in Figure 5 shows a 1™5 decrease from 19 70 
to 1971 with short-term bursts superimposed on the decline. 
The object remained at a minimum of about 17™7 through late 
1971 and 1972 with the short-term activity still present. 
In ]973 the object underwent two 1™0 bursts accompanied by a 
0.15 rise in the minimum brightness. A sequence finder is 
given by Hackney (1972). 

PKS 0202-17 

An optical identification and finder for this source 
were given by Bolton and Ekers (1966). The spectral index 
between 1410 and 2650 MHz is fairly flat at . 2 (Ekers 1969). 
Kinman et_ al_. (1967) reported a redshift of z = 0.717, thus 
confirming the object as a QSO. Stull (1972) found that PKS 
0202-17 is a probable variable at 8000 MHz (3.75 cm), with 
flux density increases of 20 and 25 percent recorded in one- 
month periods. 

A comparison sequence was calibrated at Rosemary Hill 
with a photographic transfer from SA 118 (Brun 1957). The 



64 



object is variable with 92 percent confidence, with an rms 



,m 



error of 0.10. The Florida data typically show small ampli- 
tude (<0?5) variations on a time scale of several weeks 
(Folsom 1970). The brightness shows a long-term decline of 
about 0?3 from 1969 to the present. The light curve is shown 
in Figure 5. A sequence finder is shown in Figure 6. 

PKS 0222-23 

The source was identified as a possible QSO and a finder 
was published by Bolton et_ al. (1965a). Ekers (1969) reported 
that the spectral indices between 408 and 1410 MHz and between 
1410 and 2650 MHz are -0.9 and -0.1, respectively. 

A comparison sequence surrounding the source was cali- 
brated with a photographic transfer from SA 119 (Brun 1957). 
As the light curve in Figure 7 shows, the object had a major 
flare in 1971 and has since declined 1T5 . The minor flares 
in the declining phase after the major outbursts are typical 
of an optically violent object (Hackney 1972). The average 
rms error is 0™15 and the variability is real with greater 
than 99.5 percent confidence. A sequence finder is shown in 
Figure 8. 

PKS 0336-01 

Bolton and Ekers (1966) identified the optical counter- 
part of the radio source (also designated CTA 26) and pub- 
lished a finding chart. The radio spectrum is inverted, the 
spectral indices being -0.5 and 0.7 from 400 to 1410 MHz and 
1410 to 2650 MHz, respectively (Ekers 1969). Kinman e_t al. 



Figure 6. Comparison sequence of PKS 0202-17. This 

photograph is reproduced from a plate taken 
on the night of 10 December 1971, at Rosemary 
Hill. North is at the top and east is to the 
left. Comparison star photographic magnitudes 
are (1) 16.15, (2) 16.90, (3) 16.49, 
(4) 16.66, (5) 17.83, (6) 17.13, (7) 16.06, 
(8) 18.14, (9) 18.10, (10) 16.18, (11) 17.05. 



II I 



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66 



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Figure 8. Comparison sequence of PKS 0222-23. This 

photograph is reproduced from a plate taken 
on the night of 12 December 19 71, at Rosemary 
Hill. North is at the top and east is to the 
left. Comparison star photographic magni- 
tudes are: (1) 16.77, (2) 16.57, (3) 16.08, 
(4) 15.80, (5) 15.81, (6) 15.26, (7) 15.45, 
(8) 14.68, (9) 16.82, (10) 15.70, (11) 14.50, 
(12) 15.59. 



■0«.' , 2 



70 



8 



71 

(1967) reported an accurate optical position and the results 
of photoelectric UBV photometry. Gardner and Whiteoak (1969) 
found 2 percent linear polarization between 11 and 20 cm, 
while Aller (1970) found 6 percent linear polarization at 
3.75 cm. Wills (1971) found day-to-day variations of two to 
four percent at 2700 MHz (11 cm) , and derived the angular 
diameter to be no more than 0.001 for 66 percent of the radi- 
ation. Medd e_t al_. (1972) observed month-to-month variability 
at 2 . 8 and 4.5 cm. Dent and Kojoian (1972) found PKS 0336-01 
to be variable at 7800 MHz (3.85 cm). 

A comparison sequence in the field of PKS 0336-01 was 
obtained at Rosemary Hill (Hackney 1973) by photographic 
transfer from the field of SA 95 (Brun 1957). Hackney (1973) 
reported the object to be variable. The up-to-date observa- 
tions indicate that the object is variable with greater than 
99.5 percent confidence with an average rms error of 0.13. 
It is interesting to note that the general shape of the light 
curve (Figure 7) through 19 71 agrees with the radio curves of 
Medd et_ al. (1972) and Dent and Kojoian (1972). Since the 
maximum of 19 70 the object reached an 0™8 minimum in 19 72 
and brightened by 0.4 in 1973. A sequence finder is given 
by Hackney (1973) . 

PKS 0458-02 

The optical counterpart of the radio source was identi- 
fied by Bolton and Ekers (1966). The object was identified 
as a QSO by Ekers (1969). The radio spectrum is fairly flat 



72 

with spectral indices of -0.3 and -0.1 between 408 and 1410 
MHz and between 1410 and 2650 MHz, respectively. Medd et_ al_. 
(1972) found that the object varied 25 percent over a period 
of a year at 2.8 and 4.5 cm. 

A comparison sequence was calibrated at Rosemary Hill 
with a photographic transfer from SA 97 (Brun 1957). The 
object is variable with greater than 99.5 percent confidence 
with an average 0.14 rms error. All points lie within a 0^3 
scatter, except for one plate in November, 1973. This plate 
was taken under poor sky conditions and, if eliminated, the 
probability of variation is reduced to 18 percent. Thus, the 
variability should be considered doubtful pending further 
observations. The light curve is shown in Figure 8 and the 
comparison sequence is shown in Figure 9. 

PKS 0518+16 

The optical counterpart of the radio source (also desig- 
nated 3C 138) was identified by Clarke et al. (1966). The 
spectral index is -0.7 between 408 and 1410 MHz (Ekers 1969). 
The object is at low galactic latitude (-11°) but a redshift 
of z = 0.760 was measured by Burbidge and Kinman (1966) and 
Lynds e_t al_. (1966) , confirming the object as a QSO. Aller 
(1970) reported 9 percent linear polarization at 3.75 cm. 
Medd et_ al_. (1972) found little change in brightness at 2.8 
and 4.5 cm. Peach (1969) measured a 0.1 optical variation 
over one year. 



Figure 9. Comparison sequence of PKS 0458-02. This 

photograph is reproduced from a plate taken 
on the night of 13 October 1972, at Rosemary 
Hill. North is at the top and east is to 
the left. Comparison star photographic 
magnitudes are: (1) 19.02, (2) 18.65, 
(3) 19.11, (4) 19.30, (5) 18.71, (6) 17.86, 
(7) 19.25, (8) 18.16, (9) 18.50, (10) 19.31. 



74 




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75 

PKS 0518+16 has been monitored for four years at Rosemary 
Hill. A comparison sequence was obtained (Hackney 1973) by 
photographic transfer from SA 74 (Brun 1957). The object is 
variable with greater than 99.5 percent confidence with an 
average rms error of 0?15. The light curve in Figure 10 indi- 
cates that the object should be classified as a violent vari- 
able (OW) . Changes of 0.8 within one month were noted in 
1970-1971. The object is seen to be gradually brightening 
since 1970 with repeated short-term bursts superimposed on 
the long-term behavior. A sequence finder is given by Hackney 
(1973). 

PKS 1252+11 

The optical counterpart of PKS 1252+11 was identified by 
Bolton et_ al_. (1965a). The object was definitely established 
as a QSO with the discovery of a redshift z = 0.871 by Lynds 
et_ a_l. (1965) and Schmidt (1966). The spectral index is 0.1 
between 1410 and 2650 MHz. Penston and Cannon (1970) obtained 
three plates of the object and concluded that optical varia- 
bility was doubtful. 

A comparison sequence (shown in Figure 11) was established 
in the field of the object, with a photographic transfer from 
SA 81 (Brun 1957). The data indicate that the object is vari- 
able with 9 8 percent confidence. The rms error is 0™09. 
The light curve (Figure 10) shows a total range of 0™55 with 
changes of 0.35 occurring within six weeks. There seem to be 
no long-term trends. 



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Figure 11. Comparison sequence of PKS 1252+11. This 
photograph is reproduced from a plate 
taken on the night of 3 March 19 74, at 
Rosemary Hill. North is at the top and 
east is to the left. Comparison star 
photographic magnitudes are: (1) 15.67, 
(2) 15.56, (3) 16.56, (4) 17.36, (5) 16.81 
(6) 14.89, (7) 15.41, (8) 15.29, (9) 16.40 
(10) 18.05, (11) 17.77. 



79 



7 



8 



II 



10 



9 



80 

PKS 1547+21 

Shimmins and Day (196 8) reported that the object shows 
a departure from a power law spectrum at 400 MHz and concluded 
that the object is a possible QSO. The photographic observa- 
tions of Folsom et_ al_. (1971) indicated that the object had 
shown a brightness change between the Palomar Sky Survey plate 
and 1970 but was not variable during their observations. 

A photographic transfer from SA 82 (Brun 1957) was used 
to calibrate a comparison sequence for PKS 1347+21 at Rose- 
mary Hill. The object is variable at a confidence level of 
89 percent with an average rms of 0.12. The light curve in 
Figure 10 indicates that the object is slowly varying with a 
0.4 decline observed between 1969 and 1972. A sequence 
finder is shown in Figure 12. 

PKS 1607+26 

Shimmins and Day (196 8) reported that the radio source 
has a spectral index of 0.4 at 400 MHz. An optical identifi- 
cation was published by Merkelijn (1968). Folsom et al . 
(1971) reported that the object is moderately variable 
optically. 

A comparison sequence in the field of PKS 1607+26 was 
established at Rosemary Hill (Folsom 19 70) with a photographic 
transfer from SA 61 (Brun 1957). The object is variable with 
95 percent confidence with an average rms error of 0.10. The 
object shows a total range of about 0.65 with one change of 
0?5 observed. The average brightness increased 0.2 from 



Figure 12. Comparison sequence of PKS 1347+21. This 

photograph is reproduced from a plate taken 
on the night of 29 April 1974, at Rosemary 
Hill. North is at the top and east is to 
the left. Comparison star photographic 
magnitudes are: (1) 17.02, (2) 16.07, 
(3) 13.88, (4) 16.69, (5) 16.37, (6) 15.37, 
(7) 15.98, (8) 15.71, (9) 15.57, (10) 14.40. 



82 




• Q 



3 



10 



4 



.5 




83 



1969 to 1974. The light curve is shown in Figure 13. A 
sequence finder is given by Folsom (1970). 

PKS 2209+08 

The optical counterpart of the radio source was identi- 
fied by Clarke et_ al. (1966). Ekers reported the spectral 
indices to be -0.6 and -0.2 between 40 8 and 1410 MHz and 
between 1410 and 2650 MHz, respectively. Miley (19 71) con- 
firmed the object as a QSO. Hackney (1973) gave an optical 
light curve. 

A comparison sequence in the field of PKS 2209+08 was 
established by photographic transfer from SA 114 (Brun 1957). 
The object is variable with greater than 99.5 percent confi- 
dence with an average rms error of 0^13. The light curve is 
shown in Figure 13. There is evidence of one major flare of 
about 0.5 in 1970 and an indication of another in 1973. This 
object also seems to experience antiflares in which the bright 
ness drops below the average value by as much as . 6 . In 19 70 
and 1973 this sudden drop immediately preceded a flare or 
possible flare. A sequence finder is given by Hackney (1973). 

PKS 2345-16 

Bolton and Ekers (1966) gave an optical identification 
of the radio source. A redshift of z = 0.6 was determined by 
Schmidt (Ekers 1969). The radio spectrum is flat, with a 
spectral index of -0.05 between 11 and 21 cm, and the object 
is variable at radio wavelengths (Kellerman and Pauliny-Toth 
1968). Folsom (1970) described two 1™5 flares occurring in 



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86 

late 1969 and early 1970. Hackney (1972) described a one- 
magnitude flare seen in the fall of 1971 and postulated two 
components. The first corresponds to long period (months or 
years) changes in brightness and the second corresponds to 
rapid (days or weeks) changes of slightly smaller amplitude. 
Medd et_ al_. (1972) found the source to be variable at 2.8 and 
4.5 cm with one large peak observed at a time corresponding 
to the optical flares observed by Folsom (1970). 

A comparison sequence was calibrated at Rosemary Hill 
(Folsom 1970) with a photographic transfer from SA 116 (Brun 
1957) and it is shown by Hackney (1972). The light curve is 
shown in Figure 13. The confidence level for variability is 
greater than 99.5 percent with an average error of 0.10. 
Computer refinement of the comparison sequence leads to the 
conclusion that the peak observed by Hackney (1972) in 1971 
was not as large as previously indicated. The source reached 
minimum brightness in 19 71 and began a long-term brightening 
which continued through 1972 and 1973. Short-term activity 
seems to be superimposed on the long-term behavior most of 
the time. A peak of about 0.5 was observed in September of 
1973. 

NRAO 140 



The optical counterpart of NRAO 140 was identified by 
Kristian and Sandage (1970). An alternate designation is OE 
355. A redshift of z = 1.263 was measured by Burbidge and 
Strittmatter (1972). Kraus et al. (1968) reported a radio 



87 

spectrum peaked at about 700 MHz at a level of 4.1 flux units. 
Medd et_ al_. (1972) reported the radio intensity to be variable 
at 2 . 8 and 4.5 cm with a peak occurring in 1969. 

A comparison sequence in the field of the object was 
established by photographic transfer from SA 48 (Brun 1957) 
and is shown in Figure 14. The object is variable with 
greater than 99.5 percent confidence with an rms error of 
m ll. The light curve is shown in Figure 15. It should be 
noted that only seven plates have been taken, thus leaving 
some uncertainty about the validity of the variability. 

NRAO 512 

The optical counterpart of the radio source NRAO 512 was 
identified by Folsom (1970) from the radio position. Locke 
et al . (1969) reported that the source was faint (less than 
two flux units) at 2 . 8 and 4.6 cm. These authors observed 
that the flux density declined by a factor of two in a five- 
month period in 1968 at both wavelengths after being constant 
for eight months. Stull (1972) observed a 0.5 flux unit 
increase in 1971 at 8000 MHz (3.75 cm). Medd et_ al . (1972) 
observed a one flux unit peak in 19 70 at 2.8 and 4.5 cm in 
addition to the decline previously reported by Locke et_ al . 
(1969). 

The observations at Rosemary Hill show that the source 
is a violent optical variable. The light curve is shown in 
Figure 15. The object is variable with greater than 99.5 
percent confidence with an average rms error of 0.11. The 



Figure 14. Comparison sequence of NRAO 140. This 

photograph is reproduced from a plate taken 
on the night of 30 October 19 73, at Rosemary- 
Hill. North is at the top and east is to 
the left. Comparison star photographic 
magnitudes are: (1) 17.64, (2) 16.79, 
(3) 18.08, (4) 18.24, (5) 17.71, (6) 15.42, 
(7) 16.03, (8) 17.01, (9) 16.15, (10) 17.41. 



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92 

source brightened by more than one magnitude between May 30 
and June 2, 19 70. By June 5 the optical intensity had 
returned to its previous level. This was followed by a long- 
term (five-month) decline in average brightness with oscilla- 
tory variations superimposed. The peak seems to follow the 
radio flare at 4.5 cm by four weeks and the flare at 2.8 cm 
by six weeks. This behavior was reported by Folsom (1970) 
and Folsom et al_. (1970). Hackney (1972) described the mini- 
mum and subsequent rapid 1™5 flare occurring in 19 71, again 
followed by an oscillatory decline. In 1972, the object 
remained 0.5 above minimum with two flares, one of 0™6 , 
observed. The object declined to near minimum brightness in 
1973 and appears to be brightening in 1974 with one 0™5 flare 
observed. The object does not seem to be quiescent at any 
time. There are rapid (daily or weekly) changes of 0™3 or 
more even at minimum brightness, similar to the behavior of 
PKS 2345-16 and OX 074. A sequence finder is given by Hackney 
(1972). 

3C 371 



The radio source was identified as an N galaxy by Wyndham 
(1966). Dibai and Esipov (1969) reported a redshift of z = 
0.0457. Folsom (1970) plotted a light curve from 1895 to 
date, including Usher's (1969) search of the Harvard and 
other plate collections and the Rosemary Hill data from 1968 



m r 



to 1970. Oke (1967), using the 200-inch, found a 1.0 brighten 
ing between 1965 and 1967 with 0™! to 0™2 variations 



95 

superimposed on a time scale of a few days. Medd et_ al . 
(19 72) observed a shallow 0.3 flux unit minimum in the radio 
emission at 2.8 and 4.5 cm in early 1969. The flux density 
is almost the same at both wavelengths. 

A comparison sequence in the field of the object was 
obtained by photographic transfer (Folsom 1970) from the 
globular cluster M 13 (Savedoff 1956). The Rosemary Hill 
light curve is shown in Figure 16. The object is variable 
with greater than 99.5 percent confidence with an average rms 
error of 0?13. In December 1968 and 1969, large cyclic varia- 
tions with an amplitude of 0.6 were observed (Folsom 1970). 
In 1970 these large variations were not present. Instead, a 
six-month increase in brightness of 0.7 was observed with 0.5 
oscillations superimposed. In 1971 one large flare of 0.6 
was observed while the minimum brightness had declined to the 
early 1970 level. In 1972 and 1975 the object was near a 
maximum in brightness with a 0.45 flare in 1972. The object 
appears to be declining toward a minimum in 19 74. A sequence 
finder is given by Folsom (1970). 

5C 446 

This object, which has a redshift z = 1.404, was one of 
the first QSO's classified as an optically violent variable 
(Penston and Cannon 1970). Medd e_t al_. (1972), observing at 
2.8 and 4.5 cm, found a slow 2.5 flux unit increase from 
1966 to 1970 with a suggestion of a decline beginning in 
19 70. They also report the flux density to decline with 



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96 

increasing frequency. Dent and Kojoian (1972) , observing at 
7.8 GHz, also found the flux density to increase from late 
196 8 to 19 70 with the object remaining near a maximum for the 
last half of 1970 and the first half of 1971. Penston and 
Cannon (1970) found the object to be an optical variable with 
a three-magnitude range in brightness, with the light curve 
showing flares and antif lares. 

The Rosemary Hill observations are shown in Figure 16. 
There is an indication of an increase in brightness from 1968 
to 19 70, in agreement with the radio results. In addition, a 
sharp 0.7 spike was seen in 1969. These points are photo- 
graphic magnitudes. Beginning in 1971 all exposures were 
made in the Johnson B magnitude system. The object declined 
to a minimum in 1972 and began to brighten in 1973. It 
showed a very sharp spike in 19 74, 3.0 above the minimum, 
immediately followed by a rapid decline. The object is vari- 
able with greater than 99.5 percent confidence with an rms 
error of 0.1. The comparison sequence was determined photo- 
electrically by Angione (1971) and is shown in Figure 17. 



Figure 17. Comparison sequence of 3C 446. This photo- 
graph is reproduced from a plate taken on 
the night of 12 October 1974, at Rosemary 
Hill. North is at the top and east is to 
the left. Comparison star "B" magnitudes 
are: (1) 15.79, (21 16.10, (3) 16.48, 
(4) 16.88, (5) 13.68, (6) 15.61, (7) 17.26. 
The magnitudes of stars 1, 3, 5, and 6 are 
by Angione (1971). The magnitudes of stars 
2, 4, and 7 were derived from Angione' s 
magnitudes by a photographic transfer as 
described in Chapter II. 



98 




CHAPTER IV 

POLARIMETRIC OBSERVATIONS OF 3L LACERTAE 
3C 120, AND 3C 273 



Since June, 1973, observations of the linearly polarized 
optical components of BL Lac, 3C 120, and 3C 273 have been 
conducted at Rosemary Hill Observatory by employing the 
method described in Chapter II. The results of these obser- 
vations are described in the present chapter. Previously 
published results from the literature are included for com- 
parison. A discussion of the sources of error peculiar to 
this method of determination of linear polarization follows 
in the first section. 

Error 

Since the preliminary reduction of the polarization, 
plates followed the same procedure as that for the standard 
photographic plates, the same types of error occur. These 
have been discussed thoroughly by Hackney (1972). In addi- 
tion, four other systematic errors may occur. The first is 
color changes in the comparison stars caused by the polariz- 
ing filter. This is probably small since the color responses 
of the Polaroid HN-32 and Polacoat 105UV shown in Figure 18 
are flat in the region of interest. A second possible major 



99 



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102 

source of error would be a real, short-term (^15 minute) 
change in the brightness of a source occurring during a polar- 
ization measurement. BL Lacertae is the only one of the three 
objects presently observed that has been suggested to show 
such changes. This possibility will be discussed further with 
the results of the polarization observations of BL Lac. Third, 
the comparison stars may be linearly polarized due to either 
intrinsic or interstellar mechanisms. This is probably a 
small effect. For example, the interstellar polarization is 
one percent or less in the region of BL Lac (Visvanathan 1969) . 
In any case, any permanent polarization in any of the individ- 
ual comparison stars would appear as a change in the expected 
magnitude of the star and thus is eliminated by the smoothing 
of the magnitudes of the comparison stars as described in 
Chapter II. Finally, one cannot dismiss the possibility of 
instrumental polarization, even given the symmetry of the 
Cassegrain system. This possibility was investigated by mea- 
suring stars for which other observers have found near- zero 
amounts of linear polarization. These measurements were made 
using the HN-32 polarizer and employing the method described 
in Chapter II. The results of these observations are summar- 
ized in Table 4, where P-. is the measured polarization taken 
from the literature (Serkowski 1968) and ? 2 is the percentage 
polarization measured at Rosemary Hill Observatory. 

As can be seen from Table 4, the difference of P 2 from 
P. is in every case less than the average value of e = 4.6 
percent (0?!) due to known errors in the photographic process 



103 



(Hackney 1972). Thus, one may conclude that there is no 

instrumental polarization greater than the uncertainty im- 
posed by the observational error. 

Table 4 

Sources with Known Polarization 

Object Date (U.T.) ? x (%) p 2 ^) 

a Aql 15 May 1973 0.05 4.6 

61 Vir 20 April 1973 0.03 1.4 

2 5 May 19 73 0.03 3.7 

2 5 May 19 73 0.03 2.2 



BL Lacertae 

BL Lac, the prototype for the Lacertid class of objects, 
shows a strong and variable linearly polarized component. 
Olsen (1969), observing at 3.75 cm, found that the polarized 
flux varied between 0.2 and 0.4 flux units in 1968-1969. The 
position angle showed sudden departures from and returns to 
a certain value. Visvanathan (1969; 1973a) found the optical 
linearly polarized component to vary between 3 and 11 percent 
The degree of polarization remained nearly constant between 
582 and 798 nanometers during one night's observations. 



104 

Visvanathan also observed a tendency toward an increase in 
the polarized flux accompanied by a change in position angle 
whenever the total flux increased. 

The present observations are summarized in Table 5 and 
shown graphically in Figure 19 , with the total flux in the 
Johnson B magnitude system included. The total flux gradu- 
ally declined toward a minimum during the summer of 19 73 with 
a 0.5 flare observed in late October. In 1974 the total flux 
declined from May to July and remained at a minimum after- 
wards. The observations indicate that the object exhibited 
significant linear polarization even at minimum brightness. 
The possible exception to this statement is the measurement 
of 2 7 May 19 74, where the amount of polarization was five 
percent, the magnitude of the observational error. It is 
interesting to note that the position angle, 8, on this date 
is about 50 degrees from +45 degrees, the average of the other 
observations. This is about twice the departure from the 
average shown by any other measurement. There is no indica- 
tion of increased polarized flux with increased total flux. 
However, the burst of 30 October 1973 was accompanied by a 
large change in 6. 

Visvanathan (1973a) has suggested that one observes the 
resultant polarization of several active spots, each of which 
is highly polarized with a different position angle. The 
mechanism producing the radiation is the synchrotron process. 
The change in polarized flux and position angle between 27 May 
and 11 September 1974 is suggestive of the flaring of one such 



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106 

active spot or, conversely, the possible decline of another 
region of the source so that one region is more prominent. 
The latter supposition is supported by the fact that the 
total flux had decreased between the two measurements. 











Tabl 


e 5 












Polarimetric 


Observations 




Ob- 


ject 


Date (U 


•T.) 


Mag (B) 


P (%) 


(Degrees) 


BL 


Lac 


30 


June 


1973 


15.24 


6.91 


45.0 






6 


July 


1973 


15.61 


19.93 


48.3 






5 


Sept 


1973 


16.42 


8.15 


21.4 






30 


Sept 


1973 


16.40 


12.39 


69.0 






31 


Oct 


1973 


16.01 


12.70 


21.8 






28 


May 


1974 


15.70 


5.15 


-84.9 






12 


Sept 


1974 


16.24 


11.09 


47.4 


3C 


120 


27 


Oct 


1973 


15.46 


12.43 


-45.0 






3 


Nov 


1973 


15.35 


2.91 


54.2 






20 


Nov 


1973 


15.42 


13.88 


42.1 






24 


Nov 


1973 


15.33 


1.38 


45.0 






18 


Dec 


1973 


15.50 


1.66 


-16.9 






10 


Feb 


1974 


14.71 


13.39 


76.7 






14 


Feb 


1974 


15.09 


1.46 


35.8 






19 


Mar 


1974 


14.98 


16.76 


55.5 


3C 


273 


1 


Mar 


1974 


13.07 


1.03 


-13.3 






19 


Mar 


1974 


13.06 


1.66 


-73.2 






19 


Apr 


1974 


13.03 


4.96 


-55.9 






24 


Apr 


1974 


13.03 


2.91 


-35.8 






21 


May 


1974 


13.00 


0.92 


45.0 






13 


June 


1974 


13.02 


2.35 


-84.4 



Error Due to Short-Term Variability 

There has been much discussion about strong, short-term 
optical variability in BL Lac. Weistrop (1973) reported 
changes of 1.5 in one-hour periods. Bertaud e_t al_. (1969) 



107 

observed changes of more than 0.5 within periods of a few 
hours, although these events are based on single plates. 

Extensive photographic observations of BL Lac have been 
conducted at Rosemary Hill using both the 18-inch and 30-inch 
telescopes to determine the nature of possible short-term 
variations (Scott e_t al_. 1973b). The source has been observed 
on six nights since October, 1973, with a typical time resolu- 
tion of seven minutes. These observations indicate that 
statistically significant changes in brightness have occurred 
within some of these nights (Scott 1974). However, at the 
present time there has been no confirmation of the extremely 
violent behavior seen by Weistrop. Possible changes of 0.4 
have been observed, using the 30-inch reflector, within 
fifteen-minute periods but such behavior is the exception 
rather than the norm. Some doubt about the validity of these 
changes is indicated since simultaneous observations of BL 
Lac with the 18-inch and 30-inch telescopes on three nights 
fail to show simultaneous changes of this magnitude in both 
sets of data. 

The observations using the 18-inch Ritchey-Chretien 
telescope, which have approximately the same twenty-minute 
time resolution as the polarization photographs, show changes 
of 0™15 or less between exposures. A smoothing of the possi- 
ble very short-term behavior suggested by the 30-inch data by 
integrating over longer periods is thus indicated. In addi- 
tion, systematic changes in brightness of 0.3 or more occurring 



108 

over a night's observations (Hackney 1973) will not signifi- 
cantly affect a one-hour polarization run. 

In summation, there seems to be strong indication that 
BL Lac does exhibit intraday variability at least occasion- 
ally. Since the changes noted using the 18-inch are of the 
order of the polarimetric observational error, it is the 
author's belief that the present polarimetric observations 
are valid, although some uncertainty in the exact numerical 
results is indicated. 

3C 120 

The Seyfert galaxy 3C 120 has been shown to have linearly 
polarized radio emissions. Hobbs e_t al_. (1968) reported the 
polarized component at 2.07 cm to be variable, showing a 
decrease in polarized flux from four to one percent from 1965 
to 1967, while the position angle changed by forty degrees. 
Hobbs and Waak (19 72) also found the source to be 4.6 percent 
linearly polarized at 9.55 mm. Kinman (1967) observed no evi- 
dence of a linearly polarized optical component in 3C 120 
based on one night's observations, but he did suggest further 
observations in the ultraviolet. 

The present observations indicate that the optical 
linearly polarized component of 3C 120, while smaller than 
that of BL Lac, is significant. These observations consist 
of measurements made on eight nights during the latter part 
of 19 73 and early 19 74. The results are summarized in Table 5 



109 

and are shown graphically in Figure 20. While no polarization 
significantly greater than the observational error was ob- 
served on five nights, including November 19-20, when a dark- 
room accident caused a very large error, this was not the 
case on the remaining nights. (The results of the night of 
March 18 sho\tf a large polarized component but are made less 
convincing by the large observational error.) It is interest- 
ing to note that the position angle on October 27 is approxi- 
mately 90 degrees from the value of 45 degrees measured during 
the three observations immediately following. There is a 
suggestion of a similar smaller change in position angle 
corresponding to the observation of February 9. It should be 
noted that the large errors in position angle indicated for 
those nights where a very small polarized component was 
observed are a result of the method of observation and data 
reduction employed (see Chapter II). 

The optical flux was measured in the Johnson B magnitude 
system and is also shown in Figure 20. The light curve indi- 
cates that 3C 120 was quiescent from September, 1973, to 
January, 19 74. A gradual increase in brightness appears to 
have begun in February, 1974. The measurement of February 9 
was made during a 0™3 flare recorded by one plate but there 
is no indication of any burst on October 27. 

In conclusion, 3C 120 appears to have a linearly polar- 
ized optical component. This component is variable and com- 
pletely disappears at times. The increase in the polarized 
component during the February 9th flare is again suggestive 



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112 

of the brightening of a small region of the source that is 
strongly polarized. After the flare the total flux decreased 
0.4 by February 13, while the polarized flux decreased by 12 
percent, suggesting that activity in this same region had 
declined. 

3C 273 

The QSO 3C 273B is weakly polarized in the radio region 
of the spectrum. The polarized flux seems to show a peak at 
6 cm (Hobbs and Waak 1972). Appenzeller and Hiltner (1967) 
found no linear polarization significantly greater than the 
observational error at optical wavelengths. 

The present observations are summarized in Table 5 and 
are shown graphically in Figure 21. The brightness in the 
Johnson B magnitude system is also included. This curve 
indicates that the total flux varied by 0™5 in 1974 with two 
sharp "anti-flares" recorded by single plates accounting for 
most of the change. At no time during this period was the 
polarized component significantly greater than the observa- 
tional error. The position angle 6 showed large changes but 
has a large degree of uncertainty on all dates for the same 
reason discussed in the section on 3C 120. Thus, one may 
draw the conclusion that the present data suggest no evidence 
for a linearly polarized optical component of 3C 273B. 



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CHAPTER V 
SUMMARY AND CONCLUSIONS 

As a superficial review of the data presented in Chapters 
III and IV will indicate, QSO's exhibit a wide range of differ 
ing characteristics: from optically violent variables to 
apparently non-variable sources and from highly polarized 
objects to those with little or no linear polarization. It 
is difficult to construct a model for a class of objects which 
show such differing behavior among themselves and which resist 
all efforts to define "typical" QSO characteristics. The most 
plausible explanation of this wide range of characteristics 
is that one is observing objects with differing intrinsic 
properties, including mass, magnetic field strength, age, and 
undoubtedly other factors as well. A familiar analogy might 
be provided by galactic stars, which also exhibit a wide 
range of properties. Careful study of the data, however, dis- 
closes interesting behavior and similarities among certain 
objects which may lead to a more complete understanding of 
QSO's and related objects. 

Of the 38 calibrated objects compiled in Table 1, nine- 
teen show statistically significant optical variability. 
These sources, which comprise 50 percent of the sample, each 
had a confidence level for variability of 80 percent or more. 

115 



116 



In addition, twelve (39 percent) of the sources with tertiary- 
calibrations listed in Table 2 show a total range of O^ or 
more. This value is taken as good indication that a later 
calibration and statistical test will prove these objects to 
be variable. Finally, 44 percent of the objects with no cali- 
bration listed in Table 3 show some type of variability by 
visual inspection. 

Of the calibrated objects , optical variability was 
deduced for the first time in PKS 0202-17, PKS 0222-23, PKS 
1252+11, and NRAO 140. Of these, PKS 0222-23 is the most 
interesting in terms of the type and magnitude of variability 
observed and should be the subject of more intensive study in 
the future. Variations of a magnitude or more were observed 
in the sources OB 338, OX 074, PKS 0222-23, PKS 0518+16, PKS 
2345-16, NRAO 512, and 3C 446 and the peculiar galaxies OQ 
20 8 and 3C 371. These objects have been and will continue to 
be observed frequently so that knowledge of the nature of 
their variability can be further refined. 

One valuable contribution of a monitoring program of the 
sort described in this dissertation is a description of the 
type or form of variability observed in QSO's, especially if 
this form is repeated in several objects. Several character- 
istic types of behavior have been observed at Rosemary Hill. 
The first was a slowly varying component which appeared as a 
gradual rise and fall in the luminosity of an object, usually 
occurring within a period of several months. The range of the 
magnitude change observed was usually about 0.4 to 1.0. This 



117 

was the most common type of variability observed with eight 
of the objects discussed in Chapter III, OK 290, PKS 0202-17, 
PKS 0222-23, PKS 0336-01, PKS 0458-02, PKS 1252+11, PKS 
1347+21, and PKS 1607+26, showing this type of variability 
exclusively. In addition, many objects which were not dis- 
cussed in Chapter III because their probability of variation 
did not meet the 80 percent criterion also appeared to exhibit 
this slowly varying behavior. 

Violent optical variability, changes of 0.5 or more in 
the brightness of a source, is a widely discussed character- 
istic of several QSO's and related objects. This type of 
variability was observed at Rosemary Hill in the objects OB 
338, 01 318, OQ 208, OX 074, PKS 0518+16, PKS 2209+08, PKS 
2345-16, NRAO 512, 3C 371, and 3C 446. The most active were 
OX 074, PKS 2345-16, NRAO 512, and 3C 446, which showed 
changes of a magnitude or more within a week. It should be 
noted that all of these objects except OB 338, 01 318, and 
PKS 2209+08 also exhibit a slowly varying component in addi- 
tion to the violent changes. Perhaps further and more fre- 
quent observations will show that these remaining three 
objects also conform to this pattern. 

One of the most interesting results of the present obser- 
vations is the conclusion that major flares tend to be 
followed by a succession of minor events of decreasing ampli- 
tude. Hackney (1972) has already noted that NRAO 512 and OX 
074 exhibit this property. In 1970 NRAO 512 underwent a 1?0 
flare followed within five months by three well-defined flares 



118 

of 0.3 or more which were superimposed on a 2™0 decline in 
optical brightness. A 1?5 flare which rose smoothly to a 
maximum on 22 August 19 71 was followed by a 1?0 decline and 
than a sharp 0™8 flare within two months. This pattern was 
repeated in 1972 with a 0?75 flare on 18 April being followed 
by a slightly less energetic flare on 4 June. A 0?6 event in 
November was observed just before the end of the 1972 appari- 
tion. In 19 73 the object subsided to a minimum comparable 
to that of early 1971. Study of the light curve shown in 
Figure 18 suggests the possibility that the event of August, 
1971, was the major event while the four flares following it 
were secondary peaks. The decline in brightness to minimum 
was much slower in this case, taking eighteen months, compared 
with five months to reach approximately the same minimum after 
the 19 70 event. 

The QSO PKS 2345-16 was observed to brighten by A in 
December, 1969, followed by a 1™0 decline and a subsequent 
rapid return to its December brightness in January, 1970. In 
July, during the next apparition, a 0?5 flare was observed 
which appeared to be superimposed on a smooth, slow decline 
that apparently took place during the period January to 
November, 19 70. 

The peculiar source OX 074 also exhibits this "ringing" 
behavior. A 0?7 flare in 1970 was followed within six weeks 
by a 0™3 event. In 19 73 a flare 0^8 above minimum was 
followed by an event of 0?7 three months later. The multiple 
secondary events of the previous two QSO's do not seem to be 



119 

present in this object, perhaps related to the fact that the 
initial flares were less intense than those of NRAO 512 and 
PKS 2345-16. 

The Seyfert galaxy OQ 208 also showed indications of the 
same behavior. A 0?8 flare observed in 19 74 was followed by 
three smaller events within five months. All events were 
recorded by only one plate, however. 

An indication of the same behavior on a much longer time 
scale is given by PKS 0222-23. The object reached its great- 
est brightness in January 19 70. The next event occurred in 
October, 1972, with an indication of another in 1974. Each 
was 0.5 fainter than the previous maximum. 

The QSO PKS 2209+08 showed a well-defined 0™3 flare 
following by two months a 0™55 event in 19 70. There was no 
indication of a declining minimum. 

All of these objects will continue to be observed to 
search for this type of behavior. In addition, the QSO 3C 
446, which flared by 2™5 in 1974, will be observed intensively 
to discover whether this type of behavior is present. 

It could be argued that the events seen following a major 
flare were coincidental. However, NRAO 512 showed only one 
event of the magnitude of the secondary events described 
above which did not appear during the decline in brightness 
following a major flare. The QSO PKS 2345-16 showed two 
such events and OX 74 only one recorded by one plate. The 
time span is not long enough to draw conclusions from PKS 
0222-23. Thus, there is a definite indication that the minor 



120 



flares are related to a major event, suggesting some sort of 
explosive or pulsational model. 

The results of the linear polarization observations indi- 
cate that BL Lacertae, 3C 120, and 3C 273 exhibit different 
polarized components , due perhaps to dissimilar radiation 
mechanisms. It is instructive to compare the linear polari- 
zation activity with the optical activity of each of these 
objects. The Lacertid BL Lac has an intense, rapidly varying 
linearly polarized component and is a violent optical vari- 
able (OW) (Hackney 19 73; Visvanathan 19 73a). The Seyfert 
galaxy 3C 120 exhibited a smaller polarized component than 
BL Lac, the magnitude often being no larger than the observa- 
tional error. It is also classified as OVV (Hackney 1973) 
but the Rosemary Hill light curve indicates that the object 
does not vary as rapidly and violently as BL Lac. The QSO 
3C 273 showed no significant linear polarization. Recent 
observations by the Florida group and others (Penston and 
Cannon 1970; Tritton and Selmes 1971) indicate that the total 
change in brightness in recent years was less than 0.5, a much 
smaller value than was observed in BL Lac or 3C 120. Thus, 
the observations indicate a definite but unsurprising trend 
toward an increase in strength and variability of the linearly 
polarized component with an increase in the optical activity 
of a source. 

The observations also indicate that Visvanathan' s (1973a) 
hypothesis that the observed polarized radiation is the resul- 
tant of emission from several active spots cannot be 



121 

contradicted. The large changes in position angle of polari- 
zation coincident with a change in the total flux of BL Lac, 
and the increase in polarized flux observed when the total 
flux of 3C 120 increased in February and March, 19 74, lend 
some support to this theory. 

To summarize, a wide range of optical behavior is 
observed in QSO's and related objects. The theory that these 
objects are composed of several smaller optically variable 
sources seems the most satisfying for two reasons. First, it 
could account for the short rise and decay time of flares 
observed in objects such as NRAO 512 and 3C 446 without 
requiring that the total energy be emitted from regions of 
the small size implied by the rapid variations. Secondly, 
if these sources are highly polarized, presumably by either 
the synchrotron or the inverse Compton process, the resultant 
emission of these variable sources could account for the vari- 
able linearly polarized components of BL Lac and 3C 120. The 
number of such sources composing a QSO would be about three 
to five, so that the flaring of one could cause a large in- 
crease in the observed total flux. In slowly varying sources 
the individual sources are relatively inactive and, if 3C 273 
is a good example, relatively unpolarized. The absence of 
any confirmed periodic behavior in QSO's also supports this 
model. 

This hypothesis leaves unanswered the question of the 
mechanism of the individual sources. The supernova theory 
of Colgate (1969) could account for the observed behavior of 



122 



the active objects but is not satisfactory for slowly varying 
or non-varying objects or inactive phases of OW objects, as 
a source of the observed radiated energy is not present if 
there are no supernova events. One can assume the lack of 
sharp changes in brightness is due to a very large number of 
supernova events but this implies that the slowly varying 
objects should tend to be brighter than active objects, an 
observationally incorrect statement (Ames and Bahcali 1973). 
The cause of the observed ringing behavior is also not 
explained by this model. It is difficult to understand how 
the spinar model could be reconciled to the individual source 
hypothesis. In addition, the lack of any observed periodici- 
ties rule out this model. The magnetic rotator model is 
appealing because it could explain the observed radio behavior 
and leaves open to question the source of optical activity in 
the center of the object. There is as yet no widely accepted 
evidence that QSO's are closer than cosmological distances, 
so only cosmological models have been discussed. 

In conclusion, the final explanation of the nature of 
the emission mechanism or mechanisms of QSO's will have to 
await future discoveries and observations. It is the author's 
hope that this dissertation will aid in that explanation. 



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BIOGRAPHICAL SKETCH 

Ben Q. McGimsey, Jr., was born on September 14, 1946, 
in Birmingham, Alabama, the eldest son of Ben and Carolyn 
McGimsey. He was graduated from Shades Valley High School, 
Birmingham, in June, 1964. As an undergraduate at Birmingham- 
Southern College he was a member of Theta Chi Fraternity, 
Theta Sigma Lambda Honorary Mathematics Fraternity, and Phi 
Beta Kappa. He was graduated cum laude in June, 196 8, with 
the degree of Bachelor of Science. He entered the Graduate 
School of the University of Florida in pursuit of the degree 
of Doctor of Philosophy with a major in physics, expected in 
December, 19 74. He became a member of Sigma Xi and the 
American Astronomical Society in 1974. He was married to 
the former Karen Sundback in 1969. 



128 



I certify that I have read this study and that in my 
opinion it conforms to acceptable standards of scholarly 
presentation and is fully adequate, in scope and quality, 
as a dissertation for the degree of Doctor of Philosophy. 




7C. G~. Smith, Chairman 
Professor of Physics and 
Astronomy 



I certify that 1 have read this study and that in my 
opinion it conforms to acceptable standards of scholarly 
presentation and is fully adequate, in scope and quality, 
as a dissertation for the degree of Doctor of Philosophy. 



•5" 




A^ E. S~~! Green 
Graduate Research 
of Physics 



Professor 



I certify that I have read this study and that in my 
opinion it conforms to acceptable standards of scholarly 
presentation and is fully adequate, in scope and quality, 
as a dissertation for the degree of Doctor of Philosophy. 



5 W 



F. B. Wood 
Professor o 



Astronomy 



I certify that I have read this study and that in my 
opinion it conforms to acceptable standards of scholarly 
presentation and is fully adequate, in scope and quality, 
as a dissertation for the degree of Doctor of Philosophy. 



R. C. IsTcF 



Associate Professor of Physics 



I certify that I have read this study and that in my 
opinion it conforms to acceptable standards of scholarly 
presentation and is fully adequate, in scope and quality, 
as a dissertation for the degree of Doctor of Philosophy. 

S^ T. Gottesman 

Assistant Professor of Astronomy 

This dissertation was submitted to the Graduate Faculty of 
the Department of Physics and Astronomy in the College of 
Arts and Sciences and to the Graduate Council, and was 
accepted as partial fulfillment of the requirements for i"he 
degree of Doctor of Philosophy. 

December, 1974 



Dean, Graduate School 



*« x 



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