THE MOJAVE CHANDRA SAMPLE: A CORRELATION STUDY OF BLAZARS
AND RADIO GALAXIES IN X-RAY AND RADIO WAVELENGTHS
A Dissertation
Submitted to the Faculty
of
Purdue University
by
Brandon S. Hogan
In Partial Fulfillment of the
Requirements for the Degree
of
Doctor of Philosophy
May 2011
Purdue University
West Lafayette, Indiana
ii
[I dedicate this to my lovely wife, Meredith, my wonderful group of friends, and my
supportive family. I could not have done this without the love and support of all of
you.]
iii
ACKNOWLEDGMENTS
[I would like to acknowledge Matthew Lister, Herman Marshall, Nathan Cooper,
Preeti Kharb, and Talvikki Hovatta, as they have supported and helped me throughout the duration of this project. This project was funded by grants from NASA and
NSF.]
iv
TABLE OF CONTENTS
Page
LIST OF TABLES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
vii
LIST OF FIGURES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
viii
SYMBOLS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
xi
ABBREVIATIONS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
xii
ABSTRACT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
xiv
1 INTRODUCTION . . . . . . . . . . . . . . . . . . .
1.1 Active Galactic Nuclei . . . . . . . . . . . . . .
1.1.1 Radio Quiet AGN . . . . . . . . . . . . .
1.1.2 Radio Loud AGN . . . . . . . . . . . . .
1.2 The Fanaroff Riley Classification of AGN . . . .
1.3 Relativistic Properties of AGN . . . . . . . . . .
1.3.1 Apparent Superluminal Motion . . . . .
1.3.2 Beaming . . . . . . . . . . . . . . . . . .
1.3.3 Inverse-Compton Scattering . . . . . . .
1.4 Astronomical Instruments used in the MOJAVE
1.4.1 The Very Large Array . . . . . . . . . .
1.4.2 The Very Long Baseline Array . . . . . .
1.4.3 Chandra X-ray Observatory . . . . . . .
1.5 The Status of X-ray Jet Astrophysics . . . . . .
1.6 Thesis Description and Outline . . . . . . . . .
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Chandra
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2 THE MOJAVE CHANDRA SAMPLE . .
2.1 Selection Criteria . . . . . . . . . . .
2.2 Individual Source Observations of the
2.2.1 0106+013(OC 12) . . . . . . .
2.2.2 0119+115 . . . . . . . . . . .
2.2.3 0224+671 (4C 67.05) . . . . .
2.2.4 0234+285 (CTD 20) . . . . .
2.2.5 0415+379 (3C 111) . . . . . .
2.2.6 0529+075 (OG 050) . . . . .
2.2.7 0605-085 . . . . . . . . . . . .
2.2.8 1045-188 . . . . . . . . . . . .
2.2.9 1055+018 (4C 01.28) . . . . .
2.2.10 1156+295 (4C 29.45) . . . . .
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MCS .
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3 DATA REDUCTION AND ANALYSIS . . . . . . . . . . . . . . .
3.1 X-ray Radio Overlays . . . . . . . . . . . . . . . . . . . . . .
3.2 X-ray and Radio Jet Analysis . . . . . . . . . . . . . . . . .
3.3 The Single Zone IC/CMB Model . . . . . . . . . . . . . . .
3.4 Scenarios Associated with the IC-CMB Model . . . . . . . .
3.4.1 IC/CMB model with No Jet Deceleration or Bending
3.4.2 IC/CMB model with Jet Deceleration . . . . . . . . .
3.4.3 IC/CMB model with Deceleration and Jet Bending .
3.5 Sext as an X-ray jet predictor . . . . . . . . . . . . . . . . .
3.6 Kolmogorov-Smirnov Tests . . . . . . . . . . . . . . . . . . .
3.7 Viewing Angle . . . . . . . . . . . . . . . . . . . . . . . . . .
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4 SPECTRAL ENERGY DISTRIBUTIONS . .
4.1 General Information . . . . . . . . . . .
4.2 3C 111 (0415+379) SED . . . . . . . . .
4.2.1 Obtaining the Radio Fluxes . . .
4.2.2 Obtaining the Optical Fluxes . .
4.2.3 Obtaining the X-ray Fluxes . . .
4.2.4 Obtaining the γ-ray Fluxes . . .
4.2.5 Uniqueness of the 3C 111 Hotspot
4.3 Individual SED Notes . . . . . . . . . . .
4.4 Summary . . . . . . . . . . . . . . . . .
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5 SUMMARY . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1 Goals and Results . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1.1 X-ray Detection Rate . . . . . . . . . . . . . . . . . . . . . .
84
84
84
2.2.11
2.2.12
2.2.13
2.2.14
2.2.15
2.2.16
2.2.17
2.2.18
2.2.19
2.2.20
2.2.21
2.2.22
2.2.23
2.2.24
2.2.25
2.2.26
2.2.27
1222+216 (4C 21.35) .
1226+023 (3C 273) . .
1253-055 (3C 279) . .
1334-127 . . . . . . . .
1510-089 . . . . . . . .
1641+399 (3C 345) . .
1655+077 . . . . . . .
1800+440 (S4 1800-44)
1828+487 (3C 380) . .
1849+670 (S4 1849-67)
1928+738 (4C 73.18) .
1957+405 (Cygnus A)
2155-152 . . . . . . . .
2201+315 (4C 31.63) .
2216-038 . . . . . . . .
2251+158 (3C 454.3) .
2345-167 . . . . . . . .
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SED
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LIST OF REFERENCES . . . . . . . . . . . . . . . . . . . . . . . . . . . .
89
5.2
5.1.2 IC/CMB Model . . . . . . . . . . . . . . . .
5.1.3 Misalignment Angles . . . . . . . . . . . . .
5.1.4 Spectral Energy Distributions . . . . . . . .
Future Work . . . . . . . . . . . . . . . . . . . . . .
5.2.1 Expanding the MCS . . . . . . . . . . . . .
5.2.2 Deeper X-ray Observations of MCS Sources
5.2.3 Optical Observations MCS Sources . . . . .
APPENDICES . . . . . . . . . . . . . . . . . . . . . . .
Appendix A: Radio Profiles . . . . . . . . . . . . . .
Appendix B: X-ray Profiles . . . . . . . . . . . . . .
Appendix C: Bulk Lorentz Factor vs. Viewing Angle
Appendix D: Spectral Energy Distributions . . . . .
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92
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108
VITA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
115
vii
LIST OF TABLES
Table
Page
2.1
MOJAVE CHANDRA SAMPLE . . . . . . . . . . . . . . . . . . . . .
21
2.2
OBSERVATION LOG . . . . . . . . . . . . . . . . . . . . . . . . . . .
22
2.3
MOJAVE CHANDRA SAMPLE JET MEASUREMENTS . . . . . . .
23
2.4
VLA ARCHIVAL DATA . . . . . . . . . . . . . . . . . . . . . . . . . .
25
3.1
MOJAVE CHANDRA SAMPLE BEAMING MODEL PARAMETERS
57
4.1
SED PARAMETERS . . . . . . . . . . . . . . . . . . . . . . . . . . . .
74
4.2
3C 111 SED INFORMATION . . . . . . . . . . . . . . . . . . . . . . .
79
viii
LIST OF FIGURES
Figure
Page
1.1
Visual representation of a radio-loud AGN [Urry & Padovani, 1995] . .
2
1.2
AGN Taxonomy [Urry & Padovani, 1995] . . . . . . . . . . . . . . . . .
4
1.3
The FR I/FR II Divide [Ghisellini et al., 1993] . . . . . . . . . . . . . .
7
1.4
Visual representation of superluminal motion as seen in Ghisellini [2000].
9
1.5
Relativistic beaming of radiation which is emitted isotropically in the rest
frame K′ (S′ in the text) [Rybicki & Lightman, 1979]. . . . . . . . . . .
12
1.6
Inverse-Compton Scattering . . . . . . . . . . . . . . . . . . . . . . . .
13
1.7
Chandra X-ray Observatory (Courtesy of NASA/CXC/NGST) . . . . .
15
2.1
Radio/X-ray overlay of 0106+013. . . . . . . . . . . . . . . . . . . . . .
24
2.2
Radio/X-ray overlay of 0119+115. . . . . . . . . . . . . . . . . . . . . .
27
2.3
Radio/X-ray overlay of 0224+671. . . . . . . . . . . . . . . . . . . . . .
28
2.4
Radio/X-ray overlay of 0234+285. . . . . . . . . . . . . . . . . . . . . .
29
2.5
Radio/X-ray overlay of 0415+379. . . . . . . . . . . . . . . . . . . . . .
30
2.6
Radio/X-ray overlay of 0529+075. . . . . . . . . . . . . . . . . . . . . .
31
2.7
Radio/X-ray overlay of 0605-085. . . . . . . . . . . . . . . . . . . . . .
32
2.8
Radio/X-ray overlay of 1045-188. . . . . . . . . . . . . . . . . . . . . .
33
2.9
Radio/X-ray overlay of 1055+018. . . . . . . . . . . . . . . . . . . . . .
34
2.10 Radio/X-ray overlay of 1156+295. . . . . . . . . . . . . . . . . . . . . .
35
2.11 Radio/X-ray overlay of 1222+216. . . . . . . . . . . . . . . . . . . . . .
36
2.12 Radio/X-ray overlay of 1226+023. . . . . . . . . . . . . . . . . . . . . .
37
2.13 Radio/X-ray overlay of 1253-055. . . . . . . . . . . . . . . . . . . . . .
38
2.14 Radio/X-ray overlay of 1334-127. . . . . . . . . . . . . . . . . . . . . .
39
2.15 Radio/X-ray overlay of 1510-089. . . . . . . . . . . . . . . . . . . . . .
40
2.16 Radio/X-ray overlay of 1641+399. . . . . . . . . . . . . . . . . . . . . .
41
ix
Figure
Page
2.17 Radio/X-ray overlay of 1655+077. . . . . . . . . . . . . . . . . . . . . .
42
2.18 Radio/X-ray overlay of 1800+440. . . . . . . . . . . . . . . . . . . . . .
43
2.19 Radio/X-ray overlay of 1828+487. . . . . . . . . . . . . . . . . . . . . .
44
2.20 Radio/X-ray overlay of 1849+670. . . . . . . . . . . . . . . . . . . . . .
45
2.21 Radio/X-ray overlay of 1928+738. . . . . . . . . . . . . . . . . . . . . .
46
2.22 Radio/X-ray overlay of 1957+405. . . . . . . . . . . . . . . . . . . . . .
47
2.23 Radio/X-ray overlay of 2155-152. . . . . . . . . . . . . . . . . . . . . .
48
2.24 Radio/X-ray overlay of 2201+315. . . . . . . . . . . . . . . . . . . . . .
49
2.25 Radio/X-ray overlay of 2216-038. . . . . . . . . . . . . . . . . . . . . .
50
2.26 Radio/X-ray overlay of 2251+158. . . . . . . . . . . . . . . . . . . . . .
50
2.27 Radio/X-ray overlay of 2345-167. . . . . . . . . . . . . . . . . . . . . .
51
3.1
Histogram relating the source population to the Sext value . . . . . . .
67
3.2
Histogram representing the redshift distribution of the MOJAVE sample
69
3.3
Histogram representing the redshift distribution of the MCS sample . .
69
3.4
Position Angle Misalignment Associated with the MCS . . . . . . . . .
71
4.1
Spectral Energy Distribution for the hotspot associated with the primary
jet in 3C 111 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
74
Radio, optical, and X-ray jets associated with 3C 273 [Jester et al., 2006]
81
4.2
A.1 Radio Profiles . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
93
A.2 Radio Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .
94
A.3 Radio Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .
95
A.4 Radio Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .
96
A.5 Radio Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .
97
B.1 X-ray Profiles . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
98
B.2 X-ray Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .
99
B.3 X-ray Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .
100
B.4 X-ray Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .
101
B.5 X-ray Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .
102
x
Figure
Page
C.1 Bulk Lorentz Factor vs. Viewing Angle . . . . . . . . . . . . . . . . . .
104
C.2 Bulk Lorentz Factor vs. Viewing Angle Cont. . . . . . . . . . . . . . .
105
C.3 Bulk Lorentz Factor vs. Viewing Angle Cont. . . . . . . . . . . . . . .
106
C.4 Bulk Lorentz Factor vs. Viewing Angle Cont. . . . . . . . . . . . . . .
107
D.1 Spectral Energy Distribution for the knot associated with the primary jet
in 1641+399 [Sambruna et al., 2004] . . . . . . . . . . . . . . . . . . .
108
D.2 Spectral Energy Distribution for the knots associated with the primary jet
in 3C 273 (1226+023) [Jester et al., 2006] . . . . . . . . . . . . . . . .
109
D.3 Spectral Energy Distribution for the knots associated with the primary jet
in 3C 273 (1226+023) [Sambruna et al., 2001] . . . . . . . . . . . . . .
110
D.4 Spectral Energy Distribution for the knots associated with the primary jet
in 3C 273 (1226+023) [Marshall et al., 2001] . . . . . . . . . . . . . . .
111
D.5 Spectral Energy Distribution for the knot associated with the primary jet
in 1222+216 [Jorstad & Marscher, 2006] . . . . . . . . . . . . . . . . .
112
D.6 Spectral Energy Distribution for the primary jet in 3C 279 (1253-055)
[Collmar et al., 2010] . . . . . . . . . . . . . . . . . . . . . . . . . . . .
113
D.7 Spectral Energy Distribution for the primary jet in 1928+738 [Sambruna
et al., 2004] . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
114
xi
SYMBOLS
Γ
Bulk Lorentz Factor
δ
Doppler Factor
βapp
Apparent Superluminal Speed
vapp
Apparent Superluminal Velocity
c
Speed of Light
φ
Filling Factor
k
Baryon Fraction Energy Parameter
V
Emitting Volume
L
Observed Synchrotron Luminosity
C
Weak Function of the Low Frequency Spectral Index of the Synchrotron Spectrum
B1
Non-Boosted Spatially Averaged, Minimum Energy Magnetic
Field of the Jet
R
X-ray to Radio Luminosity Ratio
νr
Radio Frequency
νx
X-ray Frequency
Sr
Radio Flux Density
Sx
X-ray Flux Density
K
K Factor
θ
Angle With Respect to the Line of Sight
µ
Cosine of θ
xii
ABBREVIATIONS
AGN
Active Galactic Nuclei
VLA
Very Large Array
VLBA
Very Large Baseline Array
VLBI
Very Large Baseline Interferometry
HST
Hubble Space Telescope
MOJAVE
Monitoring Of Jets in AGN with VLBA Experiments
MCS
MOJAVE Chandra Sample
FSRQ
Flat Spectrum Radio Quasar
BL Lac
BL Lacertae object
SSRQ
Steep Spectrum Radio Quasar
NLRG
Narrow Line Radio Galaxies
BLRG
Broad Line Radio Galaxies
NELG
Narrow Emission Line Galaxies
QSO
Quasi Stellar Object (Quasar)
BLR
Broad Line Region
NLR
Narrow Line Region
BAL
Broad Absorption Line (Quasar)
FR I
Fanaroff & Riley type I object
FR II
Fanaroff & Riley type II object
CDQ
Core Dominated Quasar
LDQ
Lobe Dominated Quasar
IGM
Inter-Galactic Medium
CMB
Cosmic Microwave Background
IC
Inverse Compton
NGST
Nortrop Grumman Space Technology
xiii
ISIM
Integrated Science Instrument Module
HRC
High Resolution Camera
CCD
Charge Collecting Device
ACIS
Advanced CCD Imaging Spectrometer
SAO
Smithsonian Astrophysical Observatory
MIT
Massachusetts Institute of Technology
IC/CMB
inverse Compton scattering off of cosmic microwave background
NASA
National Aeronautics and Space Administration
SED
Spectral Energy Distribution
FWHM
Full Width Half Maximum
WCS
World Coordinate System
xiv
ABSTRACT
Hogan, Brandon S. Ph.D., Purdue University, May 2011. The MOJAVE Chandra
Sample: A Correlation Study of Blazars and Radio Galaxies in X-ray and Radio
Wavelengths . Major Professor: Matthew L. Lister.
The Chandra X-ray observatory has increased the quality and number of detections in the X-ray regime since its launch in 1999. It is an imporant tool for studying
the jets which are associated with Active Galacitc Nuclei (AGN) and their possible
emission mechanisms. The MOJAVE Chandra Sample (MCS) is a sample of 27 AGN
which have been selected from the radio flux-limited MOJAVE (Monitoring of Jets
in AGN with VLBA Experiments) sample. The objects contained in the MOJAVE
sample are traditionally associated with relativistically beamed jets that have small
viewing angles. The MCS was created to study the correlation of X-ray and radio
emission on kiloparsec scales. The complete sample is made up of all MOJAVE Fanaroff & Riley type II objects which have over 100 mJy of extended radio emission
at 1.4 GHz and a radio structure of at least 3′′ in extent. Chandra observations
have revealed X-ray and radio correlation in 21 of the 27 jets, bringing the detection
rate to ∼78%. The selection criteria provides a quantitative method of discovering
new X-ray jets associated with AGN from radio observations. The X-ray morphologies are usually well correlated with the radio emission, except for the sources which
show extreme bending on the kiloparsec scale. The emission mechanism for these
relativisiticly beamed quasars and radio galaxies can be interpreted as inverse Compton scattering off of the consmic microwave background by the electrons in the jets
(IC/CMB). The emission mechanism is reinforced by spectral energy distributions
(SED) which model the emission mechanisms for sources with sufficient X-ray, optical, and radio data available. I have explored the effects of jet bending and jet
deceleration in conjunction with the inverse Compton emission model and used dif-
xv
ferent scenarios to derive best fit viewing angles and bulk Lorentz factors, which
were calculated by using the superluminal speeds along with parameters that were
derived from the IC/CMB model. The range of viewing angles and Lorentz factors
are examined for each scenario, as well as their implications for the other parameters
associated with models. To achieve results that are consistant with other models jet
bending and deceleration must be considered with the IC/CMB model.
1
1. INTRODUCTION
The overall goal of this dissertation is to investigate and further understand how jet
emission from Active Galactic Nuclei (AGN) correlates between the radio and X-ray
regimes. In this dissertation I describe how one can use the data associated with
the MOJAVE (Monitoring Of Jets in AGN with VLBA Experiments) sample along
with additional selection criteria to determine if X-ray jet detections are probable with
the Chandra X-ray Observatory. I also study the implications of this correlation on
the overall inverse-Compton emission mechanism which I have chosen for modeling
the X-ray emission of these objects. The introductory sections below provide the
background information necessary to form a foundation for the understanding of the
material provided in this dissertation.
1.1
Active Galactic Nuclei
An AGN is traditionally regarded as an accreting supermassive black hole, which
is located at the center of a galaxy, and has a mass on the order of 106 M⊙ to 1010
M⊙ , where M⊙ is one solar mass. This black hole is surrounded by an accretion
disk, which is made up of material that spirals inward toward the black hole, and
encompasses a flat circular region which is perpendicular to the rotation poles and/or
jets. As seen in Figure 1.1, a typical radio-loud AGN is comprised of a black hole,
an accretion disk, a torus, and a jet, where the narrow line regions are found further
away from the core than the broad line regions. The narrow line and broad line regions are where the narrow and broad emission lines are produced.
Jets are hypothesized to have been produced by a phenomenon known as magnetic
launching, which is described by the rotation of a black hole and accretion disk system
2
Figure 1.1. Visual representation of a radio-loud AGN [Urry & Padovani, 1995]
[Marscher, 2009]. This leads to the magnetic field lines twisting up into a helical
structure, which then causes a pressure gradient to occur, accelerating the plasma
flow downstream. This helical magnetic field is often thought of as the confinement
structure for the jet. Bridle & Perley [1984] define extragalactic radio jets by three
criteria.
• The jet must be four times as long as it is wide.
• It must be separable from other extended structures at high resolution.
3
• It must be aligned with the compact radio core that it protrudes from.
Extragalactic jets, which are comprised of highly energetic plasma, often appear to be
moving near the speed of light, or in some cases, even faster than the speed of light.
This is known as apparent superluminal motion and is described in further detail in
§1.3.1.
There are traditionally two ways that a jet can terminate. The jet either dissipates enough energy into the environment around it, such that it fades slowly in a
plume-like structure, or it abruptly terminates at a shock front known as a terminal
hotspot. Hotspots are often enveloped in large regions of radiation which gravitate
backwards toward the core, known as lobes. These lobes are often seen at the end of
both jets, even though the jet which is traveling away from the observer can sometimes not be seen due to relativistic beaming effects (see §1.3.2).
AGN are traditionally separated into two groups; radio-loud and radio-quiet AGN.
Radio-loud AGN make up about 15% to 20% of the total AGN population [Urry &
Padovani, 1995]. These groups are defined by their ratios of 5 GHz radio flux to
optical (B-band) flux. If the ratio of radio flux to optical flux is greater than 10
then the object is considered to be radio loud [Kellermann et al., 1989]. The radioloud classification typically includes the Narrow Line Radio Galaxies (NLRG), Broad
Line Radio Galaxies (BLRG), Steep Spectrum Radio Quasars (SSRQ), Flat Spectrum Radio Quasars (FSRQ), and Blazars (FSRQs and BL Lacertae objects), and
the radio-quiet classification includes Seyfert Galaxies (type 1 & 2), Narrow Emission
Line (X-ray) Galaxies (NELG), Infrared Quasars, and radio quiet Quasars (Quasi
Steller Objects, QSO). A visual description of the relation of these items to radio
loudness, angle to the line of sight, black hole spin, and optical emission lines is found
in Figure 1.2.
4
Figure 1.2. AGN Taxonomy [Urry & Padovani, 1995]
1.1.1
Radio Quiet AGN
Seyfert Galaxies
Seyfert galaxies have the lowest luminosities of all of the sources in the radio quiet
regime, and are usually located much nearer to us than the more powerful AGN.
They are sub-divided into two types based on the emission lines that they produce.
Type 1 Seyfert galaxies exhibit emission lines from the Broad Line Region (BLR)
and the Narrow Line Region (NLR), while their Type 2 counterparts exhibit emission
lines from only the NLR. The BLR emission lines are produced close to the core,
presumably from the interaction between the emission from the core/jet nozzle and
the clouds above the accretion disk, or perhaps by the disk itself [line width ≤10000
km/sec, Antonucci 1993]. Thus, orientation could cause the difference in the two
types of Seyfert galaxies, because at smaller angles the emission does not have to
travel through the dusty torus (i.e., The rotation axis of Type 1 Seyfert galaxies has
a smaller angle to the line of sight than the Type 2 Seyferts). The narrow emission
5
lines are produced from clouds which are downstream from the nucleus [line width
≤1000 km/sec, Antonucci 1993, Urry & Padovani 1995].
Radio-Quiet Quasars
A radio quiet quasar can show both broad and narrow absorption lines like its
Seyfert Type 1 counterpart, but is distinguished by its larger luminosity [Urry &
Padovani, 1995]. Broad Absorption Line Quasars (BALs) make up about 10% of the
population of radio quiet quasars. Interestingly, the dust clouds which are thought to
be the cause of BALs cover about 10% of the source, leading to the presumption that
all radio quiet quasars have these clouds around them [Antonucci 1993 and references
within]. The polar axis (or jet) viewing angle might account for the difference between
BAL quasars and the rest of the radio quiet quasars, as seen in Figure 1.2.
1.1.2
Radio Loud AGN
Blazars
Blazars, like most other radio-loud AGN, are described by very powerful jets
which are generated in AGN as a result of accretion onto supermassive black holes.
Blazar jets can transport energy over large distances using highly energetic plasma
as a medium. These energetic outflows are usually oriented at very small angles
with respect to the observer’s line of sight (θ ≤ 15◦ ) and tend to show apparent
superluminal velocities [Angel & Stockman, 1980]. The blazar class encompasses two
groups of objects, the flat spectrum radio quasars (FSRQs) and BL Lacertae (BL
Lac) objects. The FSRQ objects are thought to have more powerful, well collimated
jets that terminate at large shock fronts known as hotspots, and are also described
as Fanaroff-Riley type II jets (FRII; Fanaroff & Riley 1974). On the other hand,
the BL Lac objects are described as Fanaroff-Riley type I objects, are less powerful
than FR IIs, and tend to dissipate more energy into the intergalactic medium before
6
they terminate in a plume like structure [Urry & Padovani, 1995]. The FanaroffRiley classification scheme is discussed in greater detail in §1.2. In terms of the X-ray
production in the jet, the inverse Compton radiation process is suggested to be more
important in FSRQs than in the less powerful BL Lac sources, even though both have
small angles to the line of sight [Ghisellini & Tavecchio, 2008, Harris & Krawczynski,
2006]. The jets in blazars are often one-sided because of relativistic beaming, which
will be described in §1.3.2.
Radio Galaxies and Radio Quasars
One large difference between the radio galaxies and quasars and blazars is that
the blazars are viewed at a very small angle to the line of sight when compared to
the galaxies and quasars. Blazars, galaxies, and quasars have the same central engine
structure at the core and possibly the same jet structure. Because radio galaxies are
often viewed at larger angles than blazars, they are often described by their symmetric
radio lobes, as opposed to the jets that are seen in the blazar class. There are some
special cases of radio galaxies (example: Cygnus A, M87, and 3C111) that show a
well collimated jet along with the radio lobes in the radio regime, as well as correlated
X-ray jet emission (Wilson et al. 2001, Marshall et al. 2002, Hogan et al. 2011).
1.2
The Fanaroff Riley Classification of AGN
Fanaroff & Riley [1974] discovered a relationship between the the location of the
brightest portions of a radio jet and its radio luminosity. 199 sources from the 3CR
complete sample [Mackay, 1971] were studied and divided into two distinct classes
(Fanaroff & Riley Type I and Type II) which were defined by the ratio of the distance
between the brightest regions on opposite sides of the central AGN to the total extent
of the source. Any source with a value of 0.5 or less for the previous quantity was
classified as a Fanaroff & Riley Type I (or FR I) galaxy, and any source with a value
greater than 0.5 was classified as a FR II source. FR I sources tend to have the
7
brightest jet regions located closer to the core of the quasar or radio galaxy, whereas
FR II sources have the bright hot spots located further away from the core. The
FR I/FR II divide is further reinforced by a division in luminosity between the two
classes. The sources are separated by a threshold luminosity value of 2 × 1025 W Hz−1
sr−1 at 178 MHz, with the FR I class having a luminosity lower than this threshold
and FR II class having a value above it. In the optical regime the FRI sources are
more luminous than the FR II sources when viewed at the same radio luminosity
[Owen & Ledlow, 1994], which implies that the FR I/FR II break depends on the
optical as well as the radio luminosity.
Bicknell [1985] suggests that the differences in the FR I and FR II classes is from
the confinement by the pressure of the hot surrounding medium. FR I sources are
thought to be dominated by turbulence and entrainment which can slow the jet down
gradually without the need for a shock front. The more powerful FR II sources are
not in pressure equilibrium and thus are susceptible to shocks produced by Kelvin
Helmholtz instabilities within the jet.
Figure 1.3. The FR I/FR II Divide [Ghisellini et al., 1993]
FR I and FR II quasar jets should have angles to the line of sight which are not
greater than ∼ 40◦ [Ghisellini et al., 1993]. The FR I/FR II division is described below
8
as well as visually in Figure 1.3. FR II quasars can be described as lobe dominated
quasars (LDQ) and core dominated quasars (CDQ), where the CDQ usually have a
viewing angle ≤ 10◦ and are associated with the FSRQ class described earlier. The
LDQ are often associated with the SSRQ class and have viewing angles which range
from 10◦ to 40◦ . The FR I class of objects is often divided into X-ray selected BL
Lac objects and radio selected BL Lac objects. The X-ray selected BL Lac object is
now part of the classification of high synchrotron peaked blazars or HSP [Abdo et al.,
2010]. This more recent classification describes BL Lac objects that have their X-ray
emission mechanism characterized by a synchrotron spectrum instead of an inverse
Compton spectrum (see §4.1). One interpretation suggests that the radio selected BL
Lac objects are viewed at angles ≤ 15◦ while the HSP objects are viewed at angles
between 15◦ and 30◦ [Ghisellini et al., 1993].
1.3
Relativistic Properties of AGN
1.3.1
Apparent Superluminal Motion
Supermassive black holes can transport energy through massive jets which protrude from the core perpendicular to the plane of the accretion disk. This energy
is transported through the bulk motion of plasma moving at a relativistic velocity
[Rees, 1966]. If the plasma is moving at speeds very close to the speed of light and is
moving toward the observer at a very small angle, it can seem to move faster than the
speed of light. This is called apparent superluminal motion (βapp ). A mathematical
description of this phenomenon is described below [Ghisellini, 2000].
First we assume that there is an object located at point A which emits photons.
This object then moves to location B in a time interval (measured by the observer) of
∆te where it emits another photon. The second assumption is that the object has an
actual velocity which is close to the speed of light and that the velocity vector’s angle
to the line of sight (θ) is small. A visual representation of this is shown in Figure 1.4.
9
Figure 1.4. Visual representation of superluminal motion as seen in Ghisellini [2000].
The distance between points A and B is equal to βc∆te where β is just the velocity
of the object divided by the speed of light (c) and is described by Equation 1.1 below.
v
β= .
c
(1.1)
Thus, the distance between points A and C is βc∆te cos θ. The object moves at
speeds near the speed of light to point B where a second photon is released. The
distance between C and B is the projected distance that the object moves across the
10
plane of the sky and is equal to cβ∆te sin θ. c∆te represents the distance that the
initial photon travels toward the observer in the time that it takes the relativistic
object to move from point A to point B. The difference between the arrival times of
the two photons is ∆ta =∆te (1−β cos θ). The apparent speed is found by dividing the
apparent velocity (vapp ) by c, where the apparent velocity is the projected distance
on the sky divided by the difference in the arrival times of the photons (Equation 1.2)
βapp =
βc∆te sin θ
β sin θ
vapp
=
.
=
c
c∆ta
(1 − β cos θ)
(1.2)
Equation 1.2 can produce values for βapp which are greater than 1, making the
object appear to be moving faster than the speed of light as described earlier. This
can be seen analytically by increasing the value of β or decreasing the value of θ.
1.3.2
Beaming
When an object which emits radiation moves toward a stationary observer at a
relativistic speed the emission from the object may appear brighter than one would
expect. This is commonly called the lighthouse effect and is the result of aberration
of light, and is enhanced when the emitting object is moving at large velocities with
a small angle toward an observer. Assuming a point is moving with a velocity u′ in
frame S′ , the perpendicular motion of the object in the observer’s frame is described
as
u⊥ =
u′⊥
,
Γ(1 + vu′k /c2 )
(1.3)
where the Lorentz factor (Γ) is
1
.
Γ= p
1 − β2
(1.4)
A full derivation of Equation 1.3 can be found in Rybicki & Lightman [1979]. Now if
we assume that u′ ≡ |u′ |, and u′ =c, Equation 1.3 becomes
11
sin θ =
sin θ′
.
Γ(1 + β cos θ′ )
(1.5)
The cosine relativistic aberration relation is derived from the parallel motion of an
object as seen in 1.7.
uk =
u′k + v
(1 + vu′k /c2 )
.
(1.6)
Equation 1.6 can be transformed into the cosine relativistic aberration relation by
using the same assumptions as in the sine transformation in Equation 1.5.
cos θ =
cos θ′ + v/c
1 + (v/c) cos θ′
(1.7)
The Doppler factor can be introduced by rewriting Equation 1.5 as
sin θ = δ sin θ′
(1.8)
Thus, the Doppler factor is
δ=
1
,
Γ(1 − β cos θ)
(1.9)
1
.
Γ(1 + β cos θ′ )
(1.10)
where the inverse transformation (from S to S′ ) for δ is
δ=
The Doppler factor can also be used to describe the time dilation, as seen in Equation
1.11 below.
t = δt′
(1.11)
One should note the relativistic limit (β ≥ 0.7) where θ′ = 90◦ , which leads to
sin θ′ ∼ 1 and cos θ′ ∼ 0 leaving the sin θ term to approach 1/Γ which is related to β
by Equation 1.4. This would allow an observer to see the roughly half of the emission
from a relativistically moving object, which radiates isotropically in its rest frame,
12
swept into a cone which is described by a half-angle of 1/Γ (Figure 1.5). There are
very few photons which will have θ ≫ 1/Γ−1
Figure 1.5. Relativistic beaming of radiation which is emitted isotropically in the rest
frame K′ (S′ in the text) [Rybicki & Lightman, 1979].
1.3.3
Inverse-Compton Scattering
The phenomenon known as Compton scattering occurs when a photon interacts
with an electron, which has less energy than the photon. The photon loses energy and
the electron gains energy from this collision. When the previous process is reversed
it is referred to as inverse-Compton scattering. Figure 1.6 is a visual representation
of inverse-Compton scattering and shows a high energy electron which collides with a
low energy photon (ν). The electron transfers some of its energy to the photon, which
now has a higher energy than it did before the collision (ν ′ ). The example used in this
dissertation is described by an high energy electron from a blazar jet interacting with
the Cosmic Microwave Background (CMB) photons via inverse-Compton scattering.
The CMB photon is up-scattered by the electron, allowing for a net energy shift
from the electron to the photon. Synchrotron self Compton scattering (SSC) is a
second emission mechanism which can describe the emission from X-ray jets and is
also associated with inverse-Compton scattering. An SSC spectrum is observed when
13
Figure 1.6. Inverse-Compton Scattering
the synchrotron radiation produced by the jet is inverse-Compton scattered by the
same relativistic electrons which produced the initial synchrotron radiation.
1.4
Astronomical Instruments used in the MOJAVE Chandra Sample
1.4.1
The Very Large Array
The Very Large Array1 (VLA) is an array of radio antennas which can span 36
km in diameter when fully extended. The antennas can be moved radially to change
the resolution of the telescope, with each configuration having a label A, B, C, or D.
The A configuration (36 km radial antenna span) provides the best resolution while
the D configuration (0.6 km radial antenna span) provides the best sensitivity. There
are twenty seven 25 meter antennas that make up the 3 arms of the telescope, which
1
The VLA is a facility of the National Radio Astronomy Observatory, operated by Associated
Universities Inc., under cooperative agreement with the National Science Foundation
14
looks like a Y when fully extended. For this research I have chosen to use the A
configuration which provides a maximum angular resolution of 1.4′′ at a frequency of
1.4 GHz (λ = 20cm), and is referred to as the L band. This is the best configuration
for looking at extragalactic emission from blazars on the kiloparsec scale (kpc).
The angular resolution (Θ) of the telescope is related to the baseline (L) and the
wavelength (λ), as seen in Equation 1.12.
Θ≈
λ
L
(1.12)
This limits the maximum resolution that can be produced with the VLA to the kpc
scale for extragalactic objects such as blazars. Thus, to study the core and inner jet
(pc scale) structure a higher resolution is desirable.
1.4.2
The Very Long Baseline Array
The Very Long Baseline Array2 (VLBA) is an example of an instrument dedicated to performing Very Long Baseline Interferometry (VLBI). This operates on the
same principle as the VLA except that the baseline has increased, which produces
an increase in resolution such that the smaller parsec (pc) scale structure of AGN
can be studied (Θ ∼ 1 milliarcsecond at λ=2cm; Lister & Homan 2005). The VLBA
is a set of ten 25 meter antennas which are located between Hawaii and the U.S.
Virgin Islands. The entire network of antennas span a total distances of over 8500
km. Unlike the VLA, the dishes of the VLBA are not directly connected so the data
must be correlated after it has been collected digitally, with appropriate atomic clock
time stamping.
15
Figure 1.7. Chandra X-ray Observatory (Courtesy of NASA/CXC/NGST)
1.4.3
Chandra X-ray Observatory
The Chandra X-ray Observatory was launched on July 23, 1999 and has revolutionized X-ray astrophysics (Figure 1.7)3 . Originally named the Advanced X-ray
Astrophysics Facility, Chandra is a satellite which has a highly elliptical orbit and
is the one of the largest satellites ever launched. It was produced and tested in
Redondo Beach, California by TRW inc., which is now Northrop Grumman Space
Technology (NGST). Chandra itself has four nested pairs of iridium coated grazing
incidence mirrors (both paraboloid and hyperboloid) which focus the X-ray photons
on the detectors, which are located at the opposite end of the satellite (Figure 1.7).
Chandra’s Integrated Science Instrument Module (ISIM) houses the High Resolution
Camera (HRC) and the Advanced CCD (charge collecting device) Imaging Spectrometer (ACIS). The HRC and ACIS are used for the spatial detection of celestial objects,
while gratings can be moved in and out of the path of the emission to produce high
2
The VLBA is a facility of the National Radio Astronomy Observatory, operated by Associated
Universities Inc., under cooperative agreement with the National Science Foundation
3
http://chandra.harvard.edu/graphics/resources/illustrations/spacecraft labeled-72l.jpg
16
resolution spectroscopy. These instruments (HRC and ACIS) can detect X-rays from
0.2 keV to 10 keV [Garmire et al., 2003]. Chandra is currently operated by NASA at
the Smithsonian Astrophysical Observatory (SAO).
1.5
The Status of X-ray Jet Astrophysics
Prior to the development of X-ray astrophysics, jets associated with AGN were
studied using interferometric techniques with radio telescopes. The resolution of these
radio telescopes increased with time as the technology improved and the distance between the interferometer elements was increased. These telescopes were some of the
the first to image the jets associated with AGN. In the 1990s radio jet physics began
to lose its hold on the astrophysics community as newer areas of study were becoming
more tangible [Worrall, 2009]. This was short lived as jet physics was revitalized with
the launch of Chandra, an X-ray telescope which had the ability to resolve extragalactic AGN and their jets.
Before the launch of Chandra there were very few resolved jet detections associated
with AGN in the X-ray regime, predominately due to the resolution and sensitivity of
the satellites. The X-ray telescopes available at the time were Einstein and ROSAT
(Roentgen Satellite). Only bright sources with low redshifts were generally imaged
with these satellites, the majority of which were classified as radio galaxies. Examples
of early X-ray detections are M87, Centaurus A, and 3C 273 [Sambruna et al. 2004,
Marshall et al. 2005, and references within]. Since the launch of Chandra there have
been almost 100 new X-ray jet detections. Many of these detections can be found on
the Harvard University X-Jet website4 .
In 2004 and 2005 there were two major surveys that examined whether there was
a correlation between radio jet emission in QSOs and X-ray emission [Marshall et
al., 2005, Sambruna et al., 2004]. The Sambruna et al. [2004] survey was based on
surface brightness (S1.4GHz ≥5 mJy cm−2 ) of knots that were located at least 3′′ from
4
http://hea-www.harvard.edu/XJET/
17
the central nucleus of the AGN. The jet selection criteria that were applied to the
radio surveys were taken from Bridle & Perley [1984] and Liu & Xie [1992]. Their
selection criteria suggests that the sample is biased toward beamed jets and consists
of mostly FR II type quasars; 10 out of their 17 sources are considered FSRQs. The
rest of the sources are either SSRQs, BL Lacs, or radio galaxies. The Marshall et
al. [2005] sample was comprised of sources that were chosen from the Murphy et al.
[1993] and Lovell [1997] radio AGN surveys, which used the VLA and ATCA (Australian Telescope Compact Array) respectively. The selection criteria for the Marshall
et al. [2005] survey was based on the radio core flux densities (S5GHz,V LA > 1 Jy and
S2.7GHz,AT CA >0.34 Jy). Both of these surveys yielded X-ray jet detection rates of
∼ 60%.
The MOJAVE Chandra Sample (MCS) was established to study X-ray jets associated with FR II blazars and their possible X-ray emission mechanisms, and was
a subsample of the MOJAVE sample. This survey was created from selection criteria which biases the sample toward very fast, well collimated, powerful, beamed jets
which presumably have their X-ray emission presumably produced by the IC/CMB
mechanism, implying that these jets have high Doppler factors, relativistic speeds,
and small angles to the line of sight. The slection criteria required that all jets in the
sample had an extended flux greater than 100 mJy and that the terminal point of
the jet was at least 3′′ from the core. It was further culled by removing the sources
which were assumed to be less powerful (i.e., FR I objects). The MCS has an X-ray
jet detection rate of ∼ 77.78% on the kpc scale, which is almost a 20% increase from
previous X-ray jet surveys. This implies that the selection criteria, which is based on
extended flux (Sext ) and jet length, is a better predictor of X-ray jet emission than
the selection criteria associated with previous surveys [Hogan et al., 2011].
18
1.6
Thesis Description and Outline
The MCS is one of the first surveys to look specifically at the powerful FSRQ
subset of blazars on multiple wavelengths. This survey has increased the detection
rate of X-ray jets predicted by radio jet selection criteria by ∼ 20% when compared
to previous FSRQ surveys [Marshall et al., 2005, Sambruna et al., 2004]. Because
of the large redshift range of the MCS (0.033 ≤ z ≤ 2.099), I can examine the
effects of proposed X-ray mechanisms such as inverse Compton scattering off of cosmic
microwave background (IC/CMB) photons by relativistic electrons in the jets, which
is highly dependent on redshift. The selection criteria of this survey might be useful
for future surveys of blazars as well as for AGN which are located in the southern sky.
I construct spectral energy distributions for selected sources in the sample, in order to
further test the IC/CMB emission model. I also discuss jet bending and deceleration
in conjunction with the IC/CMB model, and their role in reconciling extreme bulk
Lorentz factors which are associated with some sources.
The thesis is laid out in the following manner: I describe the selection criteria for
the MCS as well as individual source observations in § 2. In § 3 the data reduction
and analysis is presented along with the implications of the IC/CMB emission model
when applied to the MCS. The spectral energy distributions for the sources in the
sample with optical, radio, and X-ray data are discussed in § 4. The thesis conclusions
are summarized in § 5. In this dissertation I use a standard cosmology with H0 = 71
km s−1 Mpc−1 , Ωm = 0.27, and ΩΛ = 0.73.
19
2. THE MOJAVE CHANDRA SAMPLE
2.1
Selection Criteria
Many of the X-ray jets that have been discovered to date, were discovered in
early surveys by Sambruna et al. [2004] & Marshall et al. [2005]. These surveys used
radio data associated with FSRQs, which were mainly selected from radio imaging
surveys, to search for X-ray jet emission with Chandra and other X-ray telescopes.
These surveys were not statistically complete and produced X-ray jet detection rates
of ∼ 60%. The MCS aims to improve extragalactic X-ray jet emission detection
by selecting targets associated with the MOJAVE sample along with other selection
criteria, thus making the MCS a complete sample of beamed FR II jets [Lister et
al., 2009b]. The original MOJAVE sample is comprised of 135 of the most powerful
AGN in the northern sky and is based on the following selection criteria [Lister et al.,
2009a]1 .
• Each source has a declination (δ) greater than −20◦
• Each source has a galactic latitude |b| > 2.5◦
• Each source has a total 2 cm VLBA flux density exceeding 1.5 Jy at any epoch
between 1994.0 and 2004.0 (>2 Jy for sources below the celestial equator)
Since the VLBA is insensitive to unbeamed radio emission, the MOJAVE sample is
highly biased toward blazar detection.
The MCS is based on the assumption that X-ray emission from extragalactic jets
with small opening angles is produced by the IC/CMB process. This leads to the
sample being focused on relativistic radio galaxys and blazars, with large Doppler
1
http://www.physics.purdue.edu/astro/MOJAVE/sample.html
20
factors. To optimize the likelihood of X-ray detection, the MOJAVE sample was
further culled by using the following criteria.
• Each source has more than 100 mJy of extended kpc emission at 1.4 GHz (VLA
A-array)
• Each source has a radio jet structure of at least 3′′ in length
• Each source is a member of the FR II class of AGN (i.e. BL Lac objects were
removed)
The BL Lac objects were removed from the sample because they are not as powerful as
the FSRQs and radio galaxies and possibly have a different X-ray emission mechanism.
This selection criteria provided a list of 27 QSOs and radio galaxies which comprise
the MCS [Hogan et al., 2011]. A complete list of the sources can be found in Table 2.1.
All of the sources were observed with Chandra, with most having integration times >
10 ks. Individual observation times as well as other information associated with the
Chandra observations are located in Table 2.2. Every source in the sample has an
1.4 GHz VLA A-array image available [Cooper et al., 2007, Kharb et al., 2010], and
a few sources have Hubble Space Telescope (HST) data available. The combination
of Chandra images, VLA (1.4 GHz) radio images, VLBA kinematic information, and
HST data sets (when available) provided the data that was used in the analysis of
the MCS.
2.2
Individual Source Observations of the MCS
The sources below have been observed in both the radio and X-ray bands. The
radio observations were made with the VLA and the X-ray observations were taken
with the Chandra X-ray Observatory. The radio/X-ray overlays are located below
and a description of how they were created is located in § 3.1. The radio and X-ray
profiles are located in Appendices A & B respectively. The position angles, which are
presented in Table 2.3, are measured from north toward east.
21
Table 2.1. MOJAVE CHANDRA SAMPLE
Source
Alias
z
Sext
β app
Reference
Obs ID
(1)
(2)
(3)
(4)
(5)
(6)
(7)
2.099
0.53
26.5 ± 4.2
Hogan et al. [2011]
9281
0.57
0.11
17.1 ± 1.1
Hogan et al. [2011]
9290
0.15
11.6 ± 0.8
Hogan et al. [2011]
9288
0106+013
OC 12
0119+115
0224+671
4C 67.05
0.523
0234+285
CTD 20
1.207
0.10
12.3 ± 1.1
Marshall et al. [2005]
4898
0415+379
3C 111
0.0491
2.70
5.9 ± 0.3
Hogan et al. [2011]
9279
0529+075
OG 050
1.254
0.13
12.7 ± 1.6
Hogan et al. [2011]
9289
0.872
0.12
19.8 ± 1.2
Sambruna et al. [2004]
2132
0605−085
0.595
0.51
8.6 ± 0.8
Hogan et al. [2011]
9280
1055+018
1045−188
4C 01.28
0.89
0.23
11.0 ± 1.2
Sambruna et al. [2004]
2137
1156+295
4C 29.45
0.729
0.20
24.9 ± 2.3
Coppi et al. [2002]
0874
1222+216
4C 21.35
0.432
0.96
21.0 ± 2.2
Jorstad & Marscher [2006]
3049
1226+023
3C 273
0.158
17.67
13.4 ± 0.8
Jester et al. [2006]
4879
1253−055
3C 279
0.536
2.10
20.6 ± 1.4
WEBT [2007]
6867
0.539
0.15
10.3 ± 1.1
Hogan et al. [2011]
9282
1334−127
1510−089
1641+399
3C 345
1655+077
0.36
0.18
20.2 ± 4.9
Sambruna et al. [2004]
2141
0.593
1.48
19.3 ± 1.2
Sambruna et al. [2004]
2143
0.621
0.20
14.4 ± 1.4
Marshall et al. [2005]
3122
1800+440
S4 1800−44
0.663
0.25
15.4 ± 1.0
Hogan et al. [2011]
9286
1828+487
3C 380
0.692
5.43
13.7 ± 0.8
Marshall et al. [2005]
3124
1849+670
S4 1849−67
0.657
0.10
30.6 ± 2.2
Hogan et al. [2011]
9291
1928+738
4C 73.18
0.302
0.36
8.4 ± 0.6
Sambruna et al. [2004]
2145
1957+405
Cygnus A
0.0561
414.18
0.2 ± 0.1
Wilson et al. [2001]
1707
0.672
0.30
18.1 ± 2.0
Hogan et al. [2011]
9284
0.295
0.37
7.9 ± 0.6
Hogan et al. [2011]
9283
0.901
0.31
5.6 ± 0.6
Hogan et al. [2011]
9285
0.859
0.88
14.2 ± 1.1
Marshall et al. [2005]
3127
0.576
0.14
13.5 ± 1.1
Hogan et al. [2011]
9328
2155−152
2201+315
4C 31.63
2216−038
2251+158
2345−167
3C 454.3
Note. — Columns are as follows: (1) IAU name (B1950.0); (2) Common Name; (3) Redshift from
NED; (4) Extended flux density (total - core) at 1.4 GHz (Jy); (5) Superluminal velocity in units of
c [Lister et al., 2009b]; (6) Reference for X-ray image; (7) Chandra observation ID number
22
Table 2.2. OBSERVATION LOG
Source
Live Time
Date
RA
DEC
(1)
(2)
(3)
(4)
(5)
0106+013
9.69
2007-11-21
1h8m38.771s
+1d35′ 0.317′′
0119+115
9.95
2008-10-27
1h21m41.595s
+11d49′ 50.413′′
0224+671
10.11
2008-06-27
2h28m50.051s
+67d21′ 3.029′′
0234+671
9.96
2004-06-24
2h37m52.40s
+28d48′ 09.00′′
0415+379
10.14
2008-12-10
4h18m21.277s
+38d1′ 35.800′′
0529+075
10.18
2007-11-16
5h32m38.998s
+7d32′ 43.345′′
0605−085
9.55
2001-05-01
6h07m59.70s
−8d34′ 50.00′′
1045−188
10.18
2008-04-01
10h48m6.621s
−19d9′ 35.727′′
1055+018
10.27
2001-01-09
10h58m29.60s
+01d33′ 59.00′′
1156+295
74.88
2000-06-29
11h59m31.80s
+29d14′ 43.80′′
1222+216
19.71
2002-11-06
12h24m54.40s
+21d22′ 47.10′′
1226+023
39.23
2004-07-28
12h29m06.20s
+02d03′ 00.40′′
1253+055
30.05
2006-01-17
12h56m11.20s
−05d47′ 21.50′′
1334−127
10.79
2008-03-09
13h37m39.783s
−12d57′ 24.693′′
1510−089
10.19
2001-03-23
15h12m50.50s
−09d06′ 00.00′′
1641+399
9.98
2001-04-27
16h42m58.80s
+39d48′ 37.00′′
1800+440
10.19
2008-01-05
18h1m32.315s
+44d4′ 21.900′′
1828+487
5.6
2002-05-20
18h29m31.80s
+48d44′ 46.60′′
1849+670
10.19
2008-02-27
18h49m16.072s
+67d5′ 41.680′′
1928+738
9.3
2001-04-27
19h27m48.50s
+73d58′ 02.00′′
1957+405
10.17
2000-05-26
19h59m28.30s
+40d44′ 02.00′′
2201+315
10.11
2008-10-12
22h3m14.976s
+31d45′ 38.270′′
2155−152
10.19
2008-07-10
21h58m6.282s
−15d1′ 9.328′′
2216−038
10.16
2007-12-02
22h18m52.038s
−3d35′ 36.879′′
2251+158
5.18
2002-11-06
22h53m57.70s
+16d08′ 53.60′′
2345−167
10.15
2008-09-01
23h48m2.609s
−16d31′ 12.022′′
Note. — Columns are as follows: (1) IAU name (B1950.0); (2) Chandra
exposure time in kiloseconds; (3) Date observed; (4) Right ascension of the
radio core position from NED(J2000); (5) Declination of the radio core position
from NED (J2000)
23
Table 2.3. MOJAVE CHANDRA SAMPLE JET MEASUREMENTS
Source
P Apc
P Akpc
Ri
Ro
Sr
νr
Count Rate
Sx
Pjet
X-Jet
(1)
(2)
(3)
(4)
(5)
(6)
(7)
(8)
(9)
(10)
(11)
0106+013
-127
180
1.5
8.0
526.7 ± 0.4
1.40
9.90 ± 1.11
9.9
< 1×10−10
Y
N
N
0119+115
6
35
1.5
8.0
22.2 ± 0.3
1.40
0.00 ± 0.38
< 1.2
5.54×10−1
0224+671
-5
-10
1.5
11.0
22.9 ± 0.7
1.40
-0.55 ± 0.45
< 0.8
9.62×10−1
1×10−10
Y
0234+285
-12
0
1.5
6.0
53.9 ± 0.4
1.40
6.09 ± 1.09
6.1
<
0415+379
65
63
1.5
100.0
50.6 ± 7.1
1.44
7.50 ± 2.49
7.5
2.44×10−6
Y
0529+075
-31
-145
1.5
8.0
69.2 ± 0.3
1.40
1.52 ± 0.65
1.5
1.99×10−4
Y
1×10−10
Y
0605−085
117
100
1.5
5.2
93.6 ± 1.1
1.42
10.85 ± 1.21
10.9
1045−188
149
125
1.5
10.0
167.7 ± 4.9
1.42
2.82 ± 0.82
2.8
<
4.74×10−8
Y
N
N
1055+018
-49
180
1.5
13.0
74.3 ± 1.8
1.42
-0.86 ± 0.97
< 2.1
9.01×10−1
1156+295
-2
-10
1.5
3.5
76.6 ± 0.7
1.52
-0.66 ± 1.15
< 2.8
8.65×10−1
1×10−10
Y
1222+216
-3
35
1.5
4.0
81.0 ± 0.3
1.40
8.22 ± 0.81
8.2
1226+023
-123
-135
1.5
20.0
4603.9 ± 11.8
1.45
115.50 ± 2.83
115.5
< 1×10−10
Y
1253−055
-124
-150
1.5
5.5
790.5 ± 2.2
1.66
3.91 ± 0.51
3.9
< 1×10−10
Y
1334−127
150
135
1.5
12.0
103.8 ± 0.2
1.49
17.07 ± 1.56
17.1
< 1×10−10
Y
14.9
<
1×10−10
Y
<
1×10−10
Y
1510−089
-31
165
1.5
6.0
63.3 ± 0.7
1.46
14.93 ± 1.63
<
1641+399
-89
-30
1.5
4.0
96.7 ± 1.8
1.51
3.86 ± 1.14
3.9
1655+077
-38
-50
1.5
5.0
46.7 ± 0.3
1.55
-0.83 ± 0.72
< 1.3
9.58×10−1
N
1×10−10
Y
1800+440
-159
-130
1.5
8.0
133.2 ± 0.5
1.51
6.28 ± 0.99
6.3
<
1828+487
-39
-40
1.5
3.0
28.4 ± 0.8
1.55
7.13 ± 1.43
7.1
< 1×10−10
Y
Y
1849+670
-45
0
5.0
20.0
8.3 ± 0.7
1.40
1.08 ± 0.40
1.1
1.36×10−6
1928+738
162
180
1.5
11.0
33.0 ± 1.3
1.42
5.24 ± 1.39
5.2
< 1×10−10
Y
1×100
N*
Y
1957+405
-79
-75
1.5
65.0
51381.4 ± 62.8
1.52
-10.40 ± 3.71
< 0.7
2155−152
-148
-170
1.5
12.0
231.3 ± 0.8
1.40
1.51 ± 0.85
1.5
5.00×10−3
Y
2201+315
-142
-110
1.5
10.0
31.1 ± 0.7
1.40
1.96 ± 1.05
2.0
1.54×10−3
2216−038
-172
135
1.5
15.5
164.2 ± 0.9
1.40
1.74 ± 0.78
1.7
4.94×10−4
Y
2251+158
-76
-50
1.5
5.5
585.0 ± 3.5
1.50
14.61 ± 2.74
14.6
< 1×10−10
Y
< 2.7
8.92×10−2
N
2345−167
124
-135
2.0
8.0
83.7 ± 0.4
1.40
0.65 ± 0.67
Note. — Columns are as follows: (1) IAU name (B1950.0) (2) Position angle of the pc-scale radio jet in (◦ ). All
position angles are measured from north through east (3) Position angle of the kpc-scale radio jet (◦ ) (4) Inner radius in
(′′ ) (5) Outer radius in (′′ ) (6) Observed flux density of the radio jet in mJy. The jet radio flux density is computed at the
given radii or the same region as for the X-ray count rate, given by the P Akpc , Ri , and Ro parameters (7) Observation
frequency of the radio image in GHz (8) Counts per kilosecond (9) The X-ray flux density (nJy) is given at 1 keV assuming
a conversion of 1 Jy s Count−1 , which is good to 10% for power law spectra with low column densities and spectral indices
(αx )near 1.5 (10) Probability of having more counts than those observed in the specified region under the null hypothesis
that the counts are background events. The jet is defined to be detected if Pjet < 0.0025 (see Section 2.2) (11) X-ray jet
detection (*) Cygnus A shows a correlation of X-ray and radio emission in the hotspot area which is not contained in the
jet detection area and thus, is considered a detection for this sample.
24
2.2.1
0106+013(OC 12)
Figure 2.1. Radio/X-ray overlay of 0106+013. The X-ray images were obtained from
Chandra with VLA 1.4 GHz radio contours overlaid in black and white (see § 3.1).
The black VLA contours are set at 5 times the rms noise level for the lowest contour,
with the exception of 0415+379 and 1849+670, which had their starting values set
to 10 and 2.5 times the rms noise respectively, and multiples of 2 greater than that
for each successive level. The white contours are offset from the black contours by
20%. The X-ray portion of each image has been energy filtered to a range of 0.5 to
7.0 keV in CIAO before being processed in DS9. The FWHM dimensions of the radio
restoring beam are denoted by a cross in the bottom corner of each image and are
also located in Table 2.4.
25
Table 2.4. VLA ARCHIVAL DATA
Source
Observation Date
Project
RMS Radio Noise
Bmaj
Bmin
Bmaj PA
(1)
(2)
(3)
(4)
(5)
(6)
(7)
0106+013
2004-09-19
AL634
1.4×10−01
1.64
1.49
100
0119+115
2004-09-19
AL634
9.8×10−02
1.53
1.44
55
AL634
1.5×10−01
1.42
1.13
78
0224+671
2004-09-19
0234+285
2004-09-19
AL634
2.4×10−01
1.20
1.11
52
0415+379
1982-06-14
LINF
1.9×10−01
1.60
1.47
168
0529+075
2004-09-19
AL634
4.6×10−02
1.70
1.35
63
0605+085
1993-01-17
AD298
3.0×10−01
1.84
1.11
166
1045−188
2007-06-30
AC874
3.4×10−01
1.00
1.00
90
1055+018
1992-11-18
AB631
1.8×10−01
2.00
1.33
52
1156+295
1984-12-24
AB310
1.2×10−01
1.39
1.32
43
1222+216
2004-11-20
AL634
1.7×10−01
1.23
1.11
54
AM297
2.0×100
1.00
1.00
90
1226+023
1990-03-23
1253+055
2001-01-11
W088D5
8.5×10−01
1.50
1.50
90
1334−127
1986-03-18
AD176
7.1×10−02
1.73
1.22
89
1510−089
1986-05-27
AB379
3.9×10−01
1.95
1.29
106
1641−399
1990-05-18
AS396
1.4×100
2.26
1.14
176
1655+077
1984-12-23
AB310
1.5×10−01
1.47
1.40
25
1800+440
1990-05-18
AS396
1.8×10−01
2.54
1.02
7
1828+487
1984-12-23
AB310
6.0×10−01
1.36
1.21
6
1849+670
2004-11-09
AL634
2.0×10−01
2.77
1.06
146
1928+738
1996-11-23
AS596
2.4×10−01
1.50
1.50
90
1957+405
1987-08-18
AC166
8.2×100
1.19
1.09
179
2155−152
2004-11-21
AL634
2.0×10−01
1.90
1.26
106
2201+315
2004-11-21
AL634
1.1×10−01
1.57
1.43
164
2216−038
2004-11-21
AL634
1.8×10−01
1.58
1.34
113
1.00
1.00
90
1.88
1.22
102
2251+158
1985-01-31
AC120
8.7×10−01
2345−167
2004-11-09
AL634
1.6×10−01
Note. — Columns are as follows: (1) IAU name (B1950.0); (2) Date observed; (3) Project code;
(4) Rms noise level of radio image in mJy beam−1 ; (5) Major axis for the radio beam in (′′ ); (6)
Minor axis for the radio beam in (′′ ); (7) Position angle of the radio beam major axis in (◦ )
26
0106+013 is a very powerful blazar, which is located farther away from us than
any other source in the MCS (z = 2.099) and has the second largest apparent speed
(βapp = 26.5). Physically this blazar shows a prominent radio jet which protrudes to
the south terminating at ∼ 5′′ from the center of the nucleus [Hogan et al., 2011].
The X-ray jet correlates well with the radio jet and shows emission until it reaches
the location associated with the terminal point of the radio jet. There is a small
amount of extraneous radio emission to the northeast with an approximate angle of
45◦ which is not correlated with the X-ray emission (Figure 2.1). This jet does not
show a lobe or hotspot in the counter-jet direction in either the X-ray or the radio
band (Appendix B & C). Kharb et al. [2011] shows that at higher radio resolution (5
GHz) along with Chandra rebinning the jet shows a gentile S shape as it progresses
from the core toward the terminal point. The jet also shows X-ray brightening and
dimming as it moves toward and away from these subtle bends respectively [Kharb
et al., 2011]. A Spectral Energy Distribution (SED) is shown in Kharb et al. [2011]
along with the deeper Chandra images.
2.2.2
0119+115
0119+115 (Figure 2.2) is an example of a FSRQ which shows no X-ray emission
above the level of the background emission besides the core structure. The radio jet
shows emission to the north, east, and west of the core, with the most prominent
radio feature located ∼ 6′′ to the north at an angle of 35◦ . The diffuse radio emission
to the east and west is located within a radius of 12′′ [Hogan et al., 2011].
27
Figure 2.2. Radio/X-ray overlay of 0119+115.
28
2.2.3
0224+671 (4C 67.05)
Figure 2.3. Radio/X-ray overlay of 0224+671.
Another example of a FSRQ with no correlation between the radio and X-ray
bands, except for the core structure, is 0224+671 (Figure 2.3). The radio image for
this AGN shows a long well collimated jet to the north ∼11′′ and a radio lobe in the
counter jet direction at a distance of ∼7′′ from the core [Hogan et al., 2011].
2.2.4
0234+285 (CTD 20)
0234+285 is a FSRQ which shows jet emission in both the radio and X-ray bands.
A radio jet is seen to the north of the core structure and shows a sharp bend before
it terminates abruptly. The X-ray emission correlates well with the radio emission
up until the bend. At and beyond the bend there is no X-ray emission above the
background level. There is an abundance of X-ray emission to the west of the core
which is most likely a readout streak. Readout streaks are produced when there is a
29
Figure 2.4. Radio/X-ray overlay of 0234+285.
pileup associated with the ACIS and is an artifact of the detector. The VLBI scale
emission also shows structured emission at a PA of −12◦ , which is closely aligned
with the P Akpc (see Table 2.3).
2.2.5
0415+379 (3C 111)
The first published image of this powerful radio galaxy was created by Linfield
& Perley [1984]. The radio/X-ray overlay shows 4 distinct radio knots in the radio
band, with three of those showing and excess of X-ray emission. The (1.4 GHz) radio
knots are only seen in the primary jet within 100′′ of the core at a position angle of
63◦ . Deeper VLBA images of this AGN show a pc scale jet with approximately the
same direction as the kpc scale jet. The lack of a counter-jet is most likely due to
Doppler boosting. Hotspots are seen in both the jet and counter jet directions even
though there is not appreciable emission between the core and the hotspot associated
with the counter jet. Both hotspots show large radio lobes which are approximately
30
Figure 2.5. Radio/X-ray overlay of 0415+379.
one quarter of the total jet length. The hotspot associated with the primary jet does
show X-ray emission above the background level, indicating an excellent correlation
between the X-ray and radio emission in this AGN [Hogan et al., 2011]. This radio
galaxy is very close to us (z = 0.0491) and thus probably does not have as small of
an angle to the line of sight as most of the more heavily beamed sources in the MCS.
The jet shows an angle of 7.9◦ according the IC/CMB calculations. Jorstad et al.
[2005] provided a value for the angle to the line of sight for 3C111 of 18.1 ± 5.0◦ ,
which is somewhat larger than the IC/CMB value that was tabulated. This jet has
a superluminal speed of 5.9c and thus, has a maximum value of the angle to the line
of sight of ∼ 19◦ for the pc scale jet [Lister et al., 2009a]. This source also shows an
extended radio flux (Sext = 2.70) which is significantly larger than the majority of
the rest of the MCS. It is interesting to note that all 5 of the sources which have a
radio Sext > 1 show X-ray jet emission associated with the jet. This is one of only
two radio galaxies in the MCS, the other being Cygnus A, and both show appreciable
X-ray emission [Wilson et al., 2001].
31
2.2.6
0529+075 (OG 050)
Figure 2.6. Radio/X-ray overlay of 0529+075.
This blazar shows radio emission in both the southeast and southwest directions.
The initial radio jet protrudes at a position angle of −145◦ and terminates at a
distances of ∼ 6′′ from the nucleus. There is also a radio feature which located to
the southeast of the core which could be an extension of the primary jet or emission
from the counter jet. The counter jet hypothesis is further supported by the lack of
X-ray emission in that area, because Doppler boosting of a jet away from the observer
would not produce much emission if the emission mechanism is IC/CMB. The main
jet emission in the X-ray band seems to coincide well with the primary radio jet, and
there is a lack of coincidence with the emission to the southeast. The pc scale jet
lies at a position angle of −45◦ , which leads to the conclusion that there is bending
between the pc and kpc scale jets in this instance [Hogan et al., 2011].
32
2.2.7
0605-085
Figure 2.7. Radio/X-ray overlay of 0605-085.
0605-085 shows an AGN with a radio jet structure which extends to the east at
an angle of 90◦ . When the X-ray image is superimposed there is a direct correlation
between the radio and X-ray jets within ∼ 4′′ of the core, with the exception of the
extra X-ray emission to the southwest. This extra emission is actually attributed to a
foreground star that in the Chandra field of view [Sambruna et al., 2004], which also
has a bright optical component. Sambruna et al. [2004] shows 5GHz radio data which
reveals two knots between the core and the termination point, which are unresolved
in the 1.4GHz data. There is HST data for this jet but it shows no emission above
the respective background level in the optical regime [Sambruna et al., 2004].
2.2.8
1045-188
This AGN shows strong, well collimated jet emission at a position angle of 125◦
and a large radio lobe directly opposed to it. This jet makes a ∼ 90◦ bend at a
33
Figure 2.8. Radio/X-ray overlay of 1045-188.
distance of ∼ 8′′ from the core. The X-ray jet emission follows the radio jet until this
bend and then shows a sharp decrease in emission. The counter jet lobe shows no
appreciable X-ray emission above the background. If the IC/CMB model is dependent
on the X-ray emission being beamed toward the observer then the jet bending in the
plane of the sky could cause unseen jet bending toward or away from the observer,
thus changing the amount of beamed emission that the observer would see [Hogan et
al., 2011].
2.2.9
1055+018 (4C 01.28)
1055+018 is an AGN that shows a radio jet extended toward the south which
ends in a hotspot with associated lobe. There is a radio lobe located directly to
the north presumably associated with the counter jet. There is no X-ray emission
associated with this source except for the core emission and pile up associated with
34
Figure 2.9. Radio/X-ray overlay of 1055+018.
the readout streak. There is no optical jet emission associated with the jet [Sambruna
et al., 2004]. It is noteworthy that Sambruna et al. [2004] classifies this sources as a
FSRQ/BL object.
2.2.10
1156+295 (4C 29.45)
1156+295 was one of the first objects to be imaged by Chandra (see Table 2.1
Column 7). There is no X-ray emission above the background level except for that
which is associated with the core.
2.2.11
1222+216 (4C 21.35)
This source was extensively studied by Jorstad & Marscher [2006].
The X-
ray/radio overlays show an elongated radio jet with significant bending which ends in
a terminal hotspot at a distance of ∼ 10′′ from the core. The radio jet hotspot has an
angle of 90◦ and there is lobe-like radio emission to the south of the core. The X-ray
35
Figure 2.10. Radio/X-ray overlay of 1156+295.
jet follows the initial nozzle of the jet (position angle ∼ 45◦ ) for a distances of ∼ 2′′
and decreases abruptly at the first bend. Again, this decrease could be attributed
to the bend in the jet, which might change the line of sight of the jet. Jorstad &
Marscher [2006] provide a 5 GHz radio image 1222+216 and cite 2 knots in the jet.
The knots are unresolved at 1.4 GHz. They also produce non-unique SEDs for one
of the knots (Jorstad & Marscher 2006 & Appendix D).
2.2.12
1226+023 (3C 273)
The jet associated with 1226+023 has been extensively studied by quite a few
groups [Marshall et al., 2001, Sambruna et al., 2001, Jester et al., 2006] due to its
unique X-ray, optical, and radio correlations. The radio/X-ray overlay that I produced
shows a radio jet extending to the southwest with an angle of −135◦ and terminating
at ∼ 22′′ from the core. The X-ray emission correlates spatially with the radio jet
rather well but peaks earlier than the radio jet does. The radio jet increases its
36
Figure 2.11. Radio/X-ray overlay of 1222+216.
emission along the jet until the hotspot where it reaches its maximum flux. The Xray emission, on the other hand, starts peaked at ∼ 14′′ from the core and decreases
its emission until it reaches the terminal region of the radio jet. This is also one of
the few jets where the pc and kpc scale position angles are very close to alignment
(Table 2.3) This source also has one of the lowest redshifts (z = 0.158) and largest
Sext (17.67) values in the MCS.
Marshall et al. [2001] showed that the optical (HST) and X-ray jets of 3C 273
are very similar in length and width, with the exception of the emission levels of
the first knot. A second major difference is that the X-ray emission fades as the
distances from the core increases, while the optical jet emission does not [Marshall
et al., 2001]. Sambruna et al. [2001] supports our initial visual inspection of this
source, and provides 2D surface brightness profiles which show how the emission
relates between the optical, X-ray, and radio wavelengths (see Fig. 2, Sambruna et
37
Figure 2.12. Radio/X-ray overlay of 1226+023.
al. 2001). Marshall et al. [2001] & Sambruna et al. [2001] differ on their choice of
emission mechanism. Marshall et al. [2001] provides an SED which shows that a single
population of electrons using a pure synchrotron emission mechanism can model the
first two knot fluxes, while Sambruna et al. [2001] shows that the external Compton
(∼IC/CMB) fits the data points used in the SED. Jester et al. [2006] believes that
the single zone IC/CMB model can only be used reliably in the first fifth of the jet.
They invoke either a two-zone IC/CMB model but warn that extreme bulk Lorentz
factors may be needed further down the jet, and or a two-zone synchrotron model,
where particle acceleration is related to a velocity shear which could produce the
X-ray emitters, to describe the jet features [Jester et al., 2006]. They also provide
SEDs to further reinforce the model choices (Figure 2, Jester et al. 2006).
38
Figure 2.13. Radio/X-ray overlay of 1253-055.
2.2.13
1253-055 (3C 279)
The image of 3C 279 shows a one sided radio jet that stretches toward the southwest until it bends very sharply toward the east at ∼ 4.5′′ from the core (initial
position angle −150◦ ). There is lobe emission to the north west at ∼ 12′′ from the
core with a position angle of −37◦ . The X-ray emission correlates well with the pri-
mary radio jet until the bend, where there is a sharp decrease in X-ray flux. Collmar
et al. [2010] shows multiwavelength analysis of the spectrum and morphology of 3C
279, by using Chandra, SWIFT, INTEGRAL, RXTE, and other telescopes associated with WEBT (Whole Earth Blazar Telescope). They produce SEDs which are
modeled using a leptonic one-zone SSC + EC model, and that the X-ray spectrum is
entirely produced by SSC emission ( Collmar et al. 2010 & Appendix D).
39
Figure 2.14. Radio/X-ray overlay of 1334-127.
2.2.14
1334-127
This blazar has an X-ray jet with a length of ∼ 6′′ that follows the radio jet
emission out to a 60◦ bend of the radio jet, then undergoes a drop in emission,
but still terminates at the same point as the radio jet. Both jets initially follow a
position angle of 135◦ . The emission characteristics in the bend region are significantly
different than the jet of 1045−188, which undergoes a sudden drop in X-ray emission
after the bend [Hogan et al., 2011].
2.2.15
1510-089
The (1.4 GHz) radio and X-ray jets extend to the southeast, bending slightly
toward the south. The X-ray jets terminates well before the radio jet, at a distance
of ∼ 5′′ , which happens to coincide with the bend of the radio jet. The radio jet
40
Figure 2.15. Radio/X-ray overlay of 1510-089.
continues past the bend for another 5′′ until it ends. Sambruna et al. [2004] also
presents the HST optical data for this source but there are no optical counterparts
for the X-ray or radio knots. The higher resolution 5GHz radio data shows three
possible X-ray knot detections [Sambruna et al., 2004].
2.2.16
1641+399 (3C 345)
3C 345 shows a one sided radio jet which protrudes to the northwest of the core.
The X-ray emission from this object follows the radio emission closely and terminates
at roughly the same spot, ∼ 3′′ from the core. The 5GHz radio data shows a more
defined knot structure in the radio jet, which is also closely aligned with the X-ray
knot [Sambruna et al., 2004]. Sambruna et al. [2004] also presents HST optical data
which corresponds to the knot structure. The optical data point helps constrain the
SED which helps determine if the IC/CMB model does indeed describe the primary
emission mechanism of X-rays in blazar FR II jets (§ 4.3 & Appendix D). Deeper
41
Figure 2.16. Radio/X-ray overlay of 1641+399.
Chandra observations of this source were obtained by Kharb et al. [2011] along with
new optical data from HST, which was used to construct the SED. The new Chandra
images indicate that the X-ray hotspot is actually located closer to the core than the
radio hotspot [Kharb et al., 2011].
2.2.17
1655+077
This AGN shows a short jet extending to the southeast and a longer radio jet
which extends to the northwest and then makes a sharp (∼ 90◦ ) bend toward the
southwest [Marshall et al., 2005]. There is no X-ray emission above the background
level other than the core region. VLBA data shows pc scale knots along the position
angle of −38◦ (Marshall et al. [2005] and reference within) which is closely aligned
with the kpc scale position angle (−50◦ ).
42
Figure 2.17. Radio/X-ray overlay of 1655+077.
2.2.18
1800+440 (S4 1800-44)
This AGN shows two sided emission. The primary radio jet is located southwest
of the core and the counter jet lobe emission is located to the northeast. The X-ray
jet follows the radio jet at a position angle of −130◦ until the first apparent bend
and then abruptly terminates (∼ 3′′ ). The radio jet continues for another 3′′ past
the bend. This is another example where the X-ray jet flux decreases beyond a radio
knot located at a bend in the jet [Hogan et al., 2011].
43
Figure 2.18. Radio/X-ray overlay of 1800+440.
44
2.2.19
1828+487 (3C 380)
Figure 2.19. Radio/X-ray overlay of 1828+487.
3C 380 is an AGN that harbors radio emission on both sides of its core with the
more prominent jet protruding toward the northwest. The position angles of the pc
and kpc scale jets are within 1◦ of alignment (see Table 2.3 columns 2 & 3). The
X-ray jet profile shows a significant difference between the primary jet and counter
jet emission (Appendix B). Marshall et al. [2005] places the X-ray emission knot at a
distance of ∼ 1.8′′ from the core. This source has the third largest Sext (5.43) in the
MCS.
2.2.20
1849+670 (S4 1849-67)
1849+670 is a FSRQ that shows a radio jet to the north and a radio lobe structure
to the south. The radio emission to the north shows emission until ∼ 15′′ from the
core. The X-ray jet follows the radio jet closely for ∼ 9′′ and then abruptly stops.
The overlays show no correlation between the radio lobe and the X-ray band when
45
Figure 2.20. Radio/X-ray overlay of 1849+670.
inspected visually. Again, this supports the IC/CMB emission mechanism as it is
assumed that the primary jet is boosted toward the observer and the counter jet is
boosted away.
2.2.21
1928+738 (4C 73.18)
4C 73.18 is an AGN which shows an elongated radio jet which extends to the south
with a slight bend toward the east. It also shows radio emission to the north centered
around ∼ 6′′ from the core. The X-ray emission briefly follows the radio jet and stops
abruptly at a distance of ∼ 7′′ from the core. The radio jet bends and continues to a
distance of ∼ 17′′ from the core until it terminates. Sambruna et al. [2004] uses 5 GHz
radio data along with Chandra images and finds only one X-ray knot, which has an
optical (HST) counterpart. The SED composed by Sambruna et al. [2004] indicates
that this sources is different than most from our sample by implying that the emission
46
Figure 2.21. Radio/X-ray overlay of 1928+738.
mechanism for the X-ray component is pure synchrotron emission instead of IC/CMB
emission. The radio jet of the source also differs morphologically from most of the
other FR II sources in the sample, as it is not well collimated and terminates in a
poorly confined lobe, which reinforces the different emission mechanism proposed by
Sambruna et al. [2004] for this source.
2.2.22
1957+405 (Cygnus A)
Cygnus A is the second closest (z = 0.0561) and has the largest extended flux
(Sext = 414.18) when compared to the rest of the MCS. This source is the second
of two radio galaxies in the sample (the other being 3C 111) meaning that its jets
are traditionally viewed at a larger angle to the line of sight than blazars. The radio
image shows a well collimated jet to the northwest and extremely large lobes that
surround the hotspots from both the jet and counter jet. The X-ray image shows a
correlation at the terminal hotspot locations, which implies that there is a jet, but
47
Figure 2.22. Radio/X-ray overlay of 1957+405.
an excess of X-ray emission near and around the core prevents any jet detections.
Visually the only cores and hotspots of both the radio and X-ray images are aligned.
This excess of X-ray emission is attributed to emission from the cavity of Cygnus A
[Wilson et al., 2006]. This cavity is interpreted as energy which the jet can not move
efficiently to the lobe, and is often referred to as a ”cocoon” (Wilson et al. [2006]
and references within). Wilson et al. [2001] proposed that the emission model for
Cygnus A is Synchrotron Self-Compton (SSC) radiation, as this is the mechanism
of radiation for most FR II type radio galaxies. This could imply that there is a
difference between the radiative processes of blazars and radio galaxies, even though
they are both members of the FR II class.
2.2.23
2155-152
2215-152 is an FSRQ which shows radio emission to the north and south of the
core. The pc scale jet is oriented to the south so I assumed that the primary jet on
the kpc scale was also oriented to the south. The southern X-ray jet is observed until
a distance of ∼ 4′′ where it abruptly stops before the radio jet which shows a sharp
48
Figure 2.23. Radio/X-ray overlay of 2155-152.
decrease in emission between ∼ 8′′ [Hogan et al., 2011]. The northern lobe is centered
around a point which is located at ∼ 5′′ from the core.
2.2.24
2201+315 (4C 31.63)
4C 31.63 shows an elongated radio jet which loses collimation before it terminates
in a hotspot at a distances of ∼ 37′′ with a position angle of −100◦ . This AGN also
has a radio hotspot associated with the counter jet which is located at ∼ 45′′ (P Akpc
= 58◦ ) from the core. The X-ray emission correlates with the radio jet up until a
distance of ∼ 5′′ from the core, where there is also a decrease in radio jet emission
and loss of collimation (see appendices A & B). This X-ray jet detection is considered
marginal due to its lack of visual correlation with the radio jet and probability of jet
detection value (Pjet , Table 2.3) which is very close to the threshold value. A longer
exposure time is needed to produce a clear visual correlation.
49
Figure 2.24. Radio/X-ray overlay of 2201+315.
2.2.25
2216-038
This blazar shows two radio knot structures and a hotspot on the kpc scale oriented
toward the southeast and terminates at ∼ 15′′ from the core. There is also a radio
lobe structure located to the north at a distance of ∼ 19′′ . The closest radio knot
is located at a distance of ∼ 5′′ from the core, while the more prominent knot is
centered around ∼ 9′′ from the core at the same P Akpc . The X-ray emission is visually
correlated well with the second knot and the hotspot but not the first knot. There is
no X-ray emission associated with the radio lobe above the level of the background.
This jet also bends significantly from the pc to the kpc scale showing a position angle
misalignment of roughly 60%.
50
Figure 2.25. Radio/X-ray overlay of 2216-038.
Figure 2.26. Radio/X-ray overlay of 2251+158.
2.2.26
2251+158 (3C 454.3)
The radio image of this FSRQ shows a well collimated jet which ends in a knot
or hotspot at a distance of ∼5′′ from the core. The X-ray emission mirrors the radio
51
emission and terminates at roughly the same point [Marshall et al., 2005]. Marshall
et al. [2005] presents VLBA data that shows the pc scale jet curving to align with
the position angle of the kpc X-ray jet (−50◦ ).
2.2.27
2345-167
Figure 2.27. Radio/X-ray overlay of 2345-167.
2345-167 is a small one sided radio blazar which shows radio jet emission in a
southwest direction until a distance of ∼ 5′′ from the core where there is a sharp decrease in flux. Visually there seems to be a marginal correlation and the X-ray profile
confirms this. This sources does have a Pjet value which is close to the threshold, but
does not surpass it and thus is not considered a detection. A longer exposure time
on this source could allow for an X-ray detection.
52
3. DATA REDUCTION AND ANALYSIS
3.1
X-ray Radio Overlays
Figures 2.1 through 2.27 present the X-ray-radio overlays for the sources in the
MCS. The following procedure was used to produce these images. 1.4 GHz data were
first obtained using the VLA A-array data from the NRAO1 data archive and our
own observations [Cooper et al., 2009, Kharb et al., 2010]. These data were reduced
following the standard procedures in the Astronomical Images Processing System
(AIPS). Standard calibrators were used to initially calibrate the amplitude and phase
of the sources. Then, the tasks IMAGR and CALIB were used iteratively to selfcalibrate the image and the sources. This self calibration was performed on both
the amplitudes (with successively decreasing solution intervals) and phases of the
visibilities (with solution interval times typically set to less than 0.5 mins in CALIB)
until convergence in image flux and structure was achieved. This procedure led to
the production of radio maps which had a typical rms noise of ∼ 0.2 mJy beam−1 .
The specific radio noise levels for each source map are located in Table 2.4. The
average FWHM restoring beam of the radio images was assumed to be ∼ 1.4′′ , which
is slightly worse than the FWHM of Chandra, which was estimated to be ∼ 0.75′′
[Hogan et al., 2011, Marshall et al., 2005].
The X-ray portions of the overlays were created by using both archival data and
proprietary data from Chandra [Hogan et al., 2011]. The observation dates and
exposure times for all of the MCS targets are listed in Table 2.2. Chandra maps
for the overlays were created using the DS9 imaging tool, by using the level 2 event
files. First, the event files were loaded into CIAO (Chandra Interactive Analysis of
1
The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.
53
Observations), and were filtered to an energy range of 0.5 to 7 keV. Once filtered to
the appropriate energy range, the files were loaded into DS9 for imaging. The files
were then smoothed using the smoothing tool, which was set to the Gaussian setting
with a kernel radius of 3 pixels, where a pixel is equivalent to 2′′ . The pseudo-color
scale was then adjusted so that the core structures were oversaturated (black), so
that the X-ray emission from the jets above the noise level could be easily identified
through the use of visual inspection. At this point both the radio and X-ray images
were ready to be merged. I started with the radio image and then aligned the X-ray
image to it using the WCS (World Coordinate System) frame matching setting in
DS9. There were a few sources which had core positions that were slightly misaligned
when the overlays were inspected. To compensate for the misalignment of the cores,
the radio and X-ray images were registered by using the Fv program in the Ftools
package provided by NASA2 [Blackburn, 1995]. The images were manually shifted by
a small amount (generally of order of less than 2 pixels) so that the nucleus of each
AGN was aligned between the X-ray and radio bands [Hogan et al., 2011]. The radio
beam size is represented by a cross in the lower right hand corner of the radio/X-ray
overlays, and has specific properties listed in Table 2.4. The beam data were extracted
from AIPS by using the IMSTAT task, which provided the length of the major and
minor axes of the beam as well as the major axis angle. The rms radio noise was
then obtained from IMSTAT (Table 2.4). The lowest radio contour level was set to 5
times the noise level, except for 0415+379 and 1849+670, which had their lowest level
contours modified to better represent the radio jet emission. 0415+379 and 1849+670
had their lowest contour set to 10 and 2.5 times the noise level, respectively. Each
successive contour past the first was set to twice the level of the previous contour.
This process was repeated for a total of 10 iterations. The black contour curves
represent these contours and the white contour curves are set 20% higher than the
black ones. Once the contour level values were known, they were drawn in DS9 by
using the contour parameter tool. The levels were manually entered into the tool and
2
Information about Ftools can be found at http://heasarc.gsfc.nasa.gov/ftools/
54
placed on the images by using a contour smoothness setting of 1. Once placed on the
radio image, the curves were then moved to the X-ray image, which was located on a
seperate frame in DS9. The images were then saved in a tiff format, as DS9 does not
have the ability to save files in a postscript format. The files were then converted to
postscript files by using GIMP2 (GNU Image Manipulation Program).
3.2
X-ray and Radio Jet Analysis
The X-ray jet detection portion of the analysis was completed using the IDL
(Interactive Data Language) programming language and scripts which were developed
by Herman Marshall and modified by Brandon Hogan [Marshall et al., 2005, Hogan
et al., 2011]. In essence, the scripts take the radio and X-ray FITS files and compare
the regions of the radio and X-ray jets which are specifically chosen based on radio
jet criteria (i.e., radio jet length and width). The script starts by identifying the core
position of both the radio and X-ray FITS files which were calculated via the method
described by Marshall et al. [2005]. The X-ray nucleus location was first deduced by
fitting Gaussians to the one dimensional histograms obtained from events which were
detected within a distance of 30′′ of the hypothesized core region, which was chosen
by visual inspection of the FITS files. This estimated centroid position was then
used to acquire a more accurate central reference point for the core by repeating the
previous process using a region which was defined by a radius of 3′′ from the previously
calculated rough centroid position. This produces a more accurate centroid position
and reduces the contamination effect of extended jet emission, which can bias the
core positions.
Poisson statistics were then used to test for the existence of an X-ray jet. The
radio images were used to create a rectangular region on the X-ray map which defines
the original shape and size of the radio jet. The rectangular region is defined by
the parameters P Akpc (kiloparsec position angle), Ri (jet inner radius), and Ro (jet
outer radius, see Table 2.3). The length of each region was allowed to vary, but the
55
width of each was fixed at 3′′ . The rectangular region was also offset from the center
of the core by 1.5′′ to eliminate the emission from the nucleus, with the exception
of 1849+670 and 2345-167, which were offset by 5′′ and 2′′ respectively. Shortening
these rectangular regions can eliminate jet contamination from the core emission
region which is often associated with the elongated restoring beams found in 1.4 GHz
radio data [Hogan et al., 2011].
The algorithm assumes that the radio jets show no bending on the kpc scale
and that the area 90◦ to the primary jet axis is free of real emission from the jet.
This assumption is valid for all of the jets in the MCS except for 0119+115, which
shows radio jet emission on both sides of the core perpendicular to the direction of
the primary jet axis, but does not show any X-ray emission above the level of the
background X-ray emission. Hence, the perpendicular jet emission is inconsequential
in this case. Then, the radio emission was plotted in two dimensions to show a profile
of the jet emission along the axis of the jet and in the region perpendicular (90◦ ) to the
jet (Appendix A). The jet emission and the background emission (emission located in
the region perpendicular to the jet axis) were then subtracted to eliminate the core
emission (bold curve). The X-ray profiles were created by using the radio jet emission
region and the region located at the opposite side (180◦ ) of the core (Appendix B).
This background region was chosen because counter-jets associated with blazars are
not seen very often in X-rays, probably due to the effects of Doppler boosting [Worrall,
2009]. Visual inspection of the MCS also reinforces this concept, as very few X-ray
counter jets are seen. The CCD readout streaks associated with Chandra also cannot
contaminate the counter jet area as they are offset from the main jet axis, and hence
also the counter jet axis. The X-ray counts in these regions were compared using
Poisson statistics, with a Poisson probability threshold for jet detection set to 0.0025
[Marshall et al., 2005, Hogan et al., 2011]. The chosen probability threshold produces
a 5% chance that a false detection will occur in 1 out of every 20 sources. The X-ray
fluxes were computed from count rates using a conversion factor of 1 µJy per count
s−1 , which is accurate to about 10% for typical jet power law spectra [Marshall et al.,
56
2005]. The sources in Table 2.3 with negative values in the count rate column have less
X-ray emission in the jet region than in the corrosponding background region. Since
visually there is often not X-ray emission in the counter jet region, these negative
count-rates are often small and are associated with a null detection in the X-ray jet
column. The analysis method indicated that there are X-ray detections in all of the
sources except for 0119+115, 0224+671, 1055+018, 1156+295, 1655+077, 1957+405
(Cygnus A), and 2345−167, although 1957+405 was treated as an X-ray detection.
Visual inspection of the latter source shows coincident hotspot emission in both the
X-ray and radio bands (see §2.2.22 for further discussion), as well as an excess of Xray emission around the core which contaminates the regions that have been chosen
for jet and counter jet emission. The sources without X-ray detections show little
to no X-ray emission above the background level except for their respective nuclei
(Appendix A, Table 2.3), despite the fact that their radio structure, redshift values,
βapp values, and Sext values are comparable to other sources in the MCS. Tables 3.1
and 2.3 list all of the relevant X-ray emission data for the sample.
3.3
The Single Zone IC/CMB Model
The analysis for the X-ray portion of the MCS was done using a method similar
to the method used by Marshall et al. [2005]. There are a few major assumptions
that are associated with the model and are as follows. The magnetic fields are in
equipartition with the particle energies for the kiloparsec scale jets. This assumption
allows for the minimum energy magnetic field strength of the jets to be calculated. If
the magnetic field strength and CMB photon density are known, the photon energy
density can be compared to the magnetic field energy density. This leads to the use
of a parameter K, which is calculated by combining the magnetic field strength with
the ratio of X-ray (inverse-Compton) and radio (synchrotron) luminosities. The K
parameter is a function of β, (Equation 1.1) and θ, the angle to the line of sight.
Another important parameter is the apparent speed (βapp ), which is a parameter
57
Table 3.1. MOJAVE CHANDRA SAMPLE BEAMING MODEL PARAMETERS
Source
αrx
R
V
B1
K
δ
θ
Γpc=kpc
Γkpc,min
Γkpc,decel
(1)
(2)
(3)
(4)
(5)
(6)
(7)
(8)
(9)
(10)
(10)
0106+013
0.94 ± 0.01
0.0726
1.6×103
148.
13 ± 2
3.6
4.2
99+33
−29
1.9
1.9+0.1
−0.1
< 0.1996
7.7×102
29.
< 19
4.3
> 6.2
36
...
...
0119+115
> 0.88
0224+671
> 0.91
< 0.1356
9.9×102
26.
< 15
3.9
> 8.9
20
...
...
0234+285
0.84 ± 0.01
0.4362
1.1×103
58.
29 ± 7
5.4
7.8
17+2
−2
2.7
3.2+0.6
−0.5
0415+379
0.83 ± 0.02
0.5590
3.7×101
19.
51 ± 12
7.1
7.9
6+1
−1
3.5
6.0+1.9
−1.9
26+7
−6
33+4
−3
11+2
−2
1.7
1.8+0.2
−0.1
3.3
3.7+0.5
−0.6
2.1
2.5+0.4
−0.3
0529+075
0.93 ± 0.02
0.0846
1.6×103
57.
11 ± 2
3.3
8.4
0605−085
0.84 ± 0.01
0.4416
7.2×102
61.
43 ± 10
6.6
5.2
1045−188
0.95 ± 0.02
0.0639
1.1×103
49.
17 ± 3
4.1
10.8
1055+018
> 0.92
< 0.1053
2.3×103
42.
< 13
3.6
> 9.3
19
...
...
1156+295
> 0.91
< 0.1308
3.3×102
65.
< 27
5.2
> 4.4
62
...
...
1222+216
0.85 ± 0.01
0.3921
1.9×102
53.
61 ± 14
7.8
6.3
32+4
−3
3.9
5.1+1.5
−1.0
1226+023
0.92 ± 0.01
0.0941
1.6×102
91.
75 ± 17
8.6
6.0
15+1
−1
4.3
6.1+1.3
−1.3
96.
17 ± 3
4.1
5.3
54+5
−5
11+2
−2
27+3
−2
35+5
−3
2.1
2.2+0.2
−0.2
3.6
5.1+1.4
−1.4
4.6
5.4+1.1
−1.0
2.9
3.2+0.3
−0.4
26
...
...
25+3
−3
16+2
−1
114+15
−17
8+1
−1
2.7
3.0+0.3
−0.4
4.1
5.6+2.1
−1.2
2.2
2.2+0.3
−0.2
3.5
5.7+1.8
−1.8
1
...
...
55+10
−13
9+1
−2
7+1
−2
21+3
−2
1.7
1.7+0.2
−0.2
2.7
3.9+0.9
−0.9
1.6
1.9+0.2
−0.1
2.8
3.2+0.4
−0.4
22
...
...
1253−055
1.02 ± 0.01
0.0166
4.3×102
1334−127
0.83 ± 0.01
0.6044
1.1×103
40.
50 ± 14
7.0
7.5
41.
82 ± 21
9.1
4.7
1510−089
0.81 ± 0.01
0.8796
2.4×102
1641+399
0.90 ± 0.02
0.1448
3.1×102
60.
33 ± 7
5.7
5.4
1655+077
> 0.92
< 0.1008
4.7×102
46.
< 19
4.3
> 7.2
52.
28 ± 6
5.3
6.6
1800+440
0.89 ± 0.01
0.1709
9.4×102
1828+487
0.81 ± 0.01
0.8936
2.3×102
52.
67 ± 16
8.2
6.2
1849+670
0.84 ± 0.02
0.5029
2.2×103
18.
18 ± 4
4.2
3.7
1928+738
0.83 ± 0.01
0.6039
3.6×102
27.
48 ± 12
6.9
8.1
< 0.0001
3.5×101
151.
<2
1.2
> 40.9
0.0253
1.6×103
53.
10 ± 2
3.2
6.1
27.
29 ± 7
5.3
9.9
1957+405
2155−152
> 1.32
0.99 ± 0.03
2201+315
0.87 ± 0.03
0.2438
3.2×102
2216−038
0.97 ± 0.02
0.0408
2.8×103
49.
9±1
3.0
15.8
0.0910
7.7×102
101.
30 ± 6
5.5
7.0
< 0.1227
7.2×102
44.
< 22
4.7
>7.6
2251+158
0.93 ± 0.01
2345−167
> 0.91
Note. — Columns are as follows: (1) IAU name (B1950.0); (2) Radio to X-ray spectral index (3) X-ray to radio luminosity
ratio (4) Synchrotron emission region volume kpc3 (5) Minimum energy magnetic field (µG) (6) K, given by Eq. 3 (7)
Angle to the line of sight determined by the IC/CMB method with no bending and deceleration (8) Doppler beaming
parameter, assuming no deceleration or bending between the pc and kpc scales (9) Jet bulk Lorentz factor assuming no
deceleration or bending between pc and kpc scales (10) Minimum value for the bulk Lorentz factor associated with the kpc
scale jet (11) Bulk Lorentz factor when the non-deceleration assumption is relaxed
58
obtained from VLBI observations and is a function of both β and θ. θ is calculated
by solving the K and βapp equations simultaneously, along with the assumption that
the value of βapp is the same for the pc scale radio jet and the kpc scale X-ray jet.
The Doppler factor (δ) and the bulk Lorentz factor (Γ) can be calculated once the
value for θ is known [e.g., see Harris & Krawczynski 2002, Marshall et al. 2005, Hogan
et al. 2011].
The single component synchrotron model has difficulties in explaining the X-ray
emission in powerful blazar jets, presumably due to the small viewing angles and
amount of Doppler boosting that occurs. Physical quantities for the X-ray emission
have been derived using a standard IC/CMB model. The calculations were obtained
by using the same IC/CMB basic assumptions as Marshall et al. [2005], which were
obtained from Harris & Krawczynski [2002], and are stated below.
• The energy density of the CMB occurs at the peak of the blackbody distribution.
• The jet frame equipartition holds between the particle energy densities and the
magnetic field, with a filling factor (Φ) of 1.
• The low energy spectral index for the synchrotron spectrum continues unchanged below the current range of the instruments used to measure it.
If the second assumption fails then relativistic protons will contribute to the particle
energy density and beaming will become much more intense. The quantity
18.85 C12 (1 + k)Lsync
B1 =
ΦV
2/7
(3.1)
is defined first, where B1 the spatially averaged, minimum energy magnetic field of
the jet in Gauss, when there is no Doppler boosting (δ = 1). C12 is a weak function of
the low frequency spectral index of the synchrotron spectrum (αr , where Sν ∝ ν −αr ),
Φ is the filling factor, Lsync is the synchrotron luminosity (calculated from the radio
flux and luminosity distance), k is the baryon energy fraction parameter, and V is the
emitting volume [Pacholczyk 1970, Harris & Krawczynski 2002, Marshall et al. 2005,
59
Hogan et al. 2011]. The values used for the constants are; k =0, C =5.7×107 , αr =0.8,
and Φ=1. The emitting volume for the jet is calculated using the Ri and Ro values
defined in Table 2.3 by taking the difference of the two values and then assuming
a cylindrical cross section given by the width associated with the Chandra FWHM
value (0.75′′ ). The VLA A-array (FWHM = 1.4′′ at 1.4GHz) radio data results in
larger derived emitting volumes than the Chandra FWHM. This discrepancy causes
the magnetic field value (B1 ) to be considered a minimum value for the tabulated
values. This magnetic field disparity can be resolved by adjusting the filling factor
Φ. If Φ is decreased from the original value of 1 by a factor of 10 the magnetic field
quantity B1 would change by roughly a factor of 2 (Marshall et al. [2005]).
The X-ray to radio luminosity ratio (R) is computed by using Equation 3.2.
αr −αrx
Sx νxαr
νx
Sx (ν/νx )−αr
=
=
,
R=
−α
α
Sr (ν/νr ) r
S r νr r
νr
(3.2)
where νr and νx are the radio and X-ray frequencies at which the flux densities
Sr and Sx are observed, respectively. The jet frame value for Lsync is affected by the
redshift and the luminosity distance which are both accounted for in the algorithm.
Equation 3.2 is valid under the assumption that the X-ray and radio frequencies are
far from the terminal points of the synchrotron and IC spectral breaks. The values for
ν r are located in Table 2.3 and ν x =2.42×1017 Hz. The equation for the K parameter
was first presented in Marshall et al. [2005], and is a quantity which is composed of
constants and observed quantities:
K = B1 (aR)1/(αr +1) (1 + z)−(αr +3)/(αr +1) b(1−αr )/(αr +1) .
(3.3)
The constants used in Equation 3.3 are a=9.947×1010 Gauss−2 and b=3.808×104
Gauss and can be found in Harris & Krawczynski [2002]. The values for these constants are found by using the equipartition assumption to equate the expected and
observed values of the ratio of X-ray to radio energy densities (R). Thus, K is a dimensionless number that is solely a function of the viewing angle and the jet speed, as
shown in Marshall et al. [2005], which can be translated into the beaming parameters:
60
K = Γδ(1 + µ′j ) =
1 − β + µ − βµ
,
(1 − βµ)2
(3.4)
where µ′j is defined in Equation A9 from Harris & Krawczynski [2002] and is described
by an angle transformation between the jet frame and the observers frame. Equation
3.4 can be solved for µ for given β and K, as seen in Equation 3.5 [Marshall et al.,
2005]. The variable µ is the cosine of θ, and used primarily to simplify the calculations.
µ=
1 − β + 2Kβ − (1 − 2β + 4Kβ + β 2 − 4Kβ 3 )1/2
.
2Kβ 2
(3.5)
Equation 3.5 is used for converting the angles to the jet frame for use later in the
IC/CMB emission model calculations, and is the negative root associated with the
solution of Equation 3.4 when solved for µ. At this point the method diverges from
the Marshall et al. [2005] and Harris & Krawczynski [2002] analysis. Marshall et al.
[2005] made the assumption that all kpc jets have Γ = 10. This assumption defined a
value for β (Equation 1.4), and made Equation 3.5 solvable for µ. I have chosen not
to use the previous assumption but to use the pc scale radio information to solve for
the values of θ and consequentially Γ and δ, with the assumption that the jets have
the same βapp values on pc and kpc scales. Equations 3.6, 3.7, and 3.8 can be used
to solve for θ, Γ and δ.
β=
βapp
p
βapp µ + 1 − µ2
θ = arctan
Γ=
2
βapp
2βapp
+ δ2 − 1
2
βapp
+ δ2 + 1
2δ
(3.6)
(3.7)
(3.8)
Specifically, β can be represented in terms of βapp and µ as seen in Equation 3.6. This
value can be substituted into the K equation (Equation 3.5), which makes µ now a
function of βapp and K. Marshall et al. [2005] showed that a change in B1 by 60% only
affects the calculated value of θ by ∼ 10%. Thus, the values for θ (which implies µ)
61
are quite reliable. Once θ is known, Equations 3.7 and 3.8 can be used to solve for Γ
and δ [Hogan et al., 2011].
The main source of error in the K parameter is the spectral index (αr ), which had
a value set to −0.8 for the IC/CMB calculations. Common observed values for αr at
kpc scale distances are between −0.7 and −0.9. Since, actual measurements αr do
not exist for the MCS, Monte Carlo error analysis was carried out on the sample, by
defining a Gaussian distribution of αr with αr = −0.8 and σαr = 0.1 [Hogan et al.,
2011]. The 1 σ error values for K are located in Table 3.1.
3.4
Scenarios Associated with the IC-CMB Model
There are three scenarios that are described below which can be associated with
the IC/CMB model.
• IC/CMB model with no jet bending and no deceleration
• IC/CMB model with deceleration and no jet bending
• IC/CMB model with both deceleration and jet bending
Each scenario is described in detail in the following sections as well as the possible
implications associated with each assumption. Jet bending with no deceleration is
not considered because a solution which allows for only vertical translation on the
θ − Γ plots in Appendix C cannot rectify the extreme values of Γ in some sources.
This is discussed more specifically in §3.5.
3.4.1
IC/CMB model with No Jet Deceleration or Bending
Equations 3.4 and 3.6 can be expressed graphically as curves on the θ − Γ plane
(see Appendix C), where β in Equation 3.6 is a function of Γ and µ is the cosine of
θ. The blue dashed curve describes the kpc equation defined by the IC/CMB model
and the black solid curve describes the pc scale, which was defined by the VLBI
62
kinematic information. The intersection of the two curves produces a viewing angle
and bulk Lorentz factor pair that satisfies both equations, under the assumption that
both the pc and kpc scale jets have the same value for βapp . The error values for
these curves are produced by attributing the error from the βapp and K values, which
defines the range of error for Γ. The error is depicted on the graphs as the dotted lines
which flank the curves (Figures C.1 through C.4). Some sources, such as 0415+379,
1800+440, and other jets in the sample, have an uncertainty associated with βapp
which produces large amounts of uncertainty on Γ (Table 3.1).
The majority of the sources in the MCS have reasonable Γ values which are agreeable with previous surveys of X-ray jet emission associated with inverse Compton
models. These models often postulate that the bulk Lorentz factors are on the order
of Γ≈10 or greater. There are other models, such as the Bayesian parameter-inference
method, which also provide Γ values for FR II jets. The Γ values provided by Mullin
& Hardcastle [2009] are significantly smaller than the ones produced by the IC/CMB
method, having values of ∼ 1.2 − 1.5. The jets in the Mullin & Hardcastle [2009]
sample, however, are selected on the basis of isotropic lobe emission, which is more
representative of the entire FR II population than the MCS. The jets in their sample
tend to have large angles to the line of sight and probably have electron populations
which are described by a different emission mechanism than the MCS. Further support of the MCS bias toward large values of Γ is presented by Lister & Marscher
[1997], which states that unbiased orientation samples of radio jets are likely to have
much lower Γ values than blazar samples. This is due to the relatively steep power
law distribution of jet speeds in the parent population. Both of the radio galaxies in
the sample show visual confirmation of two sided radio lobe emission which dominates
their 1.4 GHz radio maps. The quasars in the sample are usually dominated by core
emission instead of lobe emission.
Recently, Cooper [2010] has produced the pc scale viewing angle distribution for
the MOJAVE sample, which was derived from Monte Carlo simulations. The model
uses the luminosity function for the MOJAVE parent population [Cara & Lister,
63
2008] to model the 1000 trial populations of the 135 sources. Γ values for the population are described by a power law ranging from 3 to 50 with an index of −1.5.
The results approximate a Poisson distribution of the pc jet viewing angles, which
is peaked around 2◦ . This distribution for the viewing angles is expected because of
the highly beamed nature of the MOJAVE sample. Since the MCS is a relativistic,
highly beamed sub-sample of the MOJAVE sample, one should expect to see a small
angle bias in it also.
There are two sources (0106+013 & 1849+670) which show unusually large values
for the Γ parameter, when using the IC/CMB model along with pc scale jet kinematic information. These sources have Γ values which exceed 70. Alternatively, the
largest measured value of Γ in the Hovatta et al. [2009] sample is 65, for 1730−130.
The βapp (∼ 35 c) value attributed to 1730−130 is large when compared to the rest
of that sample. Other similar samples include the Padovani & Urry [1992] sample
and the MOJAVE sample, which contain no superluminal speeds > than 50 c [Lister
et al., 2009b]. Lister & Marscher [1997] show that βapp,max and parent population
Γmax should be fairly analogous for large flux limited blazar samples. Similar to the
Hovatta et al. [2009] result for 1730−130, the two extreme sources in the MCS have
the smallest values of θ and the largest βapp values.
3.4.2
IC/CMB model with Jet Deceleration
Deceleration between the pc and kpc scales is one way to rectify the large Γ values.
This deceleration associated with the jet is caused by the transfer of power to the IGM
or other medium, which is traditionally in the form of kinetic energy [Georganopoulos
& Kazanas, 2004]. The misalignment of knots and other jet structures between radio,
X-ray, and other bands can often be seen in one-zone models which describe the
deceleration of jets. The MCS comprises a few sources which also have misaligned
knots and hotspots. A second way that deceleration helps reconcile the large values for
Γ is by widening the beaming cone. This can be done under the assumption that jets
64
decelerate from ultra-relativistic speeds to mildly relativistic and even sub-relativistic
speeds near the terminal points of the jets (§1.3.2).
The extreme values of Γ can be lowered to a more reasonable range if deceleration
is applied to the IC/CMB model. This is done by looking for a set of horizontal lines
(solutions) which intersect the pc (black) and kpc (blue) scale curves (Appendix C).
These lines are given by the cyan shaded region on each graph. The red dashed line
shows the best fit viewing angle for the original set of assumptions. If deceleration is
allowed, then the solutions on the low Γ tail of the kpc scale curve are viable solutions
that do not require jet bending. The range of possible kpc scale Γ values are listed
in Table 3.1. These values are generally narrow and significantly smaller than the
Γ ≈ 10 assumption which is often invoked with the IC/CMB model. If jet bending
is combined with deceleration, Γ ≈ 10 can be obtained for all sources. The Γmin,decel
values are calculated for the sources in the sample with X-ray jets. Values associated
with Γmin,decel in Hogan et al. [2011] are a range of numbers with no calculated error
attributed to them. The values presented in this thesis have the error associated with
the K equation provided and are listed in Table 3.1.
3.4.3
IC/CMB model with Deceleration and Jet Bending
FR II jets can display misalignment between the pc and kpc scales [Kharb et al.
2010, Conway & Murphy 1993, Moore et al. 1981]. These non-linear morphologies
are often highly exaggerated by projection effects associated with the geometry of
system. The MCS comprises some jets in which bending between the pc and kpc scales
can lower the Γ value without changing other requirements, such as the relationship
between the bulk Lorentz factor and the superluminal speed (Γ ≥ βapp ). The results
of adding both deceleration (acceleration) and jet bending to the IC/CMB model can
be seen graphically by allowing the jet to lie anywhere on black curve for the pc scale
and anywhere on the blue curve for the kpc scale. This effectively allows for the two
curves to be connected by any linear combination of two points. If Γ ≈ 10 is required
65
then in most cases the jet bends outward from the pc to the kpc scale. The beaming
parameters can still be constrained by the IC/CMB model if both deceleration and
bending are allowed. There is a lower limit set for the bulk Lorentz factor Γmin , which
is described by Equation 3.9, when the value for µ is set to 1 in Equation 3.4. These
limits are tabulated in Table 3.1, and are usually between 1.6 and 2.7, except for
sources 0415+379 and 1334−127, which have Γmin values greater than 3.5 [Hogan et
al., 2011].
K
Γmin = √
2 K −1
(3.9)
Equation 3.4 also sets a limit for the value of θkpc,max . This can be seen graphically for
individual sources in Appendix C and is a lengthy algebraic function of K [Marshall
et al., 2005]. The θkpc,max values are between 8◦ and 20◦ , which is typical for FR II
type jets as they are associated with viewing angles which are ≤ 20◦ (see §1.2).
The bulk Lorentz factor (a function of β) is limited by by the relationship between
βapp and θ. This relationship, described by Equation 3.6, confines Γ ≥ βapp , and
−1
θ ≤ 2 tan−1 (βapp
). X-ray jet observations of blazars can provide more useful limits
on jet deceleration, if the amount of jet bending was to be constrained by future
independent observations. A second way to improve the results presented here is to
pursue the IC/CMB method with a larger sample, which could improve the statistics.
3.5
Sext as an X-ray jet predictor
The MCS shows a correlation between the radio and X-ray jet emission in 77.78%
(21/27) of its sources (assuming Cygnus A has an X-ray detection). This corresponds
to a ∼ 20% increase in the detection rate from previous FSRQ surveys done by Marshall et al. [2005] & Sambruna et al. [2004], which were based on radio surveys of
FSRQ. We have found that the extended flux densities, Sext , are closely correlated
with the detection rate of the X-ray emission. Kharb et al. [2010] have presented a
interesting trend implying that there is a relationship between parsec scale apparent
66
jet speeds and extended radio luminosity in the MOJAVE blazars. Thus, X-ray jet
detection and jet speed could also be related. I found a 100% X-ray jet detection
fraction for Sext > 300 mJy (Figure 3.1) and a significantly lower detection rate (∼
57%) for sources with Sext values below 300 mJy. Using an extended flux density
threshold value as a selection criterion could prove to be conclusive way to predict
X-ray jet detections in FR II blazars and radio galaxies when selected from previously
known radio band information.
3.6
Kolmogorov-Smirnov Tests
MCS Kolmogorov-Smirnov tests were produced for three different cases; the βapp
values with respect to the detection of sources, the βapp values with respect the Sext
threshold value (300 mJy), and the redshift value with respect to the detection of the
sources [Hogan et al., 2011]. The threshold value for the probability associated with
the K-S test (p) was set to 0.05 in each of these. Values of p which are larger than
the threshold do not reject the possibility that both populations could have the same
parent population whereas a p value below the threshold would reject the possibility.
In all three cases the p value is larger than the threshold value. [Hogan et al., 2011]
A second set of K-S tests were ran on the MCS to see if it was representative of
the total MOJAVE population. When the redshift values of the MCS and MOJAVE
samples are put into a two sample Kolmogorov-Smirnov goodness-of fit-hypothesis
test, the test produced a result that rejected the null hypothesis (p value of 0.0036),
and thus they do not originate from the same population of objects. This is most
likely because the FR I objects (presumably BL Lac objects) were removed from the
sample and changed the sample statistics. The histograms representing the MOJAVE
and MCS are shown in Figures 3.2 and 3.3 respectively.
A few more K-S tests were ran on the MCS with respect to the MOJAVE sample
where I have removed the BL Lac objects from the MOJAVE sample. The K-S test
67
10
Quasar or Radio Galaxy w/ X−ray jet
Quasar w/out X−ray jet
9
8
Number
7
6
5
4
3
2
1
0
0
0.1
0.2
0.3
0.4
0.5
0.6
0.7
0.8
0.9
1
Sext (Jy)
Figure 3.1. Histogram relating the source population to the Sext value. All sources
with a Sext ≥ 300 mJy show a correlation between the X-ray and radio bands at some
level.
68
associated with the Sext produced a result which rejects the null hypothesis that
the two populations are from the same parent population (p value = 1.1301×10−12 )
and the test associated with redshift also fails (p value = 9.5348×10−4 ), which is
slightly worse than the K-S test p value for the redshift when the BL Lacs are left
in the MOJAVE sample. These failed K-S tests are most likely due to the selection
criteria which selects only the most powerful sources with elongated radio jets from
the MOJAVE sample. Interestingly, the K-S test for the βapp values showed that
it was possible that the MOJAVE and MCS samples could originate from the same
parent population (p value of 0.0884). The K-S test for the apparent speeds had a
few less sources in the MOJAVE portion of the sample because βapp values for only
107 of the 135 sources could be calculated at this time. Thus, the selection criteria
does not change the distribution of βapp with respect to the MCS, even though it
alters the distributions of redshift and Sext .
69
Figure 3.2. Histogram representing the redshift distribution of the MOJAVE sample
Figure 3.3. Histogram representing the redshift distribution of the MCS sample
70
3.7
Viewing Angle
Equation 3.10 relates the change in position angle on the plane of the sky between
pc and kpc scale jets (∆PA) to the angle to the line of sight with respect to the
observer (θn ), the intrinsic misalignment angle between the pc and kpc scales assuming
a simple bend (ζ), and the azimuthal angle of the jet (φ) [Conway & Murphy, 1993,
Moore et al., 1981].
tan(∆P A) =
sin ζ sin φ
cos ζ sin θn + sin ζ cos θn cos φ
(3.10)
The azimuthal angle is not known in any of these sources and thus is treated as a
free parameter in this discussion. This equation can be simplified by assuming that
the line of sight of the pc scale jet is small, and that the angle between the pc and
kpc jets is also small. When small angle approximation is applied to Equation 3.10
for these two variables it becomes
tan(∆P A) ≈ sin(φ)
θn
ζ
.
+ cos(φ)
(3.11)
The small angle approximation for θn is valid because the MOJAVE sample is comprised mostly of blazars which have small angles to the line of sight on the pc scale
Cooper [2010]. There are three cases which can be examined for the MCS with the
use of Equation 3.11.
• sources where ∆PA is small (< 45◦ )
• sources where ∆PA is large (45◦ ≤ ∆ PA≤ 90◦ )
• sources where ∆PA approaches and exceeds 90◦
When a source has a small value for ∆PA (≤ 45◦ ) the denominator of Equation 3.11
must be large. This implies that θn must be large when compared to ζ, which cancels
out the effect of the azimuthal angle in most cases. Thus, any discrepancy between
Γ and δ is likely to require deceleration, and not exceptional jet bending. When the
71
value for ∆PA becomes larger (45◦ ≤ ∆ PA≤ 90◦ ) the ratio between θn and ζ also
has to change for a random value of φ. In this case ζ approaches θn . Large values
−1
of ∆PA (≥ 90◦ ) would require that ζ approaches and surpasses the value of βapp
.
Equation 3.11 can still be satisfied because most sources in the MCS have large βapp
values. Moore et al. [1981] states that large values of ∆PA can be obtained with small
values of θmax , where θmax is the largest value of θn which is likely to occur [Hogan et
al., 2011]. A value of ∆PA which approaches and exceeds 90◦ can only be obtained
when θn ≤ ζ. The unknown value of φ always plays a role in calculation because
if φ=0, Equation 3.11 always produces a value of 0 for ∆PA regardless of what the
values for θn and ζ are. Conway & Murphy [1993] also states, that for their angle misalignment calculations, they cannot obtain a scenario where there is a peak in
their distribution of misalignment angles around 90◦ . So even for a favorable ζ-θ ratio,
it still requires a very specific azimuthal angle to produce a misalignment angle ≥ 90◦ .
Figure 3.4. Position Angle Misalignment Associated with the MCS
72
A distribution for the ∆PA values (|P Akpc − P Apc |) for the MCS is located in
Figure 3.4. It is fairly obvious that the majority of the sources in the MCS have
∆PA values which are less than 60◦ and do not require the scenario where θn ≤ ζ
is needed. There are only three sources that have ∆PA values which are larger than
60◦ (0529+075, 1055+018, and 1510-089). Two of these sources show X-ray and
radio correlation for the kpc scale jet, indicating that they are not fundamentally
different from the rest of the MCS. The above discussion, combined with the figures
in Appendix C provides evidence that supports the conclusion that bending between
the pc and kpc scales cannot alone solve the problem of large bulk Lorentz factors
associated with the extreme sources in the MCS. This is not to imply that jet bending
is not needed, as it is still a viable way to lower the Γ values in the extreme sources
when combined with deceleration. Bending is very important if the assumption that
Γ ≈ 10 on kpc scales is upheld, as the combination of bending and deceleration is the
only way to reconcile the assumption.
73
4. SPECTRAL ENERGY DISTRIBUTIONS
4.1
General Information
The Spectral Energy Distribution or SED is a fundamental indicator of the kind of
emission mechanism(s) that can produce the radiation from jets in AGNs. It is widely
accepted that the radio and optical emission from extragalactic jets are predominantly
synchrotron radiation. The portions of the SED which the most controversial, is
the area associated with the X-ray and γ-ray regimes. The X-ray emission can be
described by SSC, IC/CMB, or even synchrotron radiation, and is often influenced
by the amount of beaming associated with the source. The optical component plays
a key role in constraining which emission model will best fit the SED. If the optical
point is aligned on the same spectral slope as the X-ray and radio points, the emission
is best fit with a single zone Synchrotron model. If the optical flux is below a linear
extrapolation of the radio (synchrotron) and the X-ray fluxes the model will most
likely be IC/CMB, or perhaps SSC. The emission modeling script that I have chosen
to use approximates the synchrotron, SSC and IC/CMB radiation as three separate
curves and is described in Krawczynski et al. [2004]. The solid line represents the
synchrotron radiation, while the dot dashed and dashed lines represent the IC/CMB
and SSC radiation respectively. The observed values for ν and νF (ν) are then plotted
along with the curves using an IDL plotting script.
74
Table 4.1. SED PARAMETERS
Source
Alias
z
DL
δ
radius
B
wpsoll
γmin
γmax
n
(1)
(2)
(3)
(4)
(5)
(6)
(7)
(8)
(9)
(10)
(11)
0415+379
3C 111
0.0491
2.11 × 108
1 × 1022
1.3 × 10−5
3.1 × 10−10
3 × 105
2.7
2
15
Note. — Columns are as follows: (1) IAU name (B1950.0); (2) Common Name; (3) Redshift from NED; (4) Luminosity
distance to sources (pc); (5) Doppler factor; (6) Radius of source (cm); (7) Magnetic field value (Gauss); (8) Photon
energy flux per volume (erg cm−3 ); (9) Minimum electron energy; (10) Maximum electron energy; (11) Power-law index
of the electron energy distribution
Figure 4.1. Spectral Energy Distribution for the hotspot associated with the primary
jet in 3C 111. The γ-ray data are considered upper limits and are represented as
downward arrows, while the radio, optical and X-ray data points are represented as
diamonds. The solid line represents the synchrotron radiation. The dot dashed and
dashed lines represent the IC/CMB and SSC radiation respectively.
75
4.2
3C 111 (0415+379) SED
Appendix D shows all of the previously constructed SEDs and Figure 4.1 shows
the newly constructed SED for the primary jet’s hotspot of 3C 111 (0415+379), as
it is the only source in the MCS with new X-ray and or optical data associated with
it that has not been published. The SED algorithm that I chose uses the parameters
shown in Table 4.1 to construct the curves associated with each emission type (see
§ 4.2.1 through 4.2.4). One assumption that was made when creating the SED for
0415+379, was that the Doppler factor decreases to a value of ∼ 2 at the hotspot
region, as the jet is assumed to be less relativistic at the terminal hotspot than at
the nozzle. The 3C 111 hotspot SED shows that it is possible to model the X-ray
emission with an IC/CMB emission curve. Specifically, the SED shows a synchrotron
curve which intersects the radio points and the optical point. The IC/CMB curve
intersects the X-ray point and is well below the upper limit of the γ-ray radiation
(downward arrows) which is measured by the Fermi space satellite. The γ-ray data
are considered upper limits because they represents the flux from the entire source.
Fermi does not have the capability to resolve the hotspot alone. The IC/CMB model
is often chosen because the magnetic field is close to the equipartition magnetic field,
which is usually on the order of µG. The magnetic field associated with the hotspot
of 3C 111 is 1.3 × 10−5 G, which is roughly the same order of magnitude as what is
expected (Table 3.1). The Doppler boosted equipartition magnetic field of the jet is
found by dividing B1 by δ. The difference in calculated magnetic field and the magnetic field needed to construct the SED could be attributed to the hotspot location
because the non-boosted (δ = 1) equipartition magnetic field was calculated for the
area close to the nozzle of the jet (Table 3.1). The larger magnetic field could also
be attributed to the low δ assumption that was made, implying that the jet it less
relativistic near the hotspot.
The algorithm parameters were fixed for z, luminosity distance (DL ), δ, and radius before the script was executed. The power-law index was then obtained from the
76
slope of the two radio points, while the electron energies (γmin and γmax ) were set to
values similar to the values from the Sambruna et al. [2004] sample and was adjusted
slightly to provide a better fit. The magnetic field was assumed to be similar to the
minimum energy magnetic field calculated for the jet by the IC/CMB model and
slightly manipulated. Lastly, the photon energy flux per volume (wpsoll ) was shifted
to align the curves and the data points.
4.2.1
Obtaining the Radio Fluxes
The radio fluxes were extracted for a region which mirrored the X-ray hotspot area
(∼ 2′′ radius). The data were extracted with AIPS by using the task IMSTAT for the
given region. This region was defined by using TVBOX to select the region on the
tv window. The radio fluxes were already in Jy, so they were converted to erg cm−2
sec−1 for the given frequency that they were observed at (1.4GHz and 5GHz). The
two points in the radio band constrain the power-law index for the electron energy
distribution.
4.2.2
Obtaining the Optical Fluxes
The optical data information was taken from the HST drizzle file by extracting
the region from DS9. The ACS extracted regions provide the number of electrons/sec.
The fits header has a PHOTFLAM keyword, which when multiplied by the previous
quantity, produces a flux in terms of erg cm−2 s−1 Å−1 . This information along with
the observing wavelength allowed for the procurement of the SED optical point. The
optical point constrains the well, under the assumption that it lies on the synchrotron
curve.
77
4.2.3
Obtaining the X-ray Fluxes
The X-ray information was taken from the number of counts in the selected circular region associated with the hotspot of 3C 111 (∼ 2′′ radius). After the region
was selected I used the virtual observatory, which is located under the analysis tab
to open the Chandra-Ed Archive Server1 . The counts in regions tool was accessible
once the archive server was opened. This tool was used to procure the counts in the
region which was previously defined. The number of counts was then entered into
the Chandra Proposal Planning Toolkit under the PIMMS2 tab to estimate the flux.
The Chandra cycle number, energy range, galactic NH, redshift, photon index, and
count rate for the object were needed as parameters to produce the estimation of
the flux in erg cm−2 s−1 . This point along with the γ-ray emission constrains the
IC/CMB portion of the SED. This specific source shows an X-ray hotspot which has
only 9 counts detected for the 10 ks Chandra observation. The small statistics for the
hotspot in the X-ray regime makes X-ray spectral analysis difficult, as traditionally
40 or more counts are needed, and thus spectral slope (’bowtie’) limits are not placed
on the X-ray point in the SED.
4.2.4
Obtaining the γ-ray Fluxes
The γ-ray data points were calculated by using the information from the Fermi
1FGL data set. Each of the points on the SED had an energy range defined by Emin
to Emax for a given photon flux which was observed by the Fermi space satellite. The
other given quantity was the spectral index, which I shall refer to as γ. Equation 4.1
is used to describe the relationship between the differential photon flux (dN /dE) and
γ, where A is a constant.
1
2
chandra-ed.cfa.harvard.edu/archive.html
http://cxc.harvard.edu/toolkit/pimms.jsp
78
dN
= AE γ
dE
(4.1)
When solving for A, the equation is integrated and rearranged to look like Equation
4.2, where n is the photon flux.
A=
n(γ + 1)
γ+1
− Emin
γ+1
Emax
(4.2)
After solving for A, I then found the average energy value and solved for the quantity
of νF (ν) by converting from MeV cm−2 sec−1 to erg cm−2 sec−1 , as seen in Equations
4.3 & 4.4. C is a constant with a value of 1.6021 × 10−12 ergs eV−1 used to convert
the equation into ergs.
2+γ
νF (ν) = Eavg
AC
(4.3)
and the frequency (ν) is defined as
ν=
Eavg
h
(4.4)
where h is 6.58211 × 10−16 eV/sec. The final values for ν and νF (ν) are located in
Table 4.2.
4.2.5
Uniqueness of the 3C 111 Hotspot SED
The goodness of fit was assessed by eye for the hotspot associated with the primary
jet of 3C 111. The overall fit for the SED is unique since there is optical data available
to constrain the well between the synchrotron and the IC/CMB curves. Without the
optical data point the two curves were not constrained horizontally and could be
shifted left or right. The curves were still constrained vertically from the data points.
79
Table 4.2. 3C 111 SED INFORMATION
Telescope
ν
νF (ν)
(1)
(2)
(3)
VLA
1.4×109
2.5×10−14
VLA
5.0×109
2.9×10−14
HST
3.8×1014
5.4×10−16
Chandra
2.4×1017
2.7×10−15
Fermi
3.0×1023
7.1×10−12
Fermi
9.9×1023
5.1×10−12
Fermi
3.0×1024
2.9×10−12
Fermi
9.9×1024
1.8×10−12
Note. — Columns are as follows:
(1) Telescope used for the observation (2) Flux in Hz (3) Flux multiplied by a function of the flux in erg
cm−2 sec−1
80
The SSC curve can also be fit to the data points by adjusting the magnetic field,
photon flux and other parameters, but traditionally requires more extreme values for
many of the parameters. The most common example of this is that the magnetic
field is often assumed to be far from the equipartition magnetic field value in the
SSC model. The synchrotron curve is fully constrained by the optical and radio data
points and Power law index is constrained to a unique value from the slope of the
two radio points. An example of a non-unique SED is presented in Appendix D by
Figure D.5.
4.3
Individual SED Notes
The majority of the jet knot SEDs, which are presented in Appendix D, show
that the X-ray emission mechanism is predominately IC/CMB, but there are other
sources that have a more complex or different basic SED structure. 0415+379 (3C
111), 1222+216, and 1641+399 show SEDs where the X-ray portion of their jets can
be explained as IC/CMB [Sambruna et al., 2004, Jorstad & Marscher, 2006]. These
SEDs show a radio and optical region described by a synchrotron curve which shows
a sharp cutoff at about 1015 Hz, and an X-ray curve which models emission from 1016
Hz to 1025 Hz or greater. 1253-055 (3C 279), on the other hand, was modeled by
Collmar et al. [2010] and shows that SSC emission dominates the X-ray portion of
the spectrum. 1928+738, which is the only source classified as a FSRQ/BLL source
[Sambruna et al., 2004], has an SED which approximates synchrotron radiation as the
sole emission mechanism for the radio, optical, and X-ray radiation. This is unusual
for a jet which has a small angle, as SSC and synchrotron emission tends to represent
radio galaxies and other lobe selected objects. It is expected that the majority of
relativistic beamed sources with small angles to the line of sight have their X-ray
emission embodied by the IC/CMB model. High redshift X-ray sources can be as
bright as the low redshift X-ray sources because the CMB density has a (1+z)4 dependence [Sambruna et al. 2004 and references within]. Tavecchio et al. [2000] shows
81
Figure 4.2. The jet from 3C 273 observed with Chandra (top), HST (middle, λ=620
nm), and the VLA (bottom, λ=3.6 cm) [Jester et al., 2006]. The emission levels of
the radio optical and X-ray bands have peaks located at different parts of the jet.
The jet originates at the left side of the image and terminates at the right end.
that SSC calculations require a very debeamed jet for the magnetic field to approach
equipartition for the blazar 0637−752. If δ > 1 the magnetic field diverges from equilibrium very quickly in 0637−752. This further supports the previous assumption
that FR II type blazars are most likely relativistically beamed sources.
The SED for the source 3C 273 is probably the most interesting in the MCS
because of the unique emission trends. This low redshift source shows radiation in
the optical, radio, and X-ray bands for the entire length of the jet (Figure 4.3). The
optical flux is fairly constant from the nozzle to the hotspot on the kpc scale jet,
while the X-ray image shows fluxes near the core which are larger than fluxes located
82
further downstream. The radio emission does not correlate spatially with the X-ray
emission as the jet shows more radio emission at the terminal point of the jet. Three
groups have composed SEDs for the jet of 3C 273 [Sambruna et al., 2001, Marshall
et al., 2001, Jester et al., 2006]. Sambruna et al. [2001] shows SEDs for four major structures associated with 3C 273. These features are described in Figure 2 of
Sambruna et al. [2001], where the jet is broken up into quarters and each quarter is
represented by a letter ranging from A to D. The IC/CMB emission models SEDs
for each of these positions are reproduced in Appendix D. Marshall et al. [2001] also
produced a set of SEDs representing the different areas of the jet associated with 3C
273, and is represented visually in Figure 1 of their paper. The SEDs that Marshall et
al. [2001] produced showed different emission mechanisms for the different sections of
the jet. The early portions of the jet (knot A) have points which define a synchrotron
emission curve in the SED diagram, while the portions of the jet located further from
the core show X-ray spectral softening, which changes the emission mechanism. The
more recent study by Jester et al. [2006] presents SEDs for the jet associated with
3C 273, which seems to embody IC/CMB emission for the first few regions of the jet
(region A through B2, Figure 4.3), until the X-ray spectrum is softer than the radio
spectrum [Jester et al., 2006]. For regions past the first few, Jester et al. [2006] proposed two different two-zone model interpretations for the emission mechanism. One
possibility is a two-zone model where the spine produces X-rays (Γ ∼50−100) and
the slower sheath produces the synchrotron (radio) emission. The other possibility is
a spine sheath model where the spine produces the radio emission surrounded by a
sheath which is moving faster than the spine. This faster sheath produces X-rays by
a shearing mechanism and requires less extreme bulk Lorentz factor values.
83
4.4
Summary
Despite the fact that there are very few sources in the MCS that have enough
data to create a SED, there is still a lot of valuable information that can be obtained
from the available SEDs. The majority of the SEDs presented in conjunction with the
MCS show that it is possible to model the emission with IC/CMB, without making
any unreasonable assumptions. There are a few sources which show that it is possible
to model jet knots with different emission mechanisms (synchrotron and SSC), but
these models are rare for blazars. The Sambruna et al. [2004] sample has 17 sources
in it and only 6 of them have enough detectable optical data to construct a SED.
Similarly the MCS has only 6 sources with constructed SEDs, but does not have
HST observations for most of the sources. The SED modeling associated with the
MCS further reinforces the selection of the IC/CMB emission model for use with the
sample. In the future it would be ideal to obtain HST observations on all of the
sources in the MCS with X-ray and radio correlations. This would allow for a more
comprehensive study of the possible emission mechanisms that could be associated
with relativistically beamed FR II type blazars and radio galaxies. The new hotspot
SED produced for 3C 111 shows that it is possible to model the radiation as IC/CMB
but a magnetic field which is larger than expected is required. This could mean that
at hotspots the equipartition argument does not apply, or that some other more
complicated model might be needed to describe the hotspot emission mechanisms.
84
5. SUMMARY
The selection criteria that was used to define the MCS sample has increased the
overall detection rate of X-ray jet emission which is correlated with radio jet emission
associated with relativistically beamed FR II sources. The IC/CMB model was chosen
to represent the emission associated with the sample, based on the earlier results of
Marshall et al. [2005] and Sambruna et al. [2004]. The detected X-ray jet emission is
generally well correlated spatially with the radio jet morphology, except for those radio
jets that show extreme bends. The wide range of apparent X-ray to radio ratios along
with the different available SEDs, suggests that no single overall emission model can
completely explain all of the X-ray morphologies. Follow up observations have been
proposed for Chandra and HST [Kharb et al., 2011], which allows for the investigation
of possible synchrotron and IC models for the emission beyond what I have examined
in this thesis.
5.1
Goals and Results
5.1.1
X-ray Detection Rate
The X-ray detection rate for the MCS is ∼ 77.78%, which is a 20% increase from
previous surveys that used radio selected FSRQs to search for X-ray jet emission
[Marshall et al., 2005, Sambruna et al., 2004]. The selection critera that was imposed
on the MOJAVE sample not only picked large (kpc scale) jets, but also selected bright
radio jets. As seen in Figure 3.1, Sext values are a very good predictor of X-ray jet
emission. The threshold for the selection criteria was set to Sext ≥ 100 mJy, but I
found a 100% correlation between the radio and X-ray jet emission when Sext ≥ 300
mJy [Hogan et al., 2011]. Below 300 mJy there is a significant decrease in correlation.
85
5.1.2
IC/CMB Model
The IC/CMB model produces reasonable values of Γ, θ, and δ for the MCS in
most sources, when deceleration and jet bending are ignored. This major assumption
in this model is that the βapp values are the same on the pc and kpc scales. There
are however a couple of sources which have abnormally large values attributed to
Γ. As seen in Appendix C these large Γ values can be rectified by considering jet
bending and deceleration, as neither one by itself will completely solve this problem.
When jet bending alone is considered, the value for Γ is not decreased at all, and
when deceleration is considered the jets end up having Γ values far below 10, which
conflicts with the common assumption in most other IC/CMB models. If both jet
bending and deceleration are considered in combination, then the Γ value can take
on more reasonable values.
5.1.3
Misalignment Angles
Most sources in the MCS have apparent misalignment angles between the pc and
kpc jets which are less than 60◦ . These are easily described by Equations 3.10 and
3.11 for reasonable values of inner jet viewing angles (θn ) and intrinsic bend angle (ζ)
[Conway & Murphy, 1993]. There are however three sources which have misalignment
angles between the pc and kpc scales that are larger than 90◦ . These η values can
be rectified if θn ≤ ζ (assuming a simple bend). When θn ≪ ζ the source roughly
an equal chance of having any misalignment between 0◦ and 180◦ , if the azimuthal
angle is treated as an unknown free parameter. In these cases, the large misalignment
values could be used to constrain the allowable azimuthal angle values, which implies
there is a very specific orientation associated with these extremely misaligned sources.
86
5.1.4
Spectral Energy Distributions
An important way to check on the validity of any emission mechanism assumption
is to study the SED of an object. I have chosen to model the MCS jets with IC/CMB,
as it is often associated with beamed emission. For the most part, the sources which
have optical, radio and X-ray data available show that an SED can be produced
which approximates the IC/CMB model for the X-ray and γ-ray emission. There are
only 6 sources in the sample which have enough information available at this time to
produce an adequate SED, and two of them show emission that could be attributed
to other models (synchrotron and SSC). There are also sources which have X-ray,
radio, and optical maps available, but the optical images do not show any emission
associated with the knot feature above the background level. The 3C 111 (0415+379)
SED shows that the emission for the X-ray portion of the spectrum associated with
the Eastern hotspot can be modeled by IC/CMB emission. The magnetic field for this
source is roughly the same order of magnitude as what is expected. This difference
associated with the magnetic field could be attributed to the low δ assumption or
perhaps the magnetic field decreases as the jet moves away from the core.
5.2
Future Work
5.2.1
Expanding the MCS
An increase in sample size could lead to a better understanding of properties
associated with relativistically beamed jets that have small angles to the line of sight.
To increase the statistics of the MCS, I would need to obtain a larger sample of
relativistically beamed FR II blazars. Since the MOJAVE and VLA samples focus
on sources with a declination greater than zero, it would be possible to obtain new
sources from telescopes located in the southern hemisphere. The ATCA could provide
the kpc scale images with a resolution that is similar to Chandra and the VLA.
Marshall et al. [2005] published radio and X-ray information on sources, which were
87
observed by the ATCA, that could possibly be considered for an extension of the
MCS. To further study the IC/CMB assumptions that were made, I would also need
the pc scale kinematic information for the new sources. Since these sources can only
be detected in the southern hemisphere, I could have their kinematic information
obtained by using the EVLBI Network, which has VLBI telescopes located in Europe,
South Africa, and Asia. The combination of new data sets with the MCS would
increase the sample statistics and perhaps give more insight into the validity of the
IC/CMB model when applied to a sample of relativistically beamed FR II objects with
small viewing angles. The TANAMI (Tracking Active Galactic Nuclei with Austral
Milliarcsecond Interferometry) is a sample of 43 sources that uses VLBI to compile
kinematic information for radio sources [Müller et al., 2010]. The TANAMI sample
was initially created from samples of radio and γ-ray samples, and would provide a
good starting place for selecting new sources to expand the MCS [Müller et al. 2010
and references within].
5.2.2
Deeper X-ray Observations of MCS Sources
Kharb et al. [2011] has already obtained deeper Chandra observations on two
sources associated with the MCS (0106+013 and 1641+399). The extended observations have allowed for a more comprehensive picture of the X-ray jets in these sources
to be constructed. There are also two sources in the MCS which are considered either marginally detected or marginally not detected (2201+315 and 2345−167). If
longer Chandra observation times were procured for these sources it would be much
easier to quantify the detection statistic. Longer detections also benefit sources such
as 0415+379, which have SEDs constructed for them, but no limits for the slope of
the X-ray portion of the SED. If there is a small amount of counts in a region (≤40
counts) then it is difficult to produce an X-ray spectra which confines the slope of the
IC/CMB curve Sambruna et al. [2004].
88
5.2.3
Optical Observations MCS Sources
Kharb et al. [2011] has also obtained HST data for 0106+013 and 1641+399. Optical data is crucial in the production of SEDs, which can confirm that IC/CMB is the
emission model for superluminal FR II blazars. The two SEDs that were produced by
Kharb et al. [2011] imply that IC/CMB can adequately model the emission associated
with these sources, but require a smaller magnetic field value than is expected with
the equipartition assumption. Optical data at different wavelengths can constrain
the SEDs further by refining the location of the well, which is often located between
the synchrotron peak and the IC/CMB or SSC peak on the SED plots. More optical
data are needed for the entire sample, as only 6 (∼22%) of the sources have had SEDs
produced that include points from the optical, radio, and X-ray regimes. There are,
however sources which have optical maps associated with them that show no optical
emission associated with the X-ray and radio emission above the background level
(e.g. 1510+089 from Sambruna et al. [2004]).
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APPENDICES
92
Appendix A: Radio Profiles
Below are the radio profiles for sources in the MCS. The thin solid lines give the
radio profiles along the position angle of the jets. The dashed lines indicate the radio
profile at a position angle of 90◦ counter-clockwise from the jet to avoid any nonjet emission and counter jet emission. The solid, bold line indicates the difference
between the two profiles so that core emission is removed and the effective flux can be
measured. The horizontal dot-dashed lines are set to a value five times the average
noise level and the vertical dashed lines show the inner and outer radius limits. Only
the radiation between the vertical dashed lines is considered for IC/CMB calculation
purposes. This effectively removes all of the emission from the core.
93
Figure A.1. Radio Profiles
94
Figure A.2. Radio Profiles Cont.
95
Figure A.3. Radio Profiles Cont.
96
Figure A.4. Radio Profiles Cont.
97
Figure A.5. Radio Profiles Cont.
98
Appendix B: X-ray Profiles
Below are the X-ray profiles for the sources in the MCS. These are represented
as histograms of the counts in 0.2′′ bins. The solid lines give the profile along the
position angle of the jet, as defined by the radio images. The dashed lines show the
profile along the counter-jet direction, which is defined as 180◦ opposite to the jet.
Figure B.1. X-ray Profiles
99
Figure B.2. X-ray Profiles Cont.
100
Figure B.3. X-ray Profiles Cont.
101
Figure B.4. X-ray Profiles Cont.
102
Figure B.5. X-ray Profiles Cont.
103
Appendix C: Bulk Lorentz Factor vs. Viewing Angle
The Γ-θ plots presented below are a way to combine the pc scale kinematic information with the kpc scale X-ray emission information. Equation 3.4 is plotted as the
blue curve, and Equation 3.6 is plotted as the black curve, where β is a function of
Γ and µ is a function of θ (§ 1.3.2 & § 3.3). The error in each curve is represented
by the dotted lines which flank the original curves. The pc and kpc curves intersect
at a point, which represents a singular pair of Γ and θ values, where the assumption
that there is no jet bending and no jet deceleration has been made. This produces
large values for Γ in some of the more extreme sources. To rectify this problem the
no deceleration assumption was relaxed. This allows for the jet to obtain a second
set of values for Γ and θ and is represented by the red dot-dashed curve. The error
on this red curve is represented by the cyan shaded region. These Γ values were often
very small and of the order of 1 instead of 10. If both assumptions are relaxed, then
the jet can lie anywhere on either curve. This allows for a jet to have a large Γ value
on the pc scale and a more reasonable value (Γ ∼ 10) on the kpc scale.
104
Figure C.1. Bulk Lorentz Factor vs. Viewing Angle
105
Figure C.2. Bulk Lorentz Factor vs. Viewing Angle Cont.
106
Figure C.3. Bulk Lorentz Factor vs. Viewing Angle Cont.
107
Figure C.4. Bulk Lorentz Factor vs. Viewing Angle Cont.
108
Appendix D: Spectral Energy Distributions
Figure D.1. Spectral Energy Distribution for the knot associated with the primary jet
in 1641+399 [Sambruna et al., 2004]. This SED was created using the data for the
SED modeling parameters in Sambruna et al. [2004]. An additional point associated
with the 1.4 GHz VLA data was added to further constrain the SED. The solid line
represents the synchrotron radiation, while the dot dashed and dashed lines represent
the IC/CMB and SSC emission respectively.
109
Figure D.2. Spectral Energy Distribution for the knots associated with the primary
jet in 3C 273 (1226+023) [Jester et al., 2006]
110
Figure D.3. Spectral Energy Distribution for the knots associated with the primary
jet in 3C 273 (1226+023) [Sambruna et al., 2001]
111
Figure D.4. Spectral Energy Distribution for the knots associated with the primary
jet in 3C 273 (1226+023) [Marshall et al., 2001]
112
Figure D.5. Spectral Energy Distribution for the knot associated with the primary jet
in 1222+216 [Jorstad & Marscher, 2006].This SED is non-unique and needs a optical
data point to further constrain the Synchrotron radiation curve.
113
Figure D.6. Spectral Energy Distribution for the primary jet in 3C 279 (1253-055)
[Figure 8 from Collmar et al. 2010]. The leptonic one-zone jet model that was used
fits only the near-infrared to γ-ray emission and is believed to be produced by pc
scale jet [Collmar et al., 2010]. The Chandra portion of the SED is derived from the
kpc scale emission.
114
Figure D.7. Spectral Energy Distribution for the primary jet in 1928+738 [Sambruna
et al., 2004]. This SED was created using the data for the SED modeling parameters
in Sambruna et al. [2004]. The solid line represents the synchrotron radiation, while
the dot dashed and dashed lines represent the IC/CMB and SSC emission respectively.
VITA
115
VITA
Name: Brandon Scott Hogan
Place of Birth: Indianapolis, Indiana U.S.A
Date of Birth: 31 March 1982
Educational Institutions Attended
• Purdue University, West Lafayette, IN, 2000−Present
Degrees Awarded
• B.S. in Applied Physics with Minors in Philosophy, Mathematics, and Biology,
Purdue University, 2005
Publications
• ”‘Chandra Discovery of 10 New X-Ray Jets Associated With FR II Radio CoreSelected AGNs in the MOJAVE Sample”’ Hogan, B., Lister, M., Kharb, P.,
Marshall, H., & Cooper, N. 2011, arXiv:1101.5342
• ”‘Chandra and HST observations of two Superluminal Blazars: 0106+013 &
1641+399”’ Kharb, P., Lister, M., Marshall, H., & Hogan., B., 2011 In Prep.
Presentations
• ”‘X-ray Jets in Superluminal Blazars”’ Hogan, B. S., Lister, M., Marshall, H., &
Kharb, P. 2009, American Astronomical Society Meeting Abstracts #213, 213,
#608.07
Honors & Awards
• AAPT Outstanding Graduate Teaching Assistant Award
• Outstanding Sophomore in Physics Award
116
Societies
• Society of Physics Students (2000-2005)
• American Astronomical Society