THE MOJAVE CHANDRA SAMPLE: A CORRELATION STUDY OF ...

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

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[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.]

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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.]

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TABLE OF CONTENTS Page LIST OF TABLES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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LIST OF FIGURES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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SYMBOLS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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ABBREVIATIONS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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ABSTRACT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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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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1 1 4 5 6 8 8 10 12 13 13 14 15 16 18

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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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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73 73 75 76 76 77 77 78 80 83

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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LIST OF REFERENCES . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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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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VITA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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LIST OF TABLES Table

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2.1

MOJAVE CHANDRA SAMPLE . . . . . . . . . . . . . . . . . . . . .

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2.2

OBSERVATION LOG . . . . . . . . . . . . . . . . . . . . . . . . . . .

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2.3

MOJAVE CHANDRA SAMPLE JET MEASUREMENTS . . . . . . .

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2.4

VLA ARCHIVAL DATA . . . . . . . . . . . . . . . . . . . . . . . . . .

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3.1

MOJAVE CHANDRA SAMPLE BEAMING MODEL PARAMETERS

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4.1

SED PARAMETERS . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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4.2

3C 111 SED INFORMATION . . . . . . . . . . . . . . . . . . . . . . .

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LIST OF FIGURES Figure

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1.1

Visual representation of a radio-loud AGN [Urry & Padovani, 1995] . .

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1.2

AGN Taxonomy [Urry & Padovani, 1995] . . . . . . . . . . . . . . . . .

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1.3

The FR I/FR II Divide [Ghisellini et al., 1993] . . . . . . . . . . . . . .

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1.4

Visual representation of superluminal motion as seen in Ghisellini [2000].

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1.5

Relativistic beaming of radiation which is emitted isotropically in the rest frame K′ (S′ in the text) [Rybicki & Lightman, 1979]. . . . . . . . . . .

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1.6

Inverse-Compton Scattering . . . . . . . . . . . . . . . . . . . . . . . .

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1.7

Chandra X-ray Observatory (Courtesy of NASA/CXC/NGST) . . . . .

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2.1

Radio/X-ray overlay of 0106+013. . . . . . . . . . . . . . . . . . . . . .

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2.7

Radio/X-ray overlay of 0605-085. . . . . . . . . . . . . . . . . . . . . .

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2.8

Radio/X-ray overlay of 1045-188. . . . . . . . . . . . . . . . . . . . . .

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2.10 Radio/X-ray overlay of 1156+295. . . . . . . . . . . . . . . . . . . . . .

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2.13 Radio/X-ray overlay of 1253-055. . . . . . . . . . . . . . . . . . . . . .

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3.1

Histogram relating the source population to the Sext value . . . . . . .

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3.2

Histogram representing the redshift distribution of the MOJAVE sample

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3.3

Histogram representing the redshift distribution of the MCS sample . .

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3.4

Position Angle Misalignment Associated with the MCS . . . . . . . . .

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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]

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4.2

A.1 Radio Profiles . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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A.2 Radio Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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B.1 X-ray Profiles . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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B.2 X-ray Profiles Cont. . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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C.1 Bulk Lorentz Factor vs. Viewing Angle . . . . . . . . . . . . . . . . . .

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C.2 Bulk Lorentz Factor vs. Viewing Angle Cont. . . . . . . . . . . . . . .

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D.1 Spectral Energy Distribution for the knot associated with the primary jet in 1641+399 [Sambruna et al., 2004] . . . . . . . . . . . . . . . . . . .

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D.2 Spectral Energy Distribution for the knots associated with the primary jet in 3C 273 (1226+023) [Jester et al., 2006] . . . . . . . . . . . . . . . .

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D.3 Spectral Energy Distribution for the knots associated with the primary jet in 3C 273 (1226+023) [Sambruna et al., 2001] . . . . . . . . . . . . . .

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D.4 Spectral Energy Distribution for the knots associated with the primary jet in 3C 273 (1226+023) [Marshall et al., 2001] . . . . . . . . . . . . . . .

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D.5 Spectral Energy Distribution for the knot associated with the primary jet in 1222+216 [Jorstad & Marscher, 2006] . . . . . . . . . . . . . . . . .

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D.6 Spectral Energy Distribution for the primary jet in 3C 279 (1253-055) [Collmar et al., 2010] . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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D.7 Spectral Energy Distribution for the primary jet in 1928+738 [Sambruna et al., 2004] . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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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

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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

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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.

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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

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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.

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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


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


2141

0.593

1.48

19.3 ± 1.2


2143

0.621

0.20

14.4 ± 1.4


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


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


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


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

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.

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 A−1 . This information along with quantity, produces a flux in terms of erg cm−2 s−1 ˚ 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 [M¨ uller 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 [M¨ uller 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


96


97


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


101


102


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.

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Figure C.1. Bulk Lorentz Factor vs. Viewing Angle

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Figure C.2. Bulk Lorentz Factor vs. Viewing Angle Cont.

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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.

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Figure D.2. Spectral Energy Distribution for the knots associated with the primary jet in 3C 273 (1226+023) [Jester et al., 2006]

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Figure D.3. Spectral Energy Distribution for the knots associated with the primary jet in 3C 273 (1226+023) [Sambruna et al., 2001]

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Figure D.4. Spectral Energy Distribution for the knots associated with the primary jet in 3C 273 (1226+023) [Marshall et al., 2001]

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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.

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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.

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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

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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