THE MOJAVE CHANDRA SAMPLE: A CORRELATION STUDY OF BLAZARS AND RADIO GALAXIES IN X-RAY AND RADIO WAVELENGTHS. A Dissertation. Submitted to the Faculty

THE MOJAVE CHANDRA SAMPLE: A CORRELATION STUDY OF BLAZARS AND RADIO GALAXIES IN X-RAY AND RADIO WAVELENGTHS A Dissertation Submitted to the Faculty of Purdue University by Brandon S. Hogan In Partial Fulfillment of the Requirements for the Degree of Doctor of Philosophy May 2011 Purdue University West Lafayette, Indiana

ii [I dedicate this to my lovely wife, Meredith, my wonderful group of friends, and my supportive family. I could not have done this without the love and support of all of you.]

iii ACKNOWLEDGMENTS [I would like to acknowledge Matthew Lister, Herman Marshall, Nathan Cooper, Preeti Kharb, and Talvikki Hovatta, as they have supported and helped me throughout the duration of this project. This project was funded by grants from NASA and NSF.]

iv TABLE OF CONTENTS LIST OF TABLES................................ LIST OF FIGURES............................... SYMBOLS.................................... ABBREVIATIONS................................ ABSTRACT................................... Page 1 INTRODUCTION.............................. 1 1.1 Active Galactic Nuclei......................... 1 1.1.1 Radio Quiet AGN........................ 4 1.1.2 Radio Loud AGN........................ 5 1.2 The Fanaroff Riley Classification of AGN............... 6 1.3 Relativistic Properties of AGN..................... 8 1.3.1 Apparent Superluminal Motion................ 8 1.3.2 Beaming............................. 10 1.3.3 Inverse-Compton Scattering.................. 12 1.4 Astronomical Instruments used in the MOJAVE Chandra Sample. 13 1.4.1 The Very Large Array..................... 13 1.4.2 The Very Long Baseline Array................. 14 1.4.3 Chandra X-ray Observatory.................. 15 1.5 The Status of X-ray Jet Astrophysics................. 16 1.6 Thesis Description and Outline.................... 18 2 THE MOJAVE CHANDRA SAMPLE................... 19 2.1 Selection Criteria............................ 19 2.2 Individual Source Observations of the MCS.............. 20 2.2.1 0106+013(OC 12)........................ 24 2.2.2 0119+115............................ 26 2.2.3 0224+671 (4C 67.05)...................... 28 2.2.4 0234+285 (CTD 20)...................... 28 2.2.5 0415+379 (3C 111)....................... 29 2.2.6 0529+075 (OG 050)...................... 31 2.2.7 0605-085............................. 32 2.2.8 1045-188............................. 32 2.2.9 1055+018 (4C 01.28)...................... 33 2.2.10 1156+295 (4C 29.45)...................... 34 vii viii xi xii xiv

v Page 2.2.11 1222+216 (4C 21.35)...................... 34 2.2.12 1226+023 (3C 273)....................... 35 2.2.13 1253-055 (3C 279)....................... 38 2.2.14 1334-127............................. 39 2.2.15 1510-089............................. 39 2.2.16 1641+399 (3C 345)....................... 40 2.2.17 1655+077............................ 41 2.2.18 1800+440 (S4 1800-44)..................... 42 2.2.19 1828+487 (3C 380)....................... 44 2.2.20 1849+670 (S4 1849-67)..................... 44 2.2.21 1928+738 (4C 73.18)...................... 45 2.2.22 1957+405 (Cygnus A)..................... 46 2.2.23 2155-152............................. 47 2.2.24 2201+315 (4C 31.63)...................... 48 2.2.25 2216-038............................. 49 2.2.26 2251+158 (3C 454.3)...................... 50 2.2.27 2345-167............................. 51 3 DATA REDUCTION AND ANALYSIS................... 52 3.1 X-ray Radio Overlays.......................... 52 3.2 X-ray and Radio Jet Analysis..................... 54 3.3 The Single Zone IC/CMB Model................... 56 3.4 Scenarios Associated with the IC-CMB Model............ 61 3.4.1 IC/CMB model with No Jet Deceleration or Bending.... 61 3.4.2 IC/CMB model with Jet Deceleration............. 63 3.4.3 IC/CMB model with Deceleration and Jet Bending..... 64 3.5 S ext as an X-ray jet predictor..................... 65 3.6 Kolmogorov-Smirnov Tests....................... 66 3.7 Viewing Angle.............................. 70 4 SPECTRAL ENERGY DISTRIBUTIONS................. 73 4.1 General Information.......................... 73 4.2 3C 111 (0415+379) SED........................ 75 4.2.1 Obtaining the Radio Fluxes.................. 76 4.2.2 Obtaining the Optical Fluxes................. 76 4.2.3 Obtaining the X-ray Fluxes.................. 77 4.2.4 Obtaining the γ-ray Fluxes.................. 77 4.2.5 Uniqueness of the 3C 111 Hotspot SED............ 78 4.3 Individual SED Notes.......................... 80 4.4 Summary................................ 83 5 SUMMARY.................................. 84 5.1 Goals and Results............................ 84 5.1.1 X-ray Detection Rate...................... 84

vi Page 5.1.2 IC/CMB Model......................... 85 5.1.3 Misalignment Angles...................... 85 5.1.4 Spectral Energy Distributions................. 86 5.2 Future Work............................... 86 5.2.1 Expanding the MCS...................... 86 5.2.2 Deeper X-ray Observations of MCS Sources......... 87 5.2.3 Optical Observations MCS Sources.............. 88 LIST OF REFERENCES............................ 89 APPENDICES.................................. 92 Appendix A: Radio Profiles......................... 92 Appendix B: X-ray Profiles......................... 98 Appendix C: Bulk Lorentz Factor vs. Viewing Angle........... 103 Appendix D: Spectral Energy Distributions................ 108 VITA....................................... 115

vii LIST OF TABLES Table Page 2.1 MOJAVE CHANDRA SAMPLE..................... 21 2.2 OBSERVATION LOG........................... 22 2.3 MOJAVE CHANDRA SAMPLE JET MEASUREMENTS....... 23 2.4 VLA ARCHIVAL DATA.......................... 25 3.1 MOJAVE CHANDRA SAMPLE BEAMING MODEL PARAMETERS 57 4.1 SED PARAMETERS............................ 74 4.2 3C 111 SED INFORMATION....................... 79

viii Figure LIST OF FIGURES Page 1.1 Visual representation of a radio-loud AGN [Urry & Padovani, 1995].. 2 1.2 AGN Taxonomy [Urry & Padovani, 1995]................. 4 1.3 The FR I/FR II Divide [Ghisellini et al., 1993].............. 7 1.4 Visual representation of superluminal motion as seen in Ghisellini [2000]. 9 1.5 Relativistic beaming of radiation which is emitted isotropically in the rest frame K (S in the text) [Rybicki & Lightman, 1979]........... 12 1.6 Inverse-Compton Scattering........................ 13 1.7 Chandra X-ray Observatory (Courtesy of NASA/CXC/NGST)..... 15 2.1 Radio/X-ray overlay of 0106+013...................... 24 2.2 Radio/X-ray overlay of 0119+115...................... 27 2.3 Radio/X-ray overlay of 0224+671...................... 28 2.4 Radio/X-ray overlay of 0234+285...................... 29 2.5 Radio/X-ray overlay of 0415+379...................... 30 2.6 Radio/X-ray overlay of 0529+075...................... 31 2.7 Radio/X-ray overlay of 0605-085...................... 32 2.8 Radio/X-ray overlay of 1045-188...................... 33 2.9 Radio/X-ray overlay of 1055+018...................... 34 2.10 Radio/X-ray overlay of 1156+295...................... 35 2.11 Radio/X-ray overlay of 1222+216...................... 36 2.12 Radio/X-ray overlay of 1226+023...................... 37 2.13 Radio/X-ray overlay of 1253-055...................... 38 2.14 Radio/X-ray overlay of 1334-127...................... 39 2.15 Radio/X-ray overlay of 1510-089...................... 40 2.16 Radio/X-ray overlay of 1641+399...................... 41

ix Figure Page 2.17 Radio/X-ray overlay of 1655+077...................... 42 2.18 Radio/X-ray overlay of 1800+440...................... 43 2.19 Radio/X-ray overlay of 1828+487...................... 44 2.20 Radio/X-ray overlay of 1849+670...................... 45 2.21 Radio/X-ray overlay of 1928+738...................... 46 2.22 Radio/X-ray overlay of 1957+405...................... 47 2.23 Radio/X-ray overlay of 2155-152...................... 48 2.24 Radio/X-ray overlay of 2201+315...................... 49 2.25 Radio/X-ray overlay of 2216-038...................... 50 2.26 Radio/X-ray overlay of 2251+158...................... 50 2.27 Radio/X-ray overlay of 2345-167...................... 51 3.1 Histogram relating the source population to the S ext value....... 67 3.2 Histogram representing the redshift distribution of the MOJAVE sample 69 3.3 Histogram representing the redshift distribution of the MCS sample.. 69 3.4 Position Angle Misalignment Associated with the MCS......... 71 4.1 Spectral Energy Distribution for the hotspot associated with the primary jet in 3C 111................................ 74 4.2 Radio, optical, and X-ray jets associated with 3C 273 [Jester et al., 2006] 81 A.1 Radio Profiles................................ 93 A.2 Radio Profiles Cont............................. 94 A.3 Radio Profiles Cont............................. 95 A.4 Radio Profiles Cont............................. 96 A.5 Radio Profiles Cont............................. 97 B.1 X-ray Profiles................................ 98 B.2 X-ray Profiles Cont............................. 99 B.3 X-ray Profiles Cont............................. 100 B.4 X-ray Profiles Cont............................. 101 B.5 X-ray Profiles Cont............................. 102

x Figure Page C.1 Bulk Lorentz Factor vs. Viewing Angle.................. 104 C.2 Bulk Lorentz Factor vs. Viewing Angle Cont............... 105 C.3 Bulk Lorentz Factor vs. Viewing Angle Cont............... 106 C.4 Bulk Lorentz Factor vs. Viewing Angle Cont............... 107 D.1 Spectral Energy Distribution for the knot associated with the primary jet in 1641+399 [Sambruna et al., 2004]................... 108 D.2 Spectral Energy Distribution for the knots associated with the primary jet in 3C 273 (1226+023) [Jester et al., 2006]................ 109 D.3 Spectral Energy Distribution for the knots associated with the primary jet in 3C 273 (1226+023) [Sambruna et al., 2001].............. 110 D.4 Spectral Energy Distribution for the knots associated with the primary jet in 3C 273 (1226+023) [Marshall et al., 2001]............... 111 D.5 Spectral Energy Distribution for the knot associated with the primary jet in 1222+216 [Jorstad & Marscher, 2006]................. 112 D.6 Spectral Energy Distribution for the primary jet in 3C 279 (1253-055) [Collmar et al., 2010]............................ 113 D.7 Spectral Energy Distribution for the primary jet in 1928+738 [Sambruna et al., 2004]................................. 114

xi SYMBOLS Γ Bulk Lorentz Factor δ Doppler Factor β app v app c φ k V L C B 1 R ν r ν x S r S x K θ Apparent Superluminal Speed Apparent Superluminal Velocity Speed of Light Filling Factor Baryon Fraction Energy Parameter Emitting Volume Observed Synchrotron Luminosity Weak Function of the Low Frequency Spectral Index of the Synchrotron Spectrum Non-Boosted Spatially Averaged, Minimum Energy Magnetic Field of the Jet X-ray to Radio Luminosity Ratio Radio Frequency X-ray Frequency Radio Flux Density X-ray Flux Density K Factor Angle With Respect to the Line of Sight µ Cosine of θ

xii ABBREVIATIONS AGN VLA VLBA VLBI HST MOJAVE MCS FSRQ BL Lac SSRQ NLRG BLRG NELG QSO BLR NLR BAL FR I FR II CDQ LDQ IGM CMB IC NGST Active Galactic Nuclei Very Large Array Very Large Baseline Array Very Large Baseline Interferometry Hubble Space Telescope Monitoring Of Jets in AGN with VLBA Experiments MOJAVE Chandra Sample Flat Spectrum Radio Quasar BL Lacertae object Steep Spectrum Radio Quasar Narrow Line Radio Galaxies Broad Line Radio Galaxies Narrow Emission Line Galaxies Quasi Stellar Object (Quasar) Broad Line Region Narrow Line Region Broad Absorption Line (Quasar) Fanaroff & Riley type I object Fanaroff & Riley type II object Core Dominated Quasar Lobe Dominated Quasar Inter-Galactic Medium Cosmic Microwave Background Inverse Compton Nortrop Grumman Space Technology

xiii ISIM HRC CCD ACIS SAO MIT IC/CMB NASA SED FWHM WCS Integrated Science Instrument Module High Resolution Camera Charge Collecting Device Advanced CCD Imaging Spectrometer Smithsonian Astrophysical Observatory Massachusetts Institute of Technology inverse Compton scattering off of cosmic microwave background National Aeronautics and Space Administration Spectral Energy Distribution Full Width Half Maximum World Coordinate System

xiv ABSTRACT Hogan, Brandon S. Ph.D., Purdue University, May 2011. The MOJAVE Chandra Sample: A Correlation Study of Blazars and Radio Galaxies in X-ray and Radio Wavelengths. Major Professor: Matthew L. Lister. The Chandra X-ray observatory has increased the quality and number of detections in the X-ray regime since its launch in 1999. It is an imporant tool for studying the jets which are associated with Active Galacitc Nuclei (AGN) and their possible emission mechanisms. The MOJAVE Chandra Sample (MCS) is a sample of 27 AGN which have been selected from the radio flux-limited MOJAVE (Monitoring of Jets in AGN with VLBA Experiments) sample. The objects contained in the MOJAVE sample are traditionally associated with relativistically beamed jets that have small viewing angles. The MCS was created to study the correlation of X-ray and radio emission on kiloparsec scales. The complete sample is made up of all MOJAVE Fanaroff & Riley type II objects which have over 100 mjy of extended radio emission at 1.4 GHz and a radio structure of at least 3 in extent. Chandra observations have revealed X-ray and radio correlation in 21 of the 27 jets, bringing the detection rate to 78%. The selection criteria provides a quantitative method of discovering new X-ray jets associated with AGN from radio observations. The X-ray morphologies are usually well correlated with the radio emission, except for the sources which show extreme bending on the kiloparsec scale. The emission mechanism for these relativisiticly beamed quasars and radio galaxies can be interpreted as inverse Compton scattering off of the consmic microwave background by the electrons in the jets (IC/CMB). The emission mechanism is reinforced by spectral energy distributions (SED) which model the emission mechanisms for sources with sufficient X-ray, optical, and radio data available. I have explored the effects of jet bending and jet deceleration in conjunction with the inverse Compton emission model and used dif-

xv ferent scenarios to derive best fit viewing angles and bulk Lorentz factors, which were calculated by using the superluminal speeds along with parameters that were derived from the IC/CMB model. The range of viewing angles and Lorentz factors are examined for each scenario, as well as their implications for the other parameters associated with models. To achieve results that are consistant with other models jet bending and deceleration must be considered with the IC/CMB model.

1 1. INTRODUCTION The overall goal of this dissertation is to investigate and further understand how jet emission from Active Galactic Nuclei (AGN) correlates between the radio and X-ray regimes. In this dissertation I describe how one can use the data associated with the MOJAVE (Monitoring Of Jets in AGN with VLBA Experiments) sample along with additional selection criteria to determine if X-ray jet detections are probable with the Chandra X-ray Observatory. I also study the implications of this correlation on the overall inverse-compton emission mechanism which I have chosen for modeling the X-ray emission of these objects. The introductory sections below provide the background information necessary to form a foundation for the understanding of the material provided in this dissertation. 1.1 Active Galactic Nuclei An AGN is traditionally regarded as an accreting supermassive black hole, which is located at the center of a galaxy, and has a mass on the order of 10 6 M to 10 10 M, where M is one solar mass. This black hole is surrounded by an accretion disk, which is made up of material that spirals inward toward the black hole, and encompasses a flat circular region which is perpendicular to the rotation poles and/or jets. As seen in Figure 1.1, a typical radio-loud AGN is comprised of a black hole, an accretion disk, a torus, and a jet, where the narrow line regions are found further away from the core than the broad line regions. The narrow line and broad line regions are where the narrow and broad emission lines are produced. Jets are hypothesized to have been produced by a phenomenon known as magnetic launching, which is described by the rotation of a black hole and accretion disk system

2 Figure 1.1. Visual representation of a radio-loud AGN [Urry & Padovani, 1995] [Marscher, 2009]. This leads to the magnetic field lines twisting up into a helical structure, which then causes a pressure gradient to occur, accelerating the plasma flow downstream. This helical magnetic field is often thought of as the confinement structure for the jet. Bridle & Perley [1984] define extragalactic radio jets by three criteria. The jet must be four times as long as it is wide. It must be separable from other extended structures at high resolution.

3 It must be aligned with the compact radio core that it protrudes from. Extragalactic jets, which are comprised of highly energetic plasma, often appear to be moving near the speed of light, or in some cases, even faster than the speed of light. This is known as apparent superluminal motion and is described in further detail in 1.3.1. There are traditionally two ways that a jet can terminate. The jet either dissipates enough energy into the environment around it, such that it fades slowly in a plume-like structure, or it abruptly terminates at a shock front known as a terminal hotspot. Hotspots are often enveloped in large regions of radiation which gravitate backwards toward the core, known as lobes. These lobes are often seen at the end of both jets, even though the jet which is traveling away from the observer can sometimes not be seen due to relativistic beaming effects (see 1.3.2). AGN are traditionally separated into two groups; radio-loud and radio-quiet AGN. Radio-loud AGN make up about 15% to 20% of the total AGN population [Urry & Padovani, 1995]. These groups are defined by their ratios of 5 GHz radio flux to optical (B-band) flux. If the ratio of radio flux to optical flux is greater than 10 then the object is considered to be radio loud [Kellermann et al., 1989]. The radioloud classification typically includes the Narrow Line Radio Galaxies (NLRG), Broad Line Radio Galaxies (BLRG), Steep Spectrum Radio Quasars (SSRQ), Flat Spectrum Radio Quasars (FSRQ), and Blazars (FSRQs and BL Lacertae objects), and the radio-quiet classification includes Seyfert Galaxies (type 1 & 2), Narrow Emission Line (X-ray) Galaxies (NELG), Infrared Quasars, and radio quiet Quasars (Quasi Steller Objects, QSO). A visual description of the relation of these items to radio loudness, angle to the line of sight, black hole spin, and optical emission lines is found in Figure 1.2.

4 Figure 1.2. AGN Taxonomy [Urry & Padovani, 1995] 1.1.1 Radio Quiet AGN Seyfert Galaxies Seyfert galaxies have the lowest luminosities of all of the sources in the radio quiet regime, and are usually located much nearer to us than the more powerful AGN. They are sub-divided into two types based on the emission lines that they produce. Type 1 Seyfert galaxies exhibit emission lines from the Broad Line Region (BLR) and the Narrow Line Region (NLR), while their Type 2 counterparts exhibit emission lines from only the NLR. The BLR emission lines are produced close to the core, presumably from the interaction between the emission from the core/jet nozzle and the clouds above the accretion disk, or perhaps by the disk itself [line width 10000 km/sec, Antonucci 1993]. Thus, orientation could cause the difference in the two types of Seyfert galaxies, because at smaller angles the emission does not have to travel through the dusty torus (i.e., The rotation axis of Type 1 Seyfert galaxies has a smaller angle to the line of sight than the Type 2 Seyferts). The narrow emission

5 lines are produced from clouds which are downstream from the nucleus [line width 1000 km/sec, Antonucci 1993, Urry & Padovani 1995]. Radio-Quiet Quasars A radio quiet quasar can show both broad and narrow absorption lines like its Seyfert Type 1 counterpart, but is distinguished by its larger luminosity [Urry & Padovani, 1995]. Broad Absorption Line Quasars (BALs) make up about 10% of the population of radio quiet quasars. Interestingly, the dust clouds which are thought to be the cause of BALs cover about 10% of the source, leading to the presumption that all radio quiet quasars have these clouds around them [Antonucci 1993 and references within]. The polar axis (or jet) viewing angle might account for the difference between BAL quasars and the rest of the radio quiet quasars, as seen in Figure 1.2. 1.1.2 Radio Loud AGN Blazars Blazars, like most other radio-loud AGN, are described by very powerful jets which are generated in AGN as a result of accretion onto supermassive black holes. Blazar jets can transport energy over large distances using highly energetic plasma as a medium. These energetic outflows are usually oriented at very small angles with respect to the observer s line of sight (θ 15 ) and tend to show apparent superluminal velocities [Angel & Stockman, 1980]. The blazar class encompasses two groups of objects, the flat spectrum radio quasars (FSRQs) and BL Lacertae (BL Lac) objects. The FSRQ objects are thought to have more powerful, well collimated jets that terminate at large shock fronts known as hotspots, and are also described as Fanaroff-Riley type II jets (FRII; Fanaroff & Riley 1974). On the other hand, the BL Lac objects are described as Fanaroff-Riley type I objects, are less powerful than FR IIs, and tend to dissipate more energy into the intergalactic medium before

6 they terminate in a plume like structure [Urry & Padovani, 1995]. The Fanaroff- Riley 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 10 25 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 (β app ). 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 t e 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 t e where β is just the velocity of the object divided by the speed of light (c) and is described by Equation 1.1 below. β = v c. (1.1) Thus, the distance between points A and C is βc t e 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 t e sin θ. c t e 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 t a = t e (1 β cos θ). The apparent speed is found by dividing the apparent velocity (v app ) 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) β app = v app c = βc t e sin θ c t a = β sin θ (1 β cos θ). (1.2) Equation 1.2 can produce values for β app 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 where the Lorentz factor (Γ) is u = u Γ(1 + vu /c2 ), (1.3) Γ = 1 1 β 2. (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. u = u + v (1 + vu /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) cosθ (1.7) The Doppler factor can be introduced by rewriting Equation 1.5 as Thus, the Doppler factor is sin θ = δ sin θ (1.8) δ = where the inverse transformation (from S to S ) for δ is 1 Γ(1 β cos θ), (1.9) δ = 1 Γ(1 + β cos θ ). (1.10) 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 cosθ 0 leaving the sinθ 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 Array 1 (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 Array 2 (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 website 4. 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 (S 1.4GHz 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 (S 5GHz,V LA > 1 Jy and S 2.7GHz,ATCA >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 (S ext ) 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 H 0 = 71 km s 1 Mpc 1, Ω m = 0.27, and Ω Λ = 0.73.

19 2.1 Selection Criteria 2. THE MOJAVE CHANDRA SAMPLE 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 S ext β app Reference Obs ID (1) (2) (3) (4) (5) (6) (7) 0106+013 OC 12 2.099 0.53 26.5 ± 4.2 Hogan et al. [2011] 9281 0119+115 0.57 0.11 17.1 ± 1.1 Hogan et al. [2011] 9290 0224+671 4C 67.05 0.523 0.15 11.6 ± 0.8 Hogan et al. [2011] 9288 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 0605 085 0.872 0.12 19.8 ± 1.2 Sambruna et al. [2004] 2132 1045 188 0.595 0.51 8.6 ± 0.8 Hogan et al. [2011] 9280 1055+018 4C 01.28 0.89 0.23 11.0 ± 1.2 Sambruna et al. [2004] 2137 1156+295 4C 29.45 0.729 0.20 24.9 ± 2.3 Coppi et al. [2002] 0874 1222+216 4C 21.35 0.432 0.96 21.0 ± 2.2 Jorstad & Marscher [2006] 3049 1226+023 3C 273 0.158 17.67 13.4 ± 0.8 Jester et al. [2006] 4879 1253 055 3C 279 0.536 2.10 20.6 ± 1.4 WEBT [2007] 6867 1334 127 0.539 0.15 10.3 ± 1.1 Hogan et al. [2011] 9282 1510 089 0.36 0.18 20.2 ± 4.9 Sambruna et al. [2004] 2141 1641+399 3C 345 0.593 1.48 19.3 ± 1.2 Sambruna et al. [2004] 2143 1655+077 0.621 0.20 14.4 ± 1.4 Marshall et al. [2005] 3122 1800+440 S4 1800 44 0.663 0.25 15.4 ± 1.0 Hogan et al. [2011] 9286 1828+487 3C 380 0.692 5.43 13.7 ± 0.8 Marshall et al. [2005] 3124 1849+670 S4 1849 67 0.657 0.10 30.6 ± 2.2 Hogan et al. [2011] 9291 1928+738 4C 73.18 0.302 0.36 8.4 ± 0.6 Sambruna et al. [2004] 2145 1957+405 Cygnus A 0.0561 414.18 0.2 ± 0.1 Wilson et al. [2001] 1707 2155 152 0.672 0.30 18.1 ± 2.0 Hogan et al. [2011] 9284 2201+315 4C 31.63 0.295 0.37 7.9 ± 0.6 Hogan et al. [2011] 9283 2216 038 0.901 0.31 5.6 ± 0.6 Hogan et al. [2011] 9285 2251+158 3C 454.3 0.859 0.88 14.2 ± 1.1 Marshall et al. [2005] 3127 2345 167 0.576 0.14 13.5 ± 1.1 Hogan et al. [2011] 9328 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 PA pc PA kpc R i R o S r ν r Count Rate S x P jet 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 0119+115 6 35 1.5 8.0 22.2 ± 0.3 1.40 0.00 ± 0.38 < 1.2 5.54 10 1 N 0224+671-5 -10 1.5 11.0 22.9 ± 0.7 1.40-0.55 ± 0.45 < 0.8 9.62 10 1 N 0234+285-12 0 1.5 6.0 53.9 ± 0.4 1.40 6.09 ± 1.09 6.1 < 1 10 10 Y 0415+379 65 63 1.5 100.0 50.6 ± 7.1 1.44 7.50 ± 2.49 7.5 2.44 10 6 Y 0529+075-31 -145 1.5 8.0 69.2 ± 0.3 1.40 1.52 ± 0.65 1.5 1.99 10 4 Y 0605 085 117 100 1.5 5.2 93.6 ± 1.1 1.42 10.85 ± 1.21 10.9 < 1 10 10 Y 1045 188 149 125 1.5 10.0 167.7 ± 4.9 1.42 2.82 ± 0.82 2.8 4.74 10 8 Y 1055+018-49 180 1.5 13.0 74.3 ± 1.8 1.42-0.86 ± 0.97 < 2.1 9.01 10 1 N 1156+295-2 -10 1.5 3.5 76.6 ± 0.7 1.52-0.66 ± 1.15 < 2.8 8.65 10 1 N 1222+216-3 35 1.5 4.0 81.0 ± 0.3 1.40 8.22 ± 0.81 8.2 < 1 10 10 Y 1226+023-123 -135 1.5 20.0 4603.9 ± 11.8 1.45 115.50 ± 2.83 115.5 < 1 10 10 Y 1253 055-124 -150 1.5 5.5 790.5 ± 2.2 1.66 3.91 ± 0.51 3.9 < 1 10 10 Y 1334 127 150 135 1.5 12.0 103.8 ± 0.2 1.49 17.07 ± 1.56 17.1 < 1 10 10 Y 1510 089-31 165 1.5 6.0 63.3 ± 0.7 1.46 14.93 ± 1.63 14.9 < 1 10 10 Y 1641+399-89 -30 1.5 4.0 96.7 ± 1.8 1.51 3.86 ± 1.14 3.9 < 1 10 10 Y 1655+077-38 -50 1.5 5.0 46.7 ± 0.3 1.55-0.83 ± 0.72 < 1.3 9.58 10 1 N 1800+440-159 -130 1.5 8.0 133.2 ± 0.5 1.51 6.28 ± 0.99 6.3 < 1 10 10 Y 1828+487-39 -40 1.5 3.0 28.4 ± 0.8 1.55 7.13 ± 1.43 7.1 < 1 10 10 Y 1849+670-45 0 5.0 20.0 8.3 ± 0.7 1.40 1.08 ± 0.40 1.1 1.36 10 6 Y 1928+738 162 180 1.5 11.0 33.0 ± 1.3 1.42 5.24 ± 1.39 5.2 < 1 10 10 Y 1957+405-79 -75 1.5 65.0 51381.4 ± 62.8 1.52-10.40 ± 3.71 < 0.7 1 10 0 N* 2155 152-148 -170 1.5 12.0 231.3 ± 0.8 1.40 1.51 ± 0.85 1.5 5.00 10 3 Y 2201+315-142 -110 1.5 10.0 31.1 ± 0.7 1.40 1.96 ± 1.05 2.0 1.54 10 3 Y 2216 038-172 135 1.5 15.5 164.2 ± 0.9 1.40 1.74 ± 0.78 1.7 4.94 10 4 Y 2251+158-76 -50 1.5 5.5 585.0 ± 3.5 1.50 14.61 ± 2.74 14.6 < 1 10 10 Y 2345 167 124-135 2.0 8.0 83.7 ± 0.4 1.40 0.65 ± 0.67 < 2.7 8.92 10 2 N Note. Columns are as follows: (1) IAU name (B1950.0) (2) Position angle of the pc-scale radio jet in ( ). All position angles are measured from north through east (3) Position angle of the kpc-scale radio jet ( ) (4) Inner radius in ( ) (5) Outer radius in ( ) (6) Observed flux density of the radio jet in mjy. The jet radio flux density is computed at the given radii or the same region as for the X-ray count rate, given by the PA kpc, R i, and R o parameters (7) Observation frequency of the radio image in GHz (8) Counts per kilosecond (9) The X-ray flux density (njy) is given at 1 kev assuming a conversion of 1 Jy s Count 1, which is good to 10% for power law spectra with low column densities and spectral indices (α x)near 1.5 (10) Probability of having more counts than those observed in the specified region under the null hypothesis that the counts are background events. The jet is defined to be detected if P jet < 0.0025 (see Section 2.2) (11) X-ray jet detection (*) Cygnus A shows a correlation of X-ray and radio emission in the hotspot area which is not contained in the jet detection area and thus, is considered a detection for this sample.

24 2.2.1 0106+013(OC 12) Figure 2.1. Radio/X-ray overlay of 0106+013. The X-ray images were obtained from Chandra with VLA 1.4 GHz radio contours overlaid in black and white (see 3.1). The black VLA contours are set at 5 times the rms noise level for the lowest contour, with the exception of 0415+379 and 1849+670, which had their starting values set to 10 and 2.5 times the rms noise respectively, and multiples of 2 greater than that for each successive level. The white contours are offset from the black contours by 20%. The X-ray portion of each image has been energy filtered to a range of 0.5 to 7.0 kev in CIAO before being processed in DS9. The FWHM dimensions of the radio restoring beam are denoted by a cross in the bottom corner of each image and are also located in Table 2.4.

25 Table 2.4. VLA ARCHIVAL DATA Source Observation Date Project RMS Radio Noise B maj B min B maj PA (1) (2) (3) (4) (5) (6) (7) 0106+013 2004-09-19 AL634 1.4 10 01 1.64 1.49 100 0119+115 2004-09-19 AL634 9.8 10 02 1.53 1.44 55 0224+671 2004-09-19 AL634 1.5 10 01 1.42 1.13 78 0234+285 2004-09-19 AL634 2.4 10 01 1.20 1.11 52 0415+379 1982-06-14 LINF 1.9 10 01 1.60 1.47 168 0529+075 2004-09-19 AL634 4.6 10 02 1.70 1.35 63 0605+085 1993-01-17 AD298 3.0 10 01 1.84 1.11 166 1045 188 2007-06-30 AC874 3.4 10 01 1.00 1.00 90 1055+018 1992-11-18 AB631 1.8 10 01 2.00 1.33 52 1156+295 1984-12-24 AB310 1.2 10 01 1.39 1.32 43 1222+216 2004-11-20 AL634 1.7 10 01 1.23 1.11 54 1226+023 1990-03-23 AM297 2.0 10 0 1.00 1.00 90 1253+055 2001-01-11 W088D5 8.5 10 01 1.50 1.50 90 1334 127 1986-03-18 AD176 7.1 10 02 1.73 1.22 89 1510 089 1986-05-27 AB379 3.9 10 01 1.95 1.29 106 1641 399 1990-05-18 AS396 1.4 10 0 2.26 1.14 176 1655+077 1984-12-23 AB310 1.5 10 01 1.47 1.40 25 1800+440 1990-05-18 AS396 1.8 10 01 2.54 1.02 7 1828+487 1984-12-23 AB310 6.0 10 01 1.36 1.21 6 1849+670 2004-11-09 AL634 2.0 10 01 2.77 1.06 146 1928+738 1996-11-23 AS596 2.4 10 01 1.50 1.50 90 1957+405 1987-08-18 AC166 8.2 10 0 1.19 1.09 179 2155 152 2004-11-21 AL634 2.0 10 01 1.90 1.26 106 2201+315 2004-11-21 AL634 1.1 10 01 1.57 1.43 164 2216 038 2004-11-21 AL634 1.8 10 01 1.58 1.34 113 2251+158 1985-01-31 AC120 8.7 10 01 1.00 1.00 90 2345 167 2004-11-09 AL634 1.6 10 01 1.88 1.22 102 Note. Columns are as follows: (1) IAU name (B1950.0); (2) Date observed; (3) Project code; (4) Rms noise level of radio image in mjy beam 1 ; (5) Major axis for the radio beam in ( ); (6) Minor axis for the radio beam in ( ); (7) Position angle of the radio beam major axis in ( )

26 0106+013 is a very powerful blazar, which is located farther away from us than any other source in the MCS (z = 2.099) and has the second largest apparent speed (β app = 26.5). Physically this blazar shows a prominent radio jet which protrudes to the south terminating at 5 from the center of the nucleus [Hogan et al., 2011]. The X-ray jet correlates well with the radio jet and shows emission until it reaches the location associated with the terminal point of the radio jet. There is a small amount of extraneous radio emission to the northeast with an approximate angle of 45 which is not correlated with the X-ray emission (Figure 2.1). This jet does not show a lobe or hotspot in the counter-jet direction in either the X-ray or the radio band (Appendix B & C). Kharb et al. [2011] shows that at higher radio resolution (5 GHz) along with Chandra rebinning the jet shows a gentile S shape as it progresses from the core toward the terminal point. The jet also shows X-ray brightening and dimming as it moves toward and away from these subtle bends respectively [Kharb et al., 2011]. A Spectral Energy Distribution (SED) is shown in Kharb et al. [2011] along with the deeper Chandra images. 2.2.2 0119+115 0119+115 (Figure 2.2) is an example of a FSRQ which shows no X-ray emission above the level of the background emission besides the core structure. The radio jet shows emission to the north, east, and west of the core, with the most prominent radio feature located 6 to the north at an angle of 35. The diffuse radio emission to the east and west is located within a radius of 12 [Hogan et al., 2011].

Figure 2.2. Radio/X-ray overlay of 0119+115. 27

28 2.2.3 0224+671 (4C 67.05) Figure 2.3. Radio/X-ray overlay of 0224+671. Another example of a FSRQ with no correlation between the radio and X-ray bands, except for the core structure, is 0224+671 (Figure 2.3). The radio image for this AGN shows a long well collimated jet to the north 11 and a radio lobe in the counter jet direction at a distance of 7 from the core [Hogan et al., 2011]. 2.2.4 0234+285 (CTD 20) 0234+285 is a FSRQ which shows jet emission in both the radio and X-ray bands. A radio jet is seen to the north of the core structure and shows a sharp bend before it terminates abruptly. The X-ray emission correlates well with the radio emission up until the bend. At and beyond the bend there is no X-ray emission above the background level. There is an abundance of X-ray emission to the west of the core which is most likely a readout streak. Readout streaks are produced when there is a

29 Figure 2.4. Radio/X-ray overlay of 0234+285. pileup associated with the ACIS and is an artifact of the detector. The VLBI scale emission also shows structured emission at a PA of 12, which is closely aligned with the PA kpc (see Table 2.3). 2.2.5 0415+379 (3C 111) The first published image of this powerful radio galaxy was created by Linfield & Perley [1984]. The radio/x-ray overlay shows 4 distinct radio knots in the radio band, with three of those showing and excess of X-ray emission. The (1.4 GHz) radio knots are only seen in the primary jet within 100 of the core at a position angle of 63. Deeper VLBA images of this AGN show a pc scale jet with approximately the same direction as the kpc scale jet. The lack of a counter-jet is most likely due to Doppler boosting. Hotspots are seen in both the jet and counter jet directions even though there is not appreciable emission between the core and the hotspot associated with the counter jet. Both hotspots show large radio lobes which are approximately

30 Figure 2.5. Radio/X-ray overlay of 0415+379. one quarter of the total jet length. The hotspot associated with the primary jet does show X-ray emission above the background level, indicating an excellent correlation between the X-ray and radio emission in this AGN [Hogan et al., 2011]. This radio galaxy is very close to us (z = 0.0491) and thus probably does not have as small of an angle to the line of sight as most of the more heavily beamed sources in the MCS. The jet shows an angle of 7.9 according the IC/CMB calculations. Jorstad et al. [2005] provided a value for the angle to the line of sight for 3C111 of 18.1 ± 5.0, which is somewhat larger than the IC/CMB value that was tabulated. This jet has a superluminal speed of 5.9c and thus, has a maximum value of the angle to the line of sight of 19 for the pc scale jet [Lister et al., 2009a]. This source also shows an extended radio flux (S ext = 2.70) which is significantly larger than the majority of the rest of the MCS. It is interesting to note that all 5 of the sources which have a radio S ext > 1 show X-ray jet emission associated with the jet. This is one of only two radio galaxies in the MCS, the other being Cygnus A, and both show appreciable X-ray emission [Wilson et al., 2001].

31 2.2.6 0529+075 (OG 050) Figure 2.6. Radio/X-ray overlay of 0529+075. This blazar shows radio emission in both the southeast and southwest directions. The initial radio jet protrudes at a position angle of 145 and terminates at a distances of 6 from the nucleus. There is also a radio feature which located to the southeast of the core which could be an extension of the primary jet or emission from the counter jet. The counter jet hypothesis is further supported by the lack of X-ray emission in that area, because Doppler boosting of a jet away from the observer would not produce much emission if the emission mechanism is IC/CMB. The main jet emission in the X-ray band seems to coincide well with the primary radio jet, and there is a lack of coincidence with the emission to the southeast. The pc scale jet lies at a position angle of 45, which leads to the conclusion that there is bending between the pc and kpc scale jets in this instance [Hogan et al., 2011].

32 2.2.7 0605-085 Figure 2.7. Radio/X-ray overlay of 0605-085. 0605-085 shows an AGN with a radio jet structure which extends to the east at an angle of 90. When the X-ray image is superimposed there is a direct correlation between the radio and X-ray jets within 4 of the core, with the exception of the extra X-ray emission to the southwest. This extra emission is actually attributed to a foreground star that in the Chandra field of view [Sambruna et al., 2004], which also has a bright optical component. Sambruna et al. [2004] shows 5GHz radio data which reveals two knots between the core and the termination point, which are unresolved in the 1.4GHz data. There is HST data for this jet but it shows no emission above the respective background level in the optical regime [Sambruna et al., 2004]. 2.2.8 1045-188 This AGN shows strong, well collimated jet emission at a position angle of 125 and a large radio lobe directly opposed to it. This jet makes a 90 bend at a

33 Figure 2.8. Radio/X-ray overlay of 1045-188. distance of 8 from the core. The X-ray jet emission follows the radio jet until this bend and then shows a sharp decrease in emission. The counter jet lobe shows no appreciable X-ray emission above the background. If the IC/CMB model is dependent on the X-ray emission being beamed toward the observer then the jet bending in the plane of the sky could cause unseen jet bending toward or away from the observer, thus changing the amount of beamed emission that the observer would see [Hogan et al., 2011]. 2.2.9 1055+018 (4C 01.28) 1055+018 is an AGN that shows a radio jet extended toward the south which ends in a hotspot with associated lobe. There is a radio lobe located directly to the north presumably associated with the counter jet. There is no X-ray emission associated with this source except for the core emission and pile up associated with

34 Figure 2.9. Radio/X-ray overlay of 1055+018. the readout streak. There is no optical jet emission associated with the jet [Sambruna et al., 2004]. It is noteworthy that Sambruna et al. [2004] classifies this sources as a FSRQ/BL object. 2.2.10 1156+295 (4C 29.45) 1156+295 was one of the first objects to be imaged by Chandra (see Table 2.1 Column 7). There is no X-ray emission above the background level except for that which is associated with the core. 2.2.11 1222+216 (4C 21.35) This source was extensively studied by Jorstad & Marscher [2006]. The X- ray/radio overlays show an elongated radio jet with significant bending which ends in a terminal hotspot at a distance of 10 from the core. The radio jet hotspot has an angle of 90 and there is lobe-like radio emission to the south of the core. The X-ray

35 Figure 2.10. Radio/X-ray overlay of 1156+295. jet follows the initial nozzle of the jet (position angle 45 ) for a distances of 2 and decreases abruptly at the first bend. Again, this decrease could be attributed to the bend in the jet, which might change the line of sight of the jet. Jorstad & Marscher [2006] provide a 5 GHz radio image 1222+216 and cite 2 knots in the jet. The knots are unresolved at 1.4 GHz. They also produce non-unique SEDs for one of the knots (Jorstad & Marscher 2006 & Appendix D). 2.2.12 1226+023 (3C 273) The jet associated with 1226+023 has been extensively studied by quite a few groups [Marshall et al., 2001, Sambruna et al., 2001, Jester et al., 2006] due to its unique X-ray, optical, and radio correlations. The radio/x-ray overlay that I produced shows a radio jet extending to the southwest with an angle of 135 and terminating at 22 from the core. The X-ray emission correlates spatially with the radio jet rather well but peaks earlier than the radio jet does. The radio jet increases its

36 Figure 2.11. Radio/X-ray overlay of 1222+216. emission along the jet until the hotspot where it reaches its maximum flux. The X- ray emission, on the other hand, starts peaked at 14 from the core and decreases its emission until it reaches the terminal region of the radio jet. This is also one of the few jets where the pc and kpc scale position angles are very close to alignment (Table 2.3) This source also has one of the lowest redshifts (z = 0.158) and largest S ext (17.67) values in the MCS. Marshall et al. [2001] showed that the optical (HST) and X-ray jets of 3C 273 are very similar in length and width, with the exception of the emission levels of the first knot. A second major difference is that the X-ray emission fades as the distances from the core increases, while the optical jet emission does not [Marshall et al., 2001]. Sambruna et al. [2001] supports our initial visual inspection of this source, and provides 2D surface brightness profiles which show how the emission relates between the optical, X-ray, and radio wavelengths (see Fig. 2, Sambruna et

37 Figure 2.12. Radio/X-ray overlay of 1226+023. al. 2001). Marshall et al. [2001] & Sambruna et al. [2001] differ on their choice of emission mechanism. Marshall et al. [2001] provides an SED which shows that a single population of electrons using a pure synchrotron emission mechanism can model the first two knot fluxes, while Sambruna et al. [2001] shows that the external Compton ( IC/CMB) fits the data points used in the SED. Jester et al. [2006] believes that the single zone IC/CMB model can only be used reliably in the first fifth of the jet. They invoke either a two-zone IC/CMB model but warn that extreme bulk Lorentz factors may be needed further down the jet, and or a two-zone synchrotron model, where particle acceleration is related to a velocity shear which could produce the X-ray emitters, to describe the jet features [Jester et al., 2006]. They also provide SEDs to further reinforce the model choices (Figure 2, Jester et al. 2006).

38 Figure 2.13. Radio/X-ray overlay of 1253-055. 2.2.13 1253-055 (3C 279) The image of 3C 279 shows a one sided radio jet that stretches toward the southwest until it bends very sharply toward the east at 4.5 from the core (initial position angle 150 ). There is lobe emission to the north west at 12 from the core with a position angle of 37. The X-ray emission correlates well with the primary radio jet until the bend, where there is a sharp decrease in X-ray flux. Collmar et al. [2010] shows multiwavelength analysis of the spectrum and morphology of 3C 279, by using Chandra, SWIFT, INTEGRAL, RXTE, and other telescopes associated with WEBT (Whole Earth Blazar Telescope). They produce SEDs which are modeled using a leptonic one-zone SSC + EC model, and that the X-ray spectrum is entirely produced by SSC emission ( Collmar et al. 2010 & Appendix D).

39 Figure 2.14. Radio/X-ray overlay of 1334-127. 2.2.14 1334-127 This blazar has an X-ray jet with a length of 6 that follows the radio jet emission out to a 60 bend of the radio jet, then undergoes a drop in emission, but still terminates at the same point as the radio jet. Both jets initially follow a position angle of 135. The emission characteristics in the bend region are significantly different than the jet of 1045 188, which undergoes a sudden drop in X-ray emission after the bend [Hogan et al., 2011]. 2.2.15 1510-089 The (1.4 GHz) radio and X-ray jets extend to the southeast, bending slightly toward the south. The X-ray jets terminates well before the radio jet, at a distance of 5, which happens to coincide with the bend of the radio jet. The radio jet

40 Figure 2.15. Radio/X-ray overlay of 1510-089. continues past the bend for another 5 until it ends. Sambruna et al. [2004] also presents the HST optical data for this source but there are no optical counterparts for the X-ray or radio knots. The higher resolution 5GHz radio data shows three possible X-ray knot detections [Sambruna et al., 2004]. 2.2.16 1641+399 (3C 345) 3C 345 shows a one sided radio jet which protrudes to the northwest of the core. The X-ray emission from this object follows the radio emission closely and terminates at roughly the same spot, 3 from the core. The 5GHz radio data shows a more defined knot structure in the radio jet, which is also closely aligned with the X-ray knot [Sambruna et al., 2004]. Sambruna et al. [2004] also presents HST optical data which corresponds to the knot structure. The optical data point helps constrain the SED which helps determine if the IC/CMB model does indeed describe the primary emission mechanism of X-rays in blazar FR II jets ( 4.3 & Appendix D). Deeper

41 Figure 2.16. Radio/X-ray overlay of 1641+399. Chandra observations of this source were obtained by Kharb et al. [2011] along with new optical data from HST, which was used to construct the SED. The new Chandra images indicate that the X-ray hotspot is actually located closer to the core than the radio hotspot [Kharb et al., 2011]. 2.2.17 1655+077 This AGN shows a short jet extending to the southeast and a longer radio jet which extends to the northwest and then makes a sharp ( 90 ) bend toward the southwest [Marshall et al., 2005]. There is no X-ray emission above the background level other than the core region. VLBA data shows pc scale knots along the position angle of 38 (Marshall et al. [2005] and reference within) which is closely aligned with the kpc scale position angle ( 50 ).

42 Figure 2.17. Radio/X-ray overlay of 1655+077. 2.2.18 1800+440 (S4 1800-44) This AGN shows two sided emission. The primary radio jet is located southwest of the core and the counter jet lobe emission is located to the northeast. The X-ray jet follows the radio jet at a position angle of 130 until the first apparent bend and then abruptly terminates ( 3 ). The radio jet continues for another 3 past the bend. This is another example where the X-ray jet flux decreases beyond a radio knot located at a bend in the jet [Hogan et al., 2011].

Figure 2.18. Radio/X-ray overlay of 1800+440. 43

44 2.2.19 1828+487 (3C 380) Figure 2.19. Radio/X-ray overlay of 1828+487. 3C 380 is an AGN that harbors radio emission on both sides of its core with the more prominent jet protruding toward the northwest. The position angles of the pc and kpc scale jets are within 1 of alignment (see Table 2.3 columns 2 & 3). The X-ray jet profile shows a significant difference between the primary jet and counter jet emission (Appendix B). Marshall et al. [2005] places the X-ray emission knot at a distance of 1.8 from the core. This source has the third largest S ext (5.43) in the MCS. 2.2.20 1849+670 (S4 1849-67) 1849+670 is a FSRQ that shows a radio jet to the north and a radio lobe structure to the south. The radio emission to the north shows emission until 15 from the core. The X-ray jet follows the radio jet closely for 9 and then abruptly stops. The overlays show no correlation between the radio lobe and the X-ray band when

45 Figure 2.20. Radio/X-ray overlay of 1849+670. inspected visually. Again, this supports the IC/CMB emission mechanism as it is assumed that the primary jet is boosted toward the observer and the counter jet is boosted away. 2.2.21 1928+738 (4C 73.18) 4C 73.18 is an AGN which shows an elongated radio jet which extends to the south with a slight bend toward the east. It also shows radio emission to the north centered around 6 from the core. The X-ray emission briefly follows the radio jet and stops abruptly at a distance of 7 from the core. The radio jet bends and continues to a distance of 17 from the core until it terminates. Sambruna et al. [2004] uses 5 GHz radio data along with Chandra images and finds only one X-ray knot, which has an optical (HST) counterpart. The SED composed by Sambruna et al. [2004] indicates that this sources is different than most from our sample by implying that the emission

46 Figure 2.21. Radio/X-ray overlay of 1928+738. mechanism for the X-ray component is pure synchrotron emission instead of IC/CMB emission. The radio jet of the source also differs morphologically from most of the other FR II sources in the sample, as it is not well collimated and terminates in a poorly confined lobe, which reinforces the different emission mechanism proposed by Sambruna et al. [2004] for this source. 2.2.22 1957+405 (Cygnus A) Cygnus A is the second closest (z = 0.0561) and has the largest extended flux (S ext = 414.18) when compared to the rest of the MCS. This source is the second of two radio galaxies in the sample (the other being 3C 111) meaning that its jets are traditionally viewed at a larger angle to the line of sight than blazars. The radio image shows a well collimated jet to the northwest and extremely large lobes that surround the hotspots from both the jet and counter jet. The X-ray image shows a correlation at the terminal hotspot locations, which implies that there is a jet, but

47 Figure 2.22. Radio/X-ray overlay of 1957+405. an excess of X-ray emission near and around the core prevents any jet detections. Visually the only cores and hotspots of both the radio and X-ray images are aligned. This excess of X-ray emission is attributed to emission from the cavity of Cygnus A [Wilson et al., 2006]. This cavity is interpreted as energy which the jet can not move efficiently to the lobe, and is often referred to as a cocoon (Wilson et al. [2006] and references within). Wilson et al. [2001] proposed that the emission model for Cygnus A is Synchrotron Self-Compton (SSC) radiation, as this is the mechanism of radiation for most FR II type radio galaxies. This could imply that there is a difference between the radiative processes of blazars and radio galaxies, even though they are both members of the FR II class. 2.2.23 2155-152 2215-152 is an FSRQ which shows radio emission to the north and south of the core. The pc scale jet is oriented to the south so I assumed that the primary jet on the kpc scale was also oriented to the south. The southern X-ray jet is observed until a distance of 4 where it abruptly stops before the radio jet which shows a sharp

48 Figure 2.23. Radio/X-ray overlay of 2155-152. decrease in emission between 8 [Hogan et al., 2011]. The northern lobe is centered around a point which is located at 5 from the core. 2.2.24 2201+315 (4C 31.63) 4C 31.63 shows an elongated radio jet which loses collimation before it terminates in a hotspot at a distances of 37 with a position angle of 100. This AGN also has a radio hotspot associated with the counter jet which is located at 45 (PA kpc = 58 ) from the core. The X-ray emission correlates with the radio jet up until a distance of 5 from the core, where there is also a decrease in radio jet emission and loss of collimation (see appendices A & B). This X-ray jet detection is considered marginal due to its lack of visual correlation with the radio jet and probability of jet detection value (P jet, Table 2.3) which is very close to the threshold value. A longer exposure time is needed to produce a clear visual correlation.

49 Figure 2.24. Radio/X-ray overlay of 2201+315. 2.2.25 2216-038 This blazar shows two radio knot structures and a hotspot on the kpc scale oriented toward the southeast and terminates at 15 from the core. There is also a radio lobe structure located to the north at a distance of 19. The closest radio knot is located at a distance of 5 from the core, while the more prominent knot is centered around 9 from the core at the same PA kpc. The X-ray emission is visually correlated well with the second knot and the hotspot but not the first knot. There is no X-ray emission associated with the radio lobe above the level of the background. This jet also bends significantly from the pc to the kpc scale showing a position angle misalignment of roughly 60%.

50 Figure 2.25. Radio/X-ray overlay of 2216-038. Figure 2.26. Radio/X-ray overlay of 2251+158. 2.2.26 2251+158 (3C 454.3) The radio image of this FSRQ shows a well collimated jet which ends in a knot or hotspot at a distance of 5 from the core. The X-ray emission mirrors the radio

51 emission and terminates at roughly the same point [Marshall et al., 2005]. Marshall et al. [2005] presents VLBA data that shows the pc scale jet curving to align with the position angle of the kpc X-ray jet ( 50 ). 2.2.27 2345-167 Figure 2.27. Radio/X-ray overlay of 2345-167. 2345-167 is a small one sided radio blazar which shows radio jet emission in a southwest direction until a distance of 5 from the core where there is a sharp decrease in flux. Visually there seems to be a marginal correlation and the X-ray profile confirms this. This sources does have a P jet value which is close to the threshold, but does not surpass it and thus is not considered a detection. A longer exposure time on this source could allow for an X-ray detection.

52 3. DATA REDUCTION AND ANALYSIS 3.1 X-ray Radio Overlays Figures 2.1 through 2.27 present the X-ray-radio overlays for the sources in the MCS. The following procedure was used to produce these images. 1.4 GHz data were first obtained using the VLA A-array data from the NRAO 1 data archive and our own observations [Cooper et al., 2009, Kharb et al., 2010]. These data were reduced following the standard procedures in the Astronomical Images Processing System (AIPS). Standard calibrators were used to initially calibrate the amplitude and phase of the sources. Then, the tasks IMAGR and CALIB were used iteratively to selfcalibrate the image and the sources. This self calibration was performed on both the amplitudes (with successively decreasing solution intervals) and phases of the visibilities (with solution interval times typically set to less than 0.5 mins in CALIB) until convergence in image flux and structure was achieved. This procedure led to the production of radio maps which had a typical rms noise of 0.2 mjy beam 1. The specific radio noise levels for each source map are located in Table 2.4. The average FWHM restoring beam of the radio images was assumed to be 1.4, which is slightly worse than the FWHM of Chandra, which was estimated to be 0.75 [Hogan et al., 2011, Marshall et al., 2005]. The X-ray portions of the overlays were created by using both archival data and proprietary data from Chandra [Hogan et al., 2011]. The observation dates and exposure times for all of the MCS targets are listed in Table 2.2. Chandra maps for the overlays were created using the DS9 imaging tool, by using the level 2 event files. First, the event files were loaded into CIAO (Chandra Interactive Analysis of 1 The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.

53 Observations), and were filtered to an energy range of 0.5 to 7 kev. Once filtered to the appropriate energy range, the files were loaded into DS9 for imaging. The files were then smoothed using the smoothing tool, which was set to the Gaussian setting with a kernel radius of 3 pixels, where a pixel is equivalent to 2. The pseudo-color scale was then adjusted so that the core structures were oversaturated (black), so that the X-ray emission from the jets above the noise level could be easily identified through the use of visual inspection. At this point both the radio and X-ray images were ready to be merged. I started with the radio image and then aligned the X-ray image to it using the WCS (World Coordinate System) frame matching setting in DS9. There were a few sources which had core positions that were slightly misaligned when the overlays were inspected. To compensate for the misalignment of the cores, the radio and X-ray images were registered by using the Fv program in the Ftools package provided by NASA 2 [Blackburn, 1995]. The images were manually shifted by a small amount (generally of order of less than 2 pixels) so that the nucleus of each AGN was aligned between the X-ray and radio bands [Hogan et al., 2011]. The radio beam size is represented by a cross in the lower right hand corner of the radio/x-ray overlays, and has specific properties listed in Table 2.4. The beam data were extracted from AIPS by using the IMSTAT task, which provided the length of the major and minor axes of the beam as well as the major axis angle. The rms radio noise was then obtained from IMSTAT (Table 2.4). The lowest radio contour level was set to 5 times the noise level, except for 0415+379 and 1849+670, which had their lowest level contours modified to better represent the radio jet emission. 0415+379 and 1849+670 had their lowest contour set to 10 and 2.5 times the noise level, respectively. Each successive contour past the first was set to twice the level of the previous contour. This process was repeated for a total of 10 iterations. The black contour curves represent these contours and the white contour curves are set 20% higher than the black ones. Once the contour level values were known, they were drawn in DS9 by using the contour parameter tool. The levels were manually entered into the tool and 2 Information about Ftools can be found at http://heasarc.gsfc.nasa.gov/ftools/

54 placed on the images by using a contour smoothness setting of 1. Once placed on the radio image, the curves were then moved to the X-ray image, which was located on a seperate frame in DS9. The images were then saved in a tiff format, as DS9 does not have the ability to save files in a postscript format. The files were then converted to postscript files by using GIMP2 (GNU Image Manipulation Program). 3.2 X-ray and Radio Jet Analysis The X-ray jet detection portion of the analysis was completed using the IDL (Interactive Data Language) programming language and scripts which were developed by Herman Marshall and modified by Brandon Hogan [Marshall et al., 2005, Hogan et al., 2011]. In essence, the scripts take the radio and X-ray FITS files and compare the regions of the radio and X-ray jets which are specifically chosen based on radio jet criteria (i.e., radio jet length and width). The script starts by identifying the core position of both the radio and X-ray FITS files which were calculated via the method described by Marshall et al. [2005]. The X-ray nucleus location was first deduced by fitting Gaussians to the one dimensional histograms obtained from events which were detected within a distance of 30 of the hypothesized core region, which was chosen by visual inspection of the FITS files. This estimated centroid position was then used to acquire a more accurate central reference point for the core by repeating the previous process using a region which was defined by a radius of 3 from the previously calculated rough centroid position. This produces a more accurate centroid position and reduces the contamination effect of extended jet emission, which can bias the core positions. Poisson statistics were then used to test for the existence of an X-ray jet. The radio images were used to create a rectangular region on the X-ray map which defines the original shape and size of the radio jet. The rectangular region is defined by the parameters PA kpc (kiloparsec position angle), R i (jet inner radius), and R o (jet outer radius, see Table 2.3). The length of each region was allowed to vary, but the

55 width of each was fixed at 3. The rectangular region was also offset from the center of the core by 1.5 to eliminate the emission from the nucleus, with the exception of 1849+670 and 2345-167, which were offset by 5 and 2 respectively. Shortening these rectangular regions can eliminate jet contamination from the core emission region which is often associated with the elongated restoring beams found in 1.4 GHz radio data [Hogan et al., 2011]. The algorithm assumes that the radio jets show no bending on the kpc scale and that the area 90 to the primary jet axis is free of real emission from the jet. This assumption is valid for all of the jets in the MCS except for 0119+115, which shows radio jet emission on both sides of the core perpendicular to the direction of the primary jet axis, but does not show any X-ray emission above the level of the background X-ray emission. Hence, the perpendicular jet emission is inconsequential in this case. Then, the radio emission was plotted in two dimensions to show a profile of the jet emission along the axis of the jet and in the region perpendicular (90 ) to the jet (Appendix A). The jet emission and the background emission (emission located in the region perpendicular to the jet axis) were then subtracted to eliminate the core emission (bold curve). The X-ray profiles were created by using the radio jet emission region and the region located at the opposite side (180 ) of the core (Appendix B). This background region was chosen because counter-jets associated with blazars are not seen very often in X-rays, probably due to the effects of Doppler boosting [Worrall, 2009]. Visual inspection of the MCS also reinforces this concept, as very few X-ray counter jets are seen. The CCD readout streaks associated with Chandra also cannot contaminate the counter jet area as they are offset from the main jet axis, and hence also the counter jet axis. The X-ray counts in these regions were compared using Poisson statistics, with a Poisson probability threshold for jet detection set to 0.0025 [Marshall et al., 2005, Hogan et al., 2011]. The chosen probability threshold produces a 5% chance that a false detection will occur in 1 out of every 20 sources. The X-ray fluxes were computed from count rates using a conversion factor of 1 µjy per count s 1, which is accurate to about 10% for typical jet power law spectra [Marshall et al.,

56 2005]. The sources in Table 2.3 with negative values in the count rate column have less X-ray emission in the jet region than in the corrosponding background region. Since visually there is often not X-ray emission in the counter jet region, these negative count-rates are often small and are associated with a null detection in the X-ray jet column. The analysis method indicated that there are X-ray detections in all of the sources except for 0119+115, 0224+671, 1055+018, 1156+295, 1655+077, 1957+405 (Cygnus A), and 2345 167, although 1957+405 was treated as an X-ray detection. Visual inspection of the latter source shows coincident hotspot emission in both the X-ray and radio bands (see 2.2.22 for further discussion), as well as an excess of X- ray emission around the core which contaminates the regions that have been chosen for jet and counter jet emission. The sources without X-ray detections show little to no X-ray emission above the background level except for their respective nuclei (Appendix A, Table 2.3), despite the fact that their radio structure, redshift values, β app values, and S ext values are comparable to other sources in the MCS. Tables 3.1 and 2.3 list all of the relevant X-ray emission data for the sample. 3.3 The Single Zone IC/CMB Model The analysis for the X-ray portion of the MCS was done using a method similar to the method used by Marshall et al. [2005]. There are a few major assumptions that are associated with the model and are as follows. The magnetic fields are in equipartition with the particle energies for the kiloparsec scale jets. This assumption allows for the minimum energy magnetic field strength of the jets to be calculated. If the magnetic field strength and CMB photon density are known, the photon energy density can be compared to the magnetic field energy density. This leads to the use of a parameter K, which is calculated by combining the magnetic field strength with the ratio of X-ray (inverse-compton) and radio (synchrotron) luminosities. The K parameter is a function of β, (Equation 1.1) and θ, the angle to the line of sight. Another important parameter is the apparent speed (β app ), which is a parameter

57 Table 3.1. MOJAVE CHANDRA SAMPLE BEAMING MODEL PARAMETERS Source α rx R V B 1 K δ θ Γ pc=kpc Γ kpc,min Γ kpc,decel (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (10) 0106+013 0.94 ± 0.01 0.0726 1.6 10 3 148. 13 ± 2 3.6 4.2 99 +33 29 1.9 1.9 +0.1 0.1 0119+115 > 0.88 < 0.1996 7.7 10 2 29. < 19 4.3 > 6.2 36...... 0224+671 > 0.91 < 0.1356 9.9 10 2 26. < 15 3.9 > 8.9 20...... 0234+285 0.84 ± 0.01 0.4362 1.1 10 3 58. 29 ± 7 5.4 7.8 17 +2 2 2.7 3.2 +0.6 0.5 0415+379 0.83 ± 0.02 0.5590 3.7 10 1 19. 51 ± 12 7.1 7.9 6 +1 1 3.5 6.0 +1.9 1.9 0529+075 0.93 ± 0.02 0.0846 1.6 10 3 57. 11 ± 2 3.3 8.4 26 +7 6 1.7 1.8 +0.2 0.1 0605 085 0.84 ± 0.01 0.4416 7.2 10 2 61. 43 ± 10 6.6 5.2 33 +4 3 3.3 3.7 +0.5 0.6 1045 188 0.95 ± 0.02 0.0639 1.1 10 3 49. 17 ± 3 4.1 10.8 11 +2 2 2.1 2.5 +0.4 0.3 1055+018 > 0.92 < 0.1053 2.3 10 3 42. < 13 3.6 > 9.3 19...... 1156+295 > 0.91 < 0.1308 3.3 10 2 65. < 27 5.2 > 4.4 62...... 1222+216 0.85 ± 0.01 0.3921 1.9 10 2 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 10 2 91. 75 ± 17 8.6 6.0 15 +1 1 4.3 6.1 +1.3 1.3 1253 055 1.02 ± 0.01 0.0166 4.3 10 2 96. 17 ± 3 4.1 5.3 54 +5 5 2.1 2.2 +0.2 0.2 1334 127 0.83 ± 0.01 0.6044 1.1 10 3 40. 50 ± 14 7.0 7.5 11 +2 2 3.6 5.1 +1.4 1.4 1510 089 0.81 ± 0.01 0.8796 2.4 10 2 41. 82 ± 21 9.1 4.7 27 +3 2 4.6 5.4 +1.1 1.0 1641+399 0.90 ± 0.02 0.1448 3.1 10 2 60. 33 ± 7 5.7 5.4 35 +5 3 2.9 3.2 +0.3 0.4 1655+077 > 0.92 < 0.1008 4.7 10 2 46. < 19 4.3 > 7.2 26...... 1800+440 0.89 ± 0.01 0.1709 9.4 10 2 52. 28 ± 6 5.3 6.6 25 +3 3 2.7 3.0 +0.3 0.4 1828+487 0.81 ± 0.01 0.8936 2.3 10 2 52. 67 ± 16 8.2 6.2 16 +2 1 4.1 5.6 +2.1 1.2 1849+670 0.84 ± 0.02 0.5029 2.2 10 3 18. 18 ± 4 4.2 3.7 114 +15 17 2.2 2.2 +0.3 0.2 1928+738 0.83 ± 0.01 0.6039 3.6 10 2 27. 48 ± 12 6.9 8.1 8 +1 1 3.5 5.7 +1.8 1.8 1957+405 > 1.32 < 0.0001 3.5 10 1 151. < 2 1.2 > 40.9 1...... 2155 152 0.99 ± 0.03 0.0253 1.6 10 3 53. 10 ± 2 3.2 6.1 55 +10 13 1.7 1.7 +0.2 0.2 2201+315 0.87 ± 0.03 0.2438 3.2 10 2 27. 29 ± 7 5.3 9.9 9 +1 2 2.7 3.9 +0.9 0.9 2216 038 0.97 ± 0.02 0.0408 2.8 10 3 49. 9 ± 1 3.0 15.8 7 +1 2 1.6 1.9 +0.2 0.1 2251+158 0.93 ± 0.01 0.0910 7.7 10 2 101. 30 ± 6 5.5 7.0 21 +3 2 2.8 3.2 +0.4 0.4 2345 167 > 0.91 < 0.1227 7.2 10 2 44. < 22 4.7 >7.6 22...... 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 kpc 3 (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 β app equations simultaneously, along with the assumption that the value of β app 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 [ ] 2/7 18.85 C12 (1 + k)l sync B 1 = (3.1) ΦV is defined first, where B 1 the spatially averaged, minimum energy magnetic field of the jet in Gauss, when there is no Doppler boosting (δ = 1). C 12 is a weak function of the low frequency spectral index of the synchrotron spectrum (α r, where S ν ν αr ), Φ is the filling factor, L sync 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 10 7, α r =0.8, and Φ=1. The emitting volume for the jet is calculated using the R i and R o 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 (B 1 ) 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 B 1 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. R = S x(ν/ν x ) αr S r (ν/ν r ) αr = S xνx αr S r νr αr = [ νx ν r ] αr α rx, (3.2) where ν r and ν x are the radio and X-ray frequencies at which the flux densities S r and S x are observed, respectively. The jet frame value for L sync 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 10 17 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 = B 1 (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 10 10 Gauss 2 and b=3.808 10 4 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 + β2 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 β app 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 = β app β app µ + 1 µ 2 (3.6) θ = arctan 2β app β 2 app + δ 2 1 (3.7) Γ = β2 app + δ 2 + 1 2δ (3.8) Specifically, β can be represented in terms of β app 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 β app and K. Marshall et al. [2005] showed that a change in B 1 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 (α r ), which had a value set to 0.8 for the IC/CMB calculations. Common observed values for α r at kpc scale distances are between 0.7 and 0.9. Since, actual measurements α r do not exist for the MCS, Monte Carlo error analysis was carried out on the sample, by defining a Gaussian distribution of α r with α r = 0.8 and σ αr = 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 β app. The error values for these curves are produced by attributing the error from the β app 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 β app 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 β app ( 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 β app,max and parent population Γ max 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 β app 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 Γ min,decel values are calculated for the sources in the sample with X-ray jets. Values associated with Γ min,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 β app ). 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 Γ min, 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 Γ min values greater than 3.5 [Hogan et al., 2011]. Γ min = K 2 K 1 (3.9) Equation 3.4 also sets a limit for the value of θ kpc,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 θ kpc,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 β app and θ. This relationship, described by Equation 3.6, confines Γ β app, and θ 2 tan 1 (β 1 app). 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 S ext 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, S ext, 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 S ext > 300 mjy (Figure 3.1) and a significantly lower detection rate ( 57%) for sources with S ext 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 β app values with respect to the detection of sources, the β app values with respect the S ext 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 9 Quasar or Radio Galaxy w/ X ray jet Quasar w/out X ray jet 8 7 Number 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 S ext value. All sources with a S ext 300 mjy show a correlation between the X-ray and radio bands at some level.

68 associated with the S ext 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 β app 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 β app values for only 107 of the 135 sources could be calculated at this time. Thus, the selection criteria does not change the distribution of β app with respect to the MCS, even though it alters the distributions of redshift and S ext.

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 (θ n ), 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( PA) = sin ζ sin φ cosζ sin θ n + sinζ cos θ n cosφ (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( PA) sin(φ) ( ). (3.11) θ nζ + cos(φ) The small angle approximation for θ n 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 θ n 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 θ n and ζ also has to change for a random value of φ. In this case ζ approaches θ n. Large values of PA ( 90 ) would require that ζ approaches and surpasses the value of β 1 app. Equation 3.11 can still be satisfied because most sources in the MCS have large β app values. Moore et al. [1981] states that large values of PA can be obtained with small values of θ max, where θ max is the largest value of θ n which is likely to occur [Hogan et al., 2011]. A value of PA which approaches and exceeds 90 can only be obtained when θ n 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 θ n 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 ( PA kpc PA pc ) 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 θ n 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 D L δ radius B w psoll γ min γ max n (1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) 0415+379 3C 111 0.0491 2.11 10 8 2 1 10 22 1.3 10 5 3.1 10 10 15 3 10 5 2.7 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 B 1 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 (D L ), 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 (γ min and γ max ) 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 (w psoll ) was shifted to align the curves and the data points. 4.2.1 Obtaining the Radio Fluxes The radio fluxes were extracted for a region which mirrored the X-ray hotspot area ( 2 radius). The data were extracted with AIPS by using the task IMSTAT for the given region. This region was defined by using TVBOX to select the region on the tv window. The radio fluxes were already in Jy, so they were converted to erg cm 2 sec 1 for the given frequency that they were observed at (1.4GHz and 5GHz). The two points in the radio band constrain the power-law index for the electron energy distribution. 4.2.2 Obtaining the Optical Fluxes The optical data information was taken from the HST drizzle file by extracting the region from DS9. The ACS extracted regions provide the number of electrons/sec. The fits header has a PHOTFLAM keyword, which when multiplied by the previous quantity, produces a flux in terms of erg cm 2 s 1 Å 1. This information along with the observing wavelength allowed for the procurement of the SED optical point. The optical point constrains the well, under the assumption that it lies on the synchrotron curve.

77 4.2.3 Obtaining the X-ray Fluxes The X-ray information was taken from the number of counts in the selected circular region associated with the hotspot of 3C 111 ( 2 radius). After the region was selected I used the virtual observatory, which is located under the analysis tab to open the Chandra-Ed Archive Server 1. 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 PIMMS 2 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 E min to E max 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 chandra-ed.cfa.harvard.edu/archive.html 2 http://cxc.harvard.edu/toolkit/pimms.jsp

78 dn de = AEγ (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) Emax γ+1 E γ+1 min (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. νf(ν) = E 2+γ avg AC (4.3) and the frequency (ν) is defined as ν = E avg 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 10 9 2.5 10 14 VLA 5.0 10 9 2.9 10 14 HST 3.8 10 14 5.4 10 16 Chandra 2.4 10 17 2.7 10 15 Fermi 3.0 10 23 7.1 10 12 Fermi 9.9 10 23 5.1 10 12 Fermi 3.0 10 24 2.9 10 12 Fermi 9.9 10 24 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 10 15 Hz, and an X-ray curve which models emission from 10 16 Hz to 10 25 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