Multi-Wavelength Analysis of Active Galactic Nuclei

1. Multi-Wavelength Analysis of Active Galactic Nuclei A dissertation submitted as partial ful

2. lment of the 100-hour certi

3. cate course in Astronomy & Astrophysics by Sameer Patel M.P. Birla Institute of Fundamental Research Bangalore, India December 2014

4. Declaration I, Sameer Patel, student of M.P. Birla Institute of Fundamental Research, Bangalore, hereby declare that the matter embodied in this dissertation has been compiled and prepared by me on the basis of available literature on the topic titled, Multi-Wavelength Analysis of Active Galactic Nuclei as a partial ful

5. llment of the 100 Hour Certi

6. cate Course in Astronomy and Astro-physics, 2014. This dissertation has not been submitted either partially or fully to any university or institute for the award of any degree, diploma or fellowship. Date: Place: Signature Director, M.P. Birla Institute of Fundamental Research, Bangalore i

7. M.P. Birla Institute of Fundamental Research Bangalore, India Abstract Multi-Wavelength Analysis of Active Galactic Nuclei by Sameer Patel This dissertation explores the current research methods and analysis adopted for the study of Active Galactic Nuclei in all wavelengths of the electromagnetic radiation. Being the most violent objects that one can see in the present Universe, AGNs have been attributed to emitting radiation in all wavelengths and still exhibit various unexplained phenomena, alongside with being the probes to the very early Universe. The uni

8. cation of the AGN model is also included for completeness, albeit not con

9. rmed in its entirety.

10. Acknowledgements I would never have been able to

11. nish my dissertation without help from friends, and support from the team at MPBIFR, Bangalore. I would also like to thank Dr. Babu for constantly reminding us to complete the dis-sertation timely, and Ms. Komala for guiding me to coast through countless papers online for reference. I would like to thank Rishi Dua, who as a good friend, was always willing to help me and give his best suggestions, and Aakash Masand, who helped me correct typographical errors and grammatical mistakes after painfully proofreading the

12. nal draft. I would also like to thank my parents. They were always supporting me and encouraging me with their best wishes. iii

13. Contents Declaration i Abstract ii Acknowledgements iii List of Figures vii List of Tables ix Abbreviations x 1 Introduction 1 1.1 The History of AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1 1.2 Active Galactic Nuclei . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1 1.3 The Taxonomy of AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2 1.3.1 Seyferts . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 3 1.3.2 Quasars and QSOs . . . . . . . . . . . . . . . . . . . . . . . . . . . 3 1.3.3 Radio Galaxies . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 5 1.3.3.1 Radio Quiet . . . . . . . . . . . . . . . . . . . . . . . . . 6 1.3.3.2 Radio Loud . . . . . . . . . . . . . . . . . . . . . . . . . . 6 1.3.4 Blazars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 8 1.3.4.1 BL Lacerate Objects . . . . . . . . . . . . . . . . . . . . . 8 1.3.4.2 Optically Violent Variable Quasars . . . . . . . . . . . . . 9 1.3.5 LINERs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 9 2 Non-Thermal Processes 12 2.1 Basic Radiative Transfer . . . . . . . . . . . . . . . . . . . . . . . . . . . . 12 2.2 Synchrotron Radiation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13 2.2.1 Emission by a Single Electron in a Magnetic Field . . . . . . . . . 13 2.2.2 Synchrotron Emission by a Power Law Distribution of Electrons . 14 2.2.3 Synchrotron Self-Absorption . . . . . . . . . . . . . . . . . . . . . 15 2.2.4 Polarization . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15 2.2.5 Synchrotron Sources in AGNs . . . . . . . . . . . . . . . . . . . . . 16 2.2.6 Faraday Rotation . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17 2.3 Thomson Scattering . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 18 iv

14. Contents 2.4 Compton Scattering . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19 2.4.1 Comptonization . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20 2.4.2 The Compton Parameter . . . . . . . . . . . . . . . . . . . . . . . 21 2.4.3 Inverse Compton Emission . . . . . . . . . . . . . . . . . . . . . . 22 2.4.4 Synchrotron Self-Compton . . . . . . . . . . . . . . . . . . . . . . . 23 2.5 Annihilation and Pair-Production . . . . . . . . . . . . . . . . . . . . . . . 24 2.6 Bremsstrahlung (Free-Free) Radiation . . . . . . . . . . . . . . . . . . . . 26 3 The IR and Sub-mm Regime 27 3.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27 3.2 Observations and Detection . . . . . . . . . . . . . . . . . . . . . . . . . . 28 3.3 The Dusty Torus . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 29 3.4 IR Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 31 3.4.1 The 1 m Minimum . . . . . . . . . . . . . . . . . . . . . . . . . . 32 3.4.2 IR Continuum Variability . . . . . . . . . . . . . . . . . . . . . . . 33 3.4.3 The Submillimeter Break . . . . . . . . . . . . . . . . . . . . . . . 33 4 The Radio Regime 34 4.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 34 4.2 The Loudness of AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . 35 4.3 The Fanaro-Riley Classi

15. cation . . . . . . . . . . . . . . . . . . . . . . . 36 4.3.1 Fanaro-Riley Class I (FR-I) . . . . . . . . . . . . . . . . . . . . . 36 4.3.2 Fanaro-Riley Class II (FR-II) . . . . . . . . . . . . . . . . . . . . 37 4.4 Radio Lobes and Jets . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 39 4.4.1 The Generation of Jets . . . . . . . . . . . . . . . . . . . . . . . . 40 4.4.2 The Formation of Radio Lobes . . . . . . . . . . . . . . . . . . . . 40 4.4.3 Accelerating the Charged Particles in the Jets . . . . . . . . . . . . 42 4.4.4 Superluminal Velocities . . . . . . . . . . . . . . . . . . . . . . . . 43 5 The Optical-UV Regime 44 5.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44 5.2 Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44 5.2.1 The Optical-UV Continuum and the Accretion Disk . . . . . . . . 45 5.3 Observations in the Optical-UV Region . . . . . . . . . . . . . . . . . . . 47 5.4 Discovery by Optical-UV Properties . . . . . . . . . . . . . . . . . . . . . 51 6 The X-Ray Regime 54 6.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 54 6.2 Probing the Innermost Regions . . . . . . . . . . . . . . . . . . . . . . . . 55 6.3 The X-Ray Spectrum of AGNs . . . . . . . . . . . . . . . . . . . . . . . . 56 6.4 Lineless AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59 6.5 The Central Obscuration . . . . . . . . . . . . . . . . . . . . . . . . . . . 61 6.6 Detection and Observations of AGN in X-Rays . . . . . . . . . . . . . . . 62 6.6.1 X-Ray Observations of AGNs . . . . . . . . . . . . . . . . . . . . . 62 6.6.2 Discovery by X-Ray Properties . . . . . . . . . . . . . . . . . . . . 62 7 The -Ray Regime 64 7.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 64

16. Contents 7.2 Gamma-Ray Loud AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . 67 7.3 -Ray Properties of Blazars . . . . . . . . . . . . . . . . . . . . . . . . . . 67 8 The Uni

17. ed Model of AGNs 70 8.1 The Uni

18. cation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70 8.2 Absorbed Versus Unabsorbed AGN . . . . . . . . . . . . . . . . . . . . . . 72 8.3 Radio-Loud Versus Radio-Quiet . . . . . . . . . . . . . . . . . . . . . . . . 78 8.4 Breaking the Uni

19. cation . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83

20. List of Figures 1.1 The spectrum of NGC 1275 . . . . . . . . . . . . . . . . . . . . . . . . . . 4 1.2 The visible spectrum of Mrk 1157 . . . . . . . . . . . . . . . . . . . . . . . 4 1.3 The visible spectrum of 3C 273 . . . . . . . . . . . . . . . . . . . . . . . . 5 1.4 The total intensity distribution of 3C 338 . . . . . . . . . . . . . . . . . . 7 1.5 The total intensity distribution of 3C 173P1 . . . . . . . . . . . . . . . . . 7 1.6 The X-ray image of 3C 273's jet . . . . . . . . . . . . . . . . . . . . . . . 8 1.7 The UV spectrum of NGC 4594 . . . . . . . . . . . . . . . . . . . . . . . . 9 1.8 The spread of emission-line galaxies from the SDSS . . . . . . . . . . . . . 10 1.9 Radio luminosity vs. optical luminosity . . . . . . . . . . . . . . . . . . . 11 2.1 Comparison of a synchrotron source with a blackbody source . . . . . . . 15 3.1 Composite spectrum of Type I AGNs . . . . . . . . . . . . . . . . . . . . . 29 3.2 HST image of NGC 4261 . . . . . . . . . . . . . . . . . . . . . . . . . . . 31 3.3 AGN spectrum continuum . . . . . . . . . . . . . . . . . . . . . . . . . . . 32 4.1 VLA map of 3C 449 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 38 4.2 VLA map of 3C 47 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 39 4.3 Electromagnetic out ows from an accretion disk . . . . . . . . . . . . . . 41 4.4 Contour images of Cygnus A's jet . . . . . . . . . . . . . . . . . . . . . . . 42 4.5 Superluminal motion of M87's jet . . . . . . . . . . . . . . . . . . . . . . . 43 5.1 Composite optical-UV spectra of AGNs . . . . . . . . . . . . . . . . . . . 46 5.2 General view of a typical optical-UV SED of AGNs . . . . . . . . . . . . . 46 5.3 Broadband SEDs of AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . 48 5.4 Average optical-UV SED for Type I AGNs . . . . . . . . . . . . . . . . . 48 5.5 Spectrum of LLAGN NGC 5252 . . . . . . . . . . . . . . . . . . . . . . . 49 5.6 Comparison of dierent broad-line pro

21. les in Type I AGNs . . . . . . . . 50 5.7 u-g color of a large number of SDSS AGNs with various redshifts . . . . . 52 5.8 Discovering AGNs by their broadband colours . . . . . . . . . . . . . . . . 53 6.1 Composite AGN spectrum in extreme UV based on FUSE data . . . . . . 57 6.2 Soft X-ray spectrum of NLS1 Arkelian 564 . . . . . . . . . . . . . . . . . . 58 6.3 Composite spectrum of 15 lineless AGNs with large X-ray-to-optical lu-minosity . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 60 7.1 Multiepoch, multiwavelength spectrum of 3C 279 . . . . . . . . . . . . . . 66 8.1 Schematic representation of uni

22. ed BL Lac phenomenon . . . . . . . . . . 71 8.2 Schematic representation of the uni

23. ed AGN model . . . . . . . . . . . . 82 vii

24. List of Figures 8.3 Anticorrelation between X-ray variability amplitude and black hole mass . 84

25. List of Tables 2.1 Synchrotron sources in AGNs . . . . . . . . . . . . . . . . . . . . . . . . . 16 8.1 The general uni

26. cation scheme of AGNs . . . . . . . . . . . . . . . . . . . 83 ix

27. Abbreviations AGN Active Galactic Nuclei SMBH Super Massive Black Hole QSO Quasi Stellar Objects IRAS Infrared Astronomical Satellite NLRG Narrow Line Radio Galaxies BLRG Broad Line Radio Galaxies WLRG Weak-Emission Line Radio Galaxies BLR Broad Line Region SSRQ Steep Spectrum Radio Quasars FSRQ Flat Spectrum Radio Quasars FR-I Fanaro Riley Type I FR-II Fanaro Riley Type II BL Lac BL Lacertae OVV Optically Violently Variable Quasars LINER Low Ionization Nuclear Emission-Line Region LLAGN Low Luminosity Active Galactic Nuclei SED Spectral Energy Distribution SF Star Formation RIAF Radiatively Inaccurate Accretion Flow SSC Synchrotron Self Compton BH Black Hole NIR Near Infrared MIR Mid Infrared FIR Far Infrared RM Rotation Measure x

28. Abbreviations IC Inverse Compton UV Ultra-Violet HST Hubble Space Telescope MHD Magnetohydrodynamics FWHM Full Width at Half Maximum S/N Signal to Noise Ratio NLS1 Narrow Line Seyfert Type I SDSS Sloan Digital Sky Survey BAL Broad Absorption Line BEL Broad Emission Line XRB X-Ray Binary SAS Small Astronomy Satellite OSO Observing Solar Observatory HEAO High Energy Astronomy Observatory 2MASS 2 Micron All Sky Survey XMM X-Ray Multi-Mirror Mission RGS Re ecting Grating Spectrometer CCD Charge Coupled Device ESA European Space Agency NASA National Aeronautics and Space Agency COSMOS Cosmic Evolution Survey EW Equivalent Width HIG Highly Ionized Gas BAT Burst Alert Telescope ROSAT Roentgen Satellite INTEGRAL International Gamma-Ray Astrophysics Laboratory CGRO Compton Gamma-Ray Observatory LAT Large Area Telescope VLBI Very Long Baseline Interferometry HESS High Energy Spectroscopic System MERLIN Multi-Element Radio Linked Interferometer Network VLA Very Large Array HBLR Hidden Broad Line Region

29. Abbreviations OSSE Oriented Scintillation Spectrometer Experiment EXOSAT European X-Ray Observatory Satellite PDS Planetary Data System HBL High-Frequency Peaked BL Lac Objects LBL Low-Frequency Peaked BL Lac Objects RMS Root Mean Squared

30. Chapter 1 Introduction 1.1 The History of AGNs Unusual activity in the nuclei of galaxies was

31. rst recognised by Minkowski and Humason (Mount Wilson Observatory), when in 1943 they asked a graduate student Carl Seyfert to study a class of galaxy with an emission spectrum from the compact bright nucleus. Most normal galaxies show a continuum with absorption lines, but the emission in the Seyfert galaxies betrayed the presence of hot tenuous gas. In some cases the emission lines were broad (Type 1 Seyferts) indicating gas moving with high velocities and in other objects, the emission lines were narrow (Type 2 Seyferts) indicating that the gas was moving more slowly. In the 1950s, as radio astronomy became a rapidly developing science a whole new range of discoveries were made in astronomy. Amongst these were the Radio Galaxies, which appeared to be elliptical galaxies that were inconspicuous at optical wavelengths but were shown to have dramatically large, prominent lobes at radio frequencies, stretching for millions of light years from the main galaxy 1.2 Active Galactic Nuclei The names active galaxies and active galactic nuclei (AGNs) are related to the main feature that distinguishes these objects from inactive (normal or regular) galaxies |the presence of accreting supermassive black holes (SMBHs) in their centers. As of 2011, there are approximately a million known sources of this type selected by their color and 1

32. Chapter 1. Introduction several hundred thousand by basic spectroscopy and accurate redshifts. It is estimated that in the local universe, at z 0.1, about 1 out of 50 galaxies contains a fast-accreting SMBH, and about 1 in 3 contains a slowly accreting SMBH. Detailed studies of large samples of AGNs, and the understanding of their connection with inactive galaxies and their redshift evolution, started in the late 1970s, long after the discovery of the

33. rst quasi-stellar objects in the early 1960s. Although all objects containing active SMBH are now referred to as AGNs, various other names, relics from the 1960s, 1970s, and even now, are still being used. The most powerful active galaxies were discovered with radio telescopes in the 1960's and named `Quasi-Stellar Radio Sources', later shortened to QSOs or quasars. Their huge luminosities ( 104246 erg s1) could not be attributed to starlight alone, and the rapid variability observed (from months down to days) implied that the radiation was emitted from very small volumes with characteristic linear size of the order of light days. At the time, it was proving dicult to reconcile these two properties. As more detailed observations were performed it became clear that AGNs were most likely powered by accretion of matter onto a central SMBH (105 M

34. ). It is considerable to add that not all galaxies are active. Our Milky-Way is one of the numerous galaxies that hosts a SMBH at its galactic center (Schodel et al., 2002), with MSMBH 4:6 0:7 106M

35. , but is not considered to be an active galaxy due to the fact that there is no apparent accretion on to the SMBH. On contrary, the central regions of an AGN are likely not static, but very dynamic and violent. 1.3 The Taxonomy of AGNs The observational classi

36. cation of AGNs is not so clear because of observational limi-tations, heavy source obscuration (in most cases) and usually varying accretion rate on many orders of magnitude. Classically, an object is classi

37. ed as an AGN if :- It contains a compact nuclear region emitting signi

38. cantly beyond what is expected from stellar processes typical of this type of galaxy. It shows the clear signature of a non-stellar continuum emitting process in its center. 2

39. Chapter 1. Introduction Its spectrum contains strong emission lines with line ratios that are typical of excitation by a non-stellar radiation

40. eld. It shows line and/or continuum variations. 1.3.1 Seyferts Owing the name to Seyfert (1943) who was the

41. rst to discover these types, the major-ity of AGN with visible host galaxies fall under this class, known as Seyfert Galaxies. Seyfert, in his

42. rst observation, had reported a small percentage of galaxies had very bright nuclei that were the source of broad emission lines produced by atoms in a wide range of ionization states. These nuclei were nearly stellar in appearance (no powerful telescopes at that time were available). Today, these are further divided into two more subcategories :- Type I Seyferts: Spectra contain very broad emission lines that include both allowed lines (H I, He I, He II) and narrower forbidden lines (O [III]). They generally also have narrow allowed lines albeit being comparatively broader than those exhibited by non-active galaxies. The width of these lines is attributed to Doppler broadening, indicating that the allowed lines originate from sources with speeds typically between 1000 and 5000 km s1 Type II Seyferts: Spectra contain only narrow lines (both permitted and forbid-den), with characteristic speeds of about 500 km s1 1.3.2 Quasars and QSOs The terms Quasar (Quasi Stellar Radio Source) and QSO (Quasi Stellar Object), often used interchangeably, are scaled up versions of a Type I Seyfert, where the nucleus has a luminosity MB 21:5 + 5 log h0 Schmidt Green (1983). Maarten Schmidt recognized that the pattern of the broad emission lines of 3C 273 (the

43. rst detected quasar) was the same as the pattern of the Balmer lines of Hydrogen, but were severely redshifted to z = 0.158 to unfamiliar wavelengths, thus alluding astronomers from understanding it. 3

44. Chapter 1. Introduction Figure 1.1: The spectrum of NGC 1275. The emission features seen at 5057 A and 6629 A are [O III] 5007 and H, respectively. (Sabra et al., 2000) Figure 1.2: The visible spectrum of Mrk 1157, a Seyfert 2 galaxy. (Osterbrock, 1984) 4

45. Chapter 1. Introduction In 1963, the Dutch astronomer Maarten Schmidt recognized that the pattern of the broad lines of 3C 273 was the same as the pattern of the Balmer lines of Hydrogen, only severely redshifted to z = 0:158, hence alluding astronomers from identifying its spectrum. The continuous spectrum of a quasar may span nearly 15 orders of magnitude in frequency, very broad compared with the sharply peaked blackbody spectrum of a star. Quasars emit an excess of UV light relative to stars and so are quite blue in appearance. This UV excess is indicated by the big blue bump in (nearly) every quasar spectrum. A quasar's radio emission may come either from radio lobes or from a central source in its core. Figure 1.3: The visible spectrum of 3C 273, a Quasar. (Francis et al., 1991) 1.3.3 Radio Galaxies These galaxies are very luminous at radio wavelengths, with luminosities up to 1039 W between 10 MHz and 100 GHz. The observed structure in radio emission is determined by the interaction between twin jets and the external medium, modi

46. ed by the eects of relativistic beaming. These are further subdivided into two categories. 5

47. Chapter 1. Introduction 1.3.3.1 Radio Quiet Similar in many aspects to Type I Seyferts, these galaxies show both broad and narrow lines, the only dierence being that they are much more luminous than Type I Seyferts. They are observed in the absence of relativistic jets, which contribute the most energies in the radio wavelength spectrum. Radio Quiet Type I AGNs: These have relatively low-luminosities and therefore are seen only nearby, where the host galaxy can be resolved, and the higher-luminosity radio-quiet quasars, which are typically seen at greater distances because of their relative rarity locally and thus rarely show an obvious galaxy surrounding the bright central source. Radio Quiet Type II AGNs: These include Seyfert II galaxies at low luminosities, as well as the narrow-emission-line X-ray galaxies (Mushotzky, 1982). The high-luminosity counterparts are not clearly identi

48. ed at this point but likely candidates are the infrared-luminous IRAS AGN (Hough et al., 1991, Sanders et al., 1989, Wills et al., 1992), which may show a predominance of Type II optical spectra. 1.3.3.2 Radio Loud Usually attributed to AGNs with unipolar/bipolar, relativistic jets beaming out of their centers, the radio emission from radio-loud active galaxies is synchrotron emission, as inferred from its very smooth, broad-band nature and strong polarization. This implies that the radio-emitting plasma contains, at least, electrons with relativistic speeds (Lorentz factors of 104) and magnetic

49. elds. However, synchrotron radiation not being unique to radio wavelengths, if the radio source can accelerate particles to high enough energies, features which are detected in the radio may also be seen in the infrared, optical, ultraviolet or even X-ray. Radio Loud Type I AGNs: These are called Broad-Line Radio Galaxies (BLRG) at low luminosities and radio-loud quasars at high luminosities, either Steep Spectrum Radio Quasars (SSRQ) or Flat Spectrum Radio Quasars (FSRQ) depending on radio continuum shape. 6

50. Chapter 1. Introduction Radio Loud Type II AGNs: Often called Narrow-Line Radio Galaxies (NLRG), these include two distinct morphological types: the low-luminosity Fanaro-Riley type I (Figure 1.4) radio galaxies (Fanaro Riley, 1974), which have often-symmetric radio jets whose intensity falls away from the nucleus, and the high-luminosity Fanaro-Riley type II (Figure 1.5) radio galaxies, which have more highly collimated jets leading to well-de

51. ned lobes with prominent hot spots. Figure 1.4: The total intensity distribution of 3C 338, a FR I classi

52. ed AGN. (Ge Owen, 1994) Figure 1.5: The total intensity distribution of 3C 173P1, a FR II classi

53. ed AGN. (Leahy Perley, 1991) 7

54. Chapter 1. Introduction 1.3.4 Blazars Originally named after what was thought to be an irregular, variable star BL Lacertae, these are AGNs which are characterized by rapid and large-amplitude ux variability and signi

55. cant optical polarization. When compared to quasars with strong emission lines, blazars have spectra dominated by a featureless non-thermal continuum. The most well known object in this class is the BL Lacertae. Joining the BL Lac objects in the blazar classi

56. cation are the optically violently variable quasars (OVVs), which are similar to the BL Lacs except that they are typically much more luminous, and their spectra may display broad emission lines. Blazars are AGNs viewed head on and hence often have jets associated with them (Figure 1.6) Figure 1.6: The X-ray image of 3C 273's jet. (3C273 Chandra by Chandra X-ray Observatory - NASA. Licensed under Public domain via Wikimedia Commons) 1.3.4.1 BL Lacerate Objects BL Lacertae Objects, or BL Lacs for short, are a subclass of blazars that are characterized by their rapid time-variability. Their luminosities may change by upto 30% in just 24 hours and by a factor of 100 over a longer time period. BL Lacs are also distinguished by their strongly polarized power-law continua (30% 40% linear polarization) that are nearly devoid of emission lines, suggesting that there are very powerful EM

57. elds at play. BL Lacs, like quasars, are at cosmological distances. Of all the BL Lacs that have been resolved, 90% of those appear to reside in elliptical galaxies. 8

58. Chapter 1. Introduction 1.3.4.2 Optically Violent Variable Quasars Almost similar to BL Lacs, OVVs are typically much more luminous and may display broad emission lines in their spectra. The currently best known example of an OVV is 3C 279. 1.3.5 LINERs LINERs (Low Ionization Nuclear Emission-line Regions) are types of active galaxies that have very low luminosities in their nuclei, but with fairly strong emission lines of low-ionization species, such as the forbidden lines of [O I] and [N II]. The Spectra of LINERs seem similar to the low-luminosity end of the Seyfert II class, and LINER signatures are detected in many (most of) spiral galaxies in high-sensivity studies. These low-ionization lines are also detectable in starburst galaxies and in H II regions and hence it is sometimes dicult to distinguish between LINERs and starburst galaxies. In the local universe, they are found in about one-third of all galaxies brighter Figure 1.7: The UV spectrum of NGC 4594 LINER observed using the HST FOS. (Nicholson et al., 1998) than B = 15.5 mag. This is larger than the number of local high-ionization AGNs by a factor of 10 or more. Local high-ionization AGNs and LINERs are present in galaxies with similar bulge luminosities and sizes, neutral hydrogen gas (H I) contents, optical colors, and stellar masses. Given a certain galaxy type and stellar mass, LINERs are 9

59. Chapter 1. Introduction usually the lowest-luminosity AGNs, with nuclear luminosity that can be smaller than the luminosity of high-ionization AGNs by 1-5 orders of magnitude. An alternative name for this class of objects is low-luminosity AGNs (LLAGNs). The strongest Figure 1.8: The spread of emission-line galaxies from the SDSS on one diagnostic diagram that uses four strong optical emission lines, H, H

60. , [O III] 5007, and [N II] 6584, to distinguish galaxies that are dominated by ionization from young stars (green points) from those that are ionized by a typical AGN SED (blue points for high-ionization AGNs and red points for low-ionization AGNs). The AGN and SF groups are well separated, but the division between the two AGN groups is less clear. The curves indicate empirical (solid) and theoretical (dashed) dividing lines between AGNs and star-forming galaxies. (Groves Kewley, 2008) optical emission lines in the spectrum of LINERs include [O III] 5007, [O II] 3727, [O I] 6300, [N II] 6584, and hydrogen Balmer lines. All these lines are prominent also in high-ionization AGNs, but in LINERS, their relative intensities indicate a lower mean ionization state. For example, the [O III] 5007/H

61. line ratio in LINERs is 3-5 times smaller than in high-ionization Type-II AGNs. Line diagnostic diagrams are ecient tools to separate LINERs from high-ionization AGNs. One such example is shown in Figure 1.8. The exact shape of a LINERs SED is still an open issue. In some sources, it is well represented by the SED shown in Figure 5.3. Such an SED has a clear de

62. cit at UV wavelengths compared with the spectrum of high-ionization AGNs. However, some LINERs show strong UV continua and, occasionally, UV continuum variations, and it is not entirely clear what fraction of the population they represent. 10

63. Chapter 1. Introduction Figure 1.9: (left) Radio luminosity vs. optical (B-band) luminosity for various types of AGNs. (right) The radio loudness parameter R vs. (L=LEdd). (Sikora et al., 2007) This is related to the issue of Radiatively Inecient Accretion Flows (RIAFs) and the relationship between the mass-accretion rate onto the BH and the emitted radiation. Point-like X-ray sources have been observed in a large number of LINERs. These nuclear hard X-ray sources are more luminous than expected for a normal population of X-ray binaries and must be related to the central source. Many LINERs also contain compact nuclear radio sources similar to those seen in radio-loud high-ionization AGNs but with lower luminosity comparable to WLRGs (Figure 1.9). The UV-to-X-ray luminosity ratio in LINERs is, again, not very well known. In LINERs with strong UV continua, ox is smaller than in low-redshift, high-ionization AGNs, consistent with the general trend between ox and Lbol. However, ox is not known for most LINERs because of the diculty in measuring the UV continuum. Like other AGNs, LINERs can be classi

64. ed into Type-I (broad emission lines) and Type-II (only narrow lines) sources. The broad lines, when observed, are seen almost exclusively in H and hardly ever in H

65. . This is most likely due to the weakness of the broad wings of the Balmer lines that are dicult to observe against a strong stellar continuum. Some, perhaps many, LINERs may belong to the category of real Type-II AGNs |those AGNs with no BLR. The phenomenon is expected to be more common among low-luminosity sources and hence to be seen in LINERs. Because of all this, the classi

66. cation of LINERs is ambiguous, and the relative number of Type-I and Type-II objects of this class is uncertain even at very low redshift. 11

67. Chapter 2 Non-Thermal Processes Much of the electromagnetic radiation emitted by AGNs is very dierent from a simple blackbody emission or a stellar radiation source. The general name adopted here for such processes is non-stellar emission, but the term non-thermal emission is commonly used to describe such sources. There are several types of non-stellar radiation processes. 2.1 Basic Radiative Transfer Describing the interaction of radiation with matter requires the use of three basic quantities: the

68. rst is the speci

69. c intensity I, which gives the local ux per unit time, frequency, area, and solid angle everywhere in the source. The second quantity is the monochromatic absorption cross section, (cm1), which combines all loss (absorption and scattering) processes. The third quantity is the volume emission coecient, j, which gives the locally emitted ux per unit volume, time, frequency, and solid angle. The three are combined into the equation of radiative transfer, dI ds = I + j; where ds is a path length interval. The

70. rst term on the right in this equation describes the radiation loss due to absorption, and the second gives the radiation gain due to local emission processes. One usually de

71. nes the optical depth element, d = ds. Hence, 12

72. Chapter 2. Non-Thermal Processes dI d = I + S; where S = j= is the source function. The formal solution of the equation of transfer depends on geometry. For a slab of thickness in a direction perpendicular to the slab, it is I() = I(0)e + R 0 e(t)S(t)dt: For any other direction , both and dt must be divided by cos . The general equation of radiative transfer is dicult to solve and requires numerical techniques. However, there are simple cases in which the solution is straightforward. In particular, the case of a slab and a constant source function that is independent of allows a direct integration and gives the following solution: I = I(0)e + S(1 e ): For an opaque source in full thermodynamic equilibrium (TE), the optical depth is large, and both I and S approach the Planck function B(T) = 2h3=c2 eh=kT1 2.2 Synchrotron Radiation 2.2.1 Emission by a Single Electron in a Magnetic Field Considering an electron of energy E that is moving in a uniform magnetic

73. eld B of energy density uB = B2=8, the energy loss rate, dE=dt, which is also the power emitted by the electron, P, is given by P = 2T c 2

74. 2uB sin2 ; where T is the Thomson cross section, c is the speed of light, 13

75. Chapter 2. Non-Thermal Processes = E=mc2 is the Lorentz factor,

76. = v=c, and v is the speed of the electron. The angular term sin2 re ects the direction of motion, where is the pitch angle between the direction of the motion and the magnetic

77. eld. Averaging over isotropic pitch angles gives P = (4=3)T c 2

78. 2uB: The radiation emitted by a single electron is beamed in the direction of motion. The spectral energy distribution (SED) of this radiation is obtained by considering the gyro frequency of the electrons around the

79. eld lines (!B = eB= mec) and the mean interval between pulses (2=!B). The calculation of the pulse width is obtained by considering the relativistic time transformation between the electron frame and the observer frame. This involves an additional factor of 2. Thus, the pulse width is proportional to 3 or, expressed with the Larmor angular frequency, !L = eB=mec (which diers from !B by a factor of ), to 2. Fourier transforming these expressions gives the mean emitted spectrum of a single electron, P , which peaks at a frequency near 2!L. 2.2.2 Synchrotron Emission by a Power Law Distribution of Electrons Assuming now a collection of electrons with an energy distribution n( )d that gives the number of electrons per unit volume with in the range ( + d ), the emission coecient due to the electrons is obtained by summing P ( ) over all energies: j = 1 4 1 R 1 P( )n( )d : There is no general analytical solution to this expression since n( ) can take various dierent forms. However, there are several cases of interest where n( ) can be presented as a power law in energy: n( )d = n0 pd : The additional assumption that all the radiation peaks around a characteristic frequency, 2L, where L is the Larmor frequency, gives the following solution for j: 14

80. Chapter 2. Non-Thermal Processes 3T n0uB1 4j = 2 L L p1 2 : Figure 2.1: A comparison of a synchrotron source with p = 2:5 (solid line) and a 105 K blackbody source (dotted line). (Netzer, 2013) 2.2.3 Synchrotron Self-Absorption The source of fast electrons can be opaque to its own radiation. This results in a signi

81. cant modi

82. cation of the emergent spectrum especially at low frequencies, where the opacity is the largest. It can be shown that in this case, / p+4 2 ; that is, the largest absorption is at the lowest frequencies. Using the equation of radiative transfer for a uniform homogeneous medium, we get the solution at the large optical depth limit, I / 5=2, which describes the synchrotron SED at low energies. This function drops faster toward low energies than the low-energy drop of a blackbody spectrum (I / 2). The overall shape of such a source is shown in Figure 2.1. 2.2.4 Polarization Synchrotron radiation is highly linearly polarized. The intrinsic polarization can reach 70%. However, what is normally observed is a much smaller level of polarization, 15

83. Chapter 2. Non-Thermal Processes Source B (G) (Hz) tcool (yr) E (erg) Extended radio sources 105 109 104 107 1059 Radio jets 103 109 103 104 1057 Compact jets 101 109 102 101 1054 BH magnetosphere 104 1018 104 1010 1047 Table 2.1: Synchrotron Sources in AGNs. (Netzer, 2013) typically 3-15%. This indicates a mixture of the highly polarized synchrotron source with a strong non-polarized source. For AGNs, especially radio-loud sources, this polarization is clearly observed. There is also a correlation between high-percentage polarization and large-amplitude variations. AGNs showing such properties go under the name blazars. In the NIR-optical-UV spectrum of radio-loud AGNs, the region around 1 m shows most of the polarization. The percentage polarization seems to drop toward shorter wavelengths, in contrast to what is expected from a pure synchrotron source. This is interpreted as an indication of an additional thermal, non-polarized source at those wavelengths. 2.2.5 Synchrotron Sources in AGNs It is thought that most of the non-thermal radio emission in AGNs is due to synchrotron emission. There are various ways to classify such radio sources using the slope, (p 1)=2, and the break frequency below which it is optically thick to its own radiation. Table 2.1 gives a summary of the properties of several observed and expected synchrotron sources in AGNs. It includes the typical strength of the magnetic

84. eld, B (in gauss), the Lorentz factor, , and the total energy generated in the source, E, which is obtained by integrating uB over the volume of such sources. The table also shows the typical cooling time of the source, tcool, which is a characteristic lifetime de

85. ned by tcool = mec2 P ' 5 108B2 1sec: 16

86. Chapter 2. Non-Thermal Processes 2.2.6 Faraday Rotation Michael Faraday discovered in 1845 that the angle of polarization of an electromagnetic wave changes when the wave is sent through a medium with a magnetic

87. eld. The so-called Faraday rotation can also aect the synchrotron emission. Faraday rotation can be understood as the dierent eect the magnetized plasma has on the left and right circularly polarized light. Depending on the orientation with respect to the magnetic

88. eld, the components will see a dierent refractive index. Thus, the phase velocity of the two components will be aected slightly dierently and lead to a shift of their relative phases. This causes the plane of polarization to rotate, depending on how strong the magnetic

89. eld is and what distance the wave has to travel through the plasma. A similar eect is also observed with linearly polarized light. Once the linearly polarized synchrotron light is emitted and travelling towards the observer, it can pass through magnetized material causing Faraday rotation. This can be the emitting plasma itself, or any magnetized gas along the line of sight. In astrophysical applications, one can simplify the problem by considering only free electrons in magnetic

90. elds. The amount of rotation in the polarization angle depends on the magnetic

91. eld strength and density of the electrons along the line of sight, but also on the frequency of the electromagnetic wave one observes: = 2RM: Here, is the wavelength of the polarized radiation, and RM is the rotation measure which is a function of the electron density ne and of the component of the magnetic

92. eld Bjj parallel to the line of sight: = 2 e3 2m2c4 R ne(s)Bjj(s)ds: Thus, the rotation is larger for low frequencies. This is because the frequency of the wave is much larger than the gyro-frequency of the electron. The closer the light and the electron are to a resonant state, and thus the larger the energy transfer from the wave to the electron. The light from extragalactic sources will not only have to cross the intergalactic medium, but the interstellar medium of our galaxy as well on its path 17

93. Chapter 2. Non-Thermal Processes to the observer. The magnetic

94. eld along the line of sight will not be constant, and importantly, it will not be of the same orientation throughout the path of light. To determine the net eect of Faraday rotation, it is necessary to measure polarization at closely spaced frequency interval over many frequencies. Because the rotation aects the high frequency the least, the best way to get an estimate of the intrinsic polarization of a synchrotron source is to measure at high frequencies. 2.3 Thomson Scattering Thomson scattering describes the non-relativistic case of an interaction between an electromagnetic wave and a free charged particle. The eect was

95. rst describe by Sir Joseph John Thomson, who discovered the electron when studying cathode rays in the late nineteenth century. The process can be understood as elastic or coherent scattering, as the photon and the particle will have the same energy after the interaction as before. For this process of the energy E of the photon has to be much smaller than the rest energy of the particle: E = h mc2: Another requirement for Thomson scattering is that the particle must be moving at non-relativistic speed (v c). In the classical view of this process, the incoming photon is absorbed by the particle with charge q, which is set into motion and then re-emits a photon of the same energy. Using the classical electron radius r0 = q2=mc2 (Bohr radius), the dierential cross-section of this elastic scattering process can be written as d d = 1 2(1 + cos2 )r2 0: This is symmetric with respect to the angle , thus the amount of radiation scattered in the forward and backward direction is equal. The total cross-section is then given by T = 2 R 0 d d sin d = 8 3 r2 0 = 8 3 q2 mc2 2 : 18

96. Chapter 2. Non-Thermal Processes In the case of electrons, this gives a Thomson cross-section of T ' 6:652 1025 cm2. The cross-section for a photon scattering on a photon is a factor of (mp=me)2 ' 3:4 106 smaller. Since in the classical view of this process, the electron has no preferred orientation, the cross-section is independent of the incoming electromagnetic wave. The polarization of the scattered radiation depends, however, on the polarization of the incoming photon wave. Unpolarized radiation becomes linearly polarized in the Thomson scattering process with the degree of polarization being = 1cos2 1+cos2 : Therefore, polarization of the observed emission can be a sign that the emergent radiation has been scattered. Thomson scattering is important in may astrophysical sources. Any photon which will be produced inside a plasma can be Thomson scattered before escaping in the direction of the observer. The chance for the single photon to be Thomson scattered and how many of the photons will be scattered out of or into the line of sight is quanti

97. ed in terms of the optical depth of the plasma: = R T nedx; where ne is the electron density, and dx is the dierential line element. The mean free path T of the photon, that is, the mean distance traveled between scatterings will thus be T = (T ne)1. 2.4 Compton Scattering The interaction between an electron and a beam of photons is described by the classical Compton scattering theory. For stationary or slow electrons, one uses energy and momentum conservation to obtain the relationship between the frequencies of the coming ( 0 ) and scattered () photons. If ~n and ~n0 are unit vectors in the directions of these photons, and cos = ~n ~n0 , we get 19

98. Chapter 2. Non-Thermal Processes = mec2 0 mec2+h0 (1cos ) : For non-relativistic electrons, the cross section for this process is given by d d = 1 2r2 e [1 + cos2 ]; where re = e2=mec2 is the classical electron radius. Integrating over angles gives the Thomson cross section, T . In the high-energy limit, the cross section is replaced by the Klein-Nishina cross section, KN, which is normally expressed using = h=mec2. The approach to the low-energy limit is given roughly by KN T (1 2); and for 1, KN 3 8 T h ln 2 + 1 2 i : 2.4.1 Comptonization The term Comptonization refers to the way photons and electrons reach equilibrium. The fractional amount of energy lost by the photon in every scattering is ' h mec2 = : Considering a distance r from a point source of monochromatic luminosity L in an optically thin medium where the electron density is Ne, the ux at this location is L=4r2, and the heating due to Compton scattering is HCS = R L 4r2NeT h h mec2 i d: The cooling of the electron gas is the result of inverse Compton scattering. Like Compton scattering, this process is a collision between a photon and an electron, except that in this case, the electron has more energy that can be transfered to the radiation

99. eld. In this case, the typical gain in the photon energy is a factor of 2 larger than the one considered earlier. This factor is obtained by

100. rst transforming to the electron's rest frame and then back to the laboratory frame. If x is the fraction of the electron energy kT which is transferred to the photon, 20

101. Chapter 2. Non-Thermal Processes * + = x kTe mec2 ; where Te is the electron temperature. Using this terminology, one can write the cooling term for the electron gas as CCS = R L 4r2NeT h xkTe mec2 i d: A simple thermodynamical argument suggests that if Compton heating and Compton cooling are the only heating-cooling processes, and if the radiation

102. eld is given by the Planck function (L = B), the equilibrium requirement, HCS = CCS, gives x = 4. Because this is a general relation between a physical process and its inverse, the result must also hold for any radiation

103. eld. The radiation

104. eld in luminous AGNs can be very intense, and the energy density of the photons normally exceeds the energy density due to electrons. The requirement HCS = CCS gives, in this case, a Compton equilibrium temperature of TC = h 4k ; where the mean frequency, , is de

105. ned by integrating over the SED of the source, = R RLd Ld : 2.4.2 The Compton Parameter The emitted spectrum of thermal and non-thermal radiation sources that are embedded in gas with a thermal distribution of velocities is modi

106. ed due to Compton and inverse Compton scattering. For high-energy electrons, inverse Compton is the dominant process, and the resulting collisions will up-scatter the photon energy. The emergent spectrum is modi

107. ed, and its spectral shape will depend on the original shape, the electron temperature, and the Compton depth of the source, which determines the number of scattering before escape. Considering an initial photon energy of hi and the case of thermal electrons with temperature Te such that hi 4kTe, the scattering of such photons by a fast electron will result in energy gain 21

108. Chapter 2. Non-Thermal Processes per scattering (inverse Compton scattering). The photon continues to gain energy, during successive scatterings, as long as hi 4kTe. If the

109. nal photon energy is hf , and the number of scatterings is N, we get hf ' hie h N 4kTe mec2 i : For a medium with Compton depth e, the mean number of scatterings is roughly max(e; 2 e ). Using this, one can de

110. ne a Compton parameter y, y = max(e; 2 h 4kTe mec2 e ) i ; such that hf hiey: The factor ey is an energy ampli

111. cation factor. For y 1, one is in the regime of unsaturated inverse Comptonization. For y 1, the process reaches a limit where the average photon energy equals the electron thermal energy. This is saturated Compton scattering. 2.4.3 Inverse Compton Emission An important example is the case of a source whose spectrum is due to scattering of soft photons onto relativistic electrons. Again, considering

112. rst the typical energy following a single scattering and then averaging over the energy distribution of the photons and electrons, a simple way to estimate the power emitted in the preceding process is to consider a beam of photons with number density nph and mean energy before scattering h0. The energy density of these photons is nphh0, and the energy ux of photons incident on a stationary electron is curad = nphh0c. The mean energy after scattering, h, is larger than the mean energy before scattering by a factor of order 2. In the rest frame of the electron, the process can be considered as a simple Thomson scattering with radiated power given by the classical expression P = T curad. Thus, the simple L / (p1)=2 estimate for the laboratory frame emitted power is P = 2T curad. A more accurate derivation of the emitted power must take 22

113. Chapter 2. Non-Thermal Processes into account the scattering angle and its transformation between frames. The

114. nal expression in this case is P = (4=3)T c 2

115. 2urad; which diers from the simple estimate by a factor of order of unity. The expression for the power emitted due to inverse Compton (IC) scattering is basically identical to the power emitted by synchrotron radiation, except that the energy density of the magnetic

116. eld, uB, was replaced by the energy density of the radiation

117. eld, urad. Thus, the mean power of the two processes, assuming they take place in the same volume of space, is simply uB=urad. Also, for the same volume of space, the energy distribution of the relativistic electrons is given by the same power-law function used in the synchrotron case, n( ) / p. Thus, one also gets a similar dependence of the monochromatic luminosity on the parameter p: L(IC) / (p1)=2: 2.4.4 Synchrotron Self-Compton In a compact synchrotron source, the emitted photons can be inverse Compton scattered by the relativistic electrons that emit the synchrotron radiation. This gives the photon a big boost in energy. The emergent radiation is synchrotron self-Compton (SSC) emission. The ux emitted by this process can be calculated by integrating over the synchrotron radiation spectrum and the electron velocity distribution. To a good approximation, the resulting spectral index is identical to the spectral index of the synchrotron source. The synchrotron self-Compton process can repeat itself, in the same source, by additional scattering of the emergent photons, which results in an additional boosting, by a factor 2, to the photons. The natural limit for the process is when the scattered photon energy extends into the -ray and the condition of h mec2 (the condition for no Compton recoil of the electron) no longer holds. At this limit, the resulting radiation density decreases dramatically. 23

118. Chapter 2. Non-Thermal Processes 2.5 Annihilation and Pair-Production The observations of -ray jets in many AGNs suggest that, under some conditions, the density of high-energy photons is large enough to result in ecient pair production and a concentration of both electrons and positrons in some parts of the central source. Under these conditions, energetic -ray photons, with energies much above the rest energy of the electron, can react with lower-energy photons to create electron-positron pairs. Short-time-scale variations of the X-ray spectrum, in the lower-luminosity AGN, indicate extremely small dimensions; -ray photons that are associated with the X-ray source would not be able to escape these regions and would create electron-positron pairs. Likely locations where such processes take place are in the corona of the central accretion disk or inside the -ray jet. The process of pair-production and its reverse process (for e e+ pair) is given by e + e+ + : Considering the interaction between a -ray photon with frequency , above the rest mass frequency of the electron, with an X-ray photon of frequency X below this frequency and using the notation of unit vectors for the photons, one can write the threshold frequency for pair production as v = mec2 h 2 2 X(1~n ~nX) : The cross section is given by = 3 16T (1

119. 2) h (3

120. 4) ln 1+

121. 1

122. 2

123. (2

124. 2) i ; where the value of

125. for the electron and the positron is measured in the center of momentum frame. The typical value of near threshold is 0:2T , and it declines with frequency as 1 . The size of the radiation source, R, plays an important role in determining the optical depth of the source and hence the probability of pair-production taking place. This dependence is usually described by de

126. ning a compactness parameter for the -ray source, l , using the source size and its luminosity, L . There is an equivalent 24

127. Chapter 2. Non-Thermal Processes compactness parameter for the X-ray source, lX. Assuming that the typical -ray photon energy is = mec2 and that the photon number density is N = L 4R2c : The mean free path of the photons for the pair-production is = N T 1 and for unit optical depths, R , which gives L T 4mec3R 1: This leads to the following expression for the compactness parameter: l = L T 4mec3R; which is equivalent to the pair production optical depth of the source. In principle, l can be measured from the variability time scale of the -ray source. In reality, however, this is dicult to measure and is occasionally replaced by lX and the X-ray variability time scale. When lX hX=mec2, it will be dicult for the -rays to escape the source without creating pairs. The rate of the inverse process, pair annihilation, in the non-relativistic limit is independent of temperature and is roughly 0.4 NeT c per unit volume, where Ne is the combined electron-positron density. In a steady state, pair production is balanced by annihilation, mec3l 4TR2h = 0:4NeTc; where l is the compactness parameter for those -ray photons for which the source is optically thick to pair production. This equation can be solved for the mean Thomson depth in the source, T . For large T , the electrons and positrons thermalize because their interaction time is short compared with the annihilation time. In AGN gas, where the conditions allow this thermalization, the temperature of the hot, Compton thick pair plasma can reach 109 K. Such gas can contribute to the observed high-energy spectrum. It can up-scatter soft (UV) emitted photons and even produce some free-free electron-positron radiation. 25

128. Chapter 2. Non-Thermal Processes 2.6 Bremsstrahlung (Free-Free) Radiation Free-free radiation, formally, is thermal radiation. However in the case of AGNs, the spectral shape is very dierent from that of a blackbody. The free-free emissivity due to ion i of an element of charge Z whose number density is Ni is given by 4j = 6:8 1038Z2T1=2 e NeNigff (; Te;Z)eh=kTe ; where gff (; Te;Z) is the velocity-averaged Gaunt factor, which accounts for quantum-mechanical eects. This factor is always of the order unity and can change slightly with frequency, in particular, at X-ray energies gff / 0:1. The Bremsstrahlung radiation extends over a large range of energies and resembles, over most of this range, a very at (small spectral index) power law. One can integrate the free-free emissivity over frequencies to obtain the total energy per unit volume per second, Cff , where C indicates that this is also the cooling rate due to free-free emission. The integration gives Cff = 1:42 1027Z2T1=2 e NeNigffNeNi erg s1 cm3; where gff is now the frequency average of the velocity-averaged Gaunt factor. This is typically in the range 1.1-1.5. 26

129. Chapter 3 The IR and Sub-mm Regime 3.1 History The use of IR techniques to measure AGN continua started in the 1970s with the advent of the

130. rst sensitive IR detectors (Low Kleinmann, 1968). However, the IR colours of Seyfert galaxies are only subtly dierent than those of normal galaxies (Kuraszkiewicz et al., 2003), and the equivalent widths of the IR lines are not sucient to use as a

131. nding mechanism. Thus, IR color surveys can have a large fraction of false AGN, unless great care is taken. The

132. rst large-scale attempt to

133. nd AGN in the IR was based on Infrared Astronomical Satellite (IRAS) data. de Grijp et al. (1987) showed that AGN had systematically dierent 60 m / 25 m colours than normal galaxies. An alternative approach Spinoglio Malkan (1989) was to obtain optical spectra of every IR-selected galaxy. This was a follow-up of the idea of Huchra Burg (1992) to obtain optical spectra of every optically-selected galaxy, but was not really a survey technique. The latest use of the IR to

134. nd active galaxies is with the Two Micron All Sky Survey (2MASS; Cutri et al. (2002)). In this survey, 60% of the objects with J - K 2 are found to have the optical properties of AGN. This selection criterion is bootstrapped by using the near-IR colors of known radio and optically-selected AGN (Elvis et al., 1994), and thus will tend to

135. nd objects with similar properties. The large space density of these IR-selected objects makes them a major contributor to the AGN population. 27

136. Chapter 3. The IR and Sub-mm Regime The far-IR (FIR) band of thousands of AGNs has been observed by IRAS, with limited spatial resolution, and by Spitzer, with much improved resolution. The 2009 launch of Herschel is the most recent development in this area. Broadband images with much improved spatial resolution are now available between 70 and 500 m. Systematic surveys have already produced high-quality photometry of hundreds of AGNs and their host galaxies, up to redshift of 5 and beyond. Lower-sensitivity, high-resolution spectroscopy over the FIR range is also provided by the Herschel instruments. 3.2 Observations and Detection Most of the emission in the NIR and MIR bands is due to secondary dust emission. Secondary in this context refers to emission by cold, warm, or hot dust grains that are heated by the primary AGN radiation source.Primary refers to radiation that is the direct result of the accretion process itself. The temperature of the NIR- and MIR-emitting dust is between 100 and 2000 K. The dimensions of the dusty structure emitting this radiation, in intermediate luminosity AGNs, is of order 1 pc. Most of the thermal FIR emission is thought to be due to colder dust that is being heated by young stars in large star-forming regions in the host galaxy. In powerful radio sources, at least part of the FIR emission is due to non-thermal processes much closer to the center. Broad and narrow emission lines are seen in the NIR-FIR part of the spectrum of many AGNs. They are thought to originate in the broad- and narrow-line regions. A very important aspect techniques which use one IR band or a combination of two IR bands is the ability to detect highly obscured (Compton thick) AGNs. A large fraction of such objects, especially at high redshift, do not show detectable X-ray emission, and being type II sources, their optical spectrum is completely dominated by the host galaxy. Such sources would not be classi

137. ed as AGNs based on their optical and X-ray continuum properties. However, their mid-IR (MIR) spectrum is dominated by warm dust emission, the result of the heating of the central torus by the central source. A luminosity ratio like L(24 m)/L(R), where R is the red optical band, will be much larger in such sources compared with inactive galaxies because the AGN light is heavily obscured at the R-band. Spectroscopic follow-up of such objects can be used to look for the unique emission-line spectrum of the AGN. Indeed, systematic searches in uniformly scanned Spitzer

138. elds reveal a large number of Compton thick AGNs. 28

139. Chapter 3. The IR and Sub-mm Regime Figure 3.1 shows a composite 0.3-30 m spectrum of intermediate-luminosity type I AGNs. The emission longword of 1 m is due primarily to secondary radiation from dust. The dip at 1 m is due to the decline of the disk-produced continuum on the short-wavelength side and the rise of the emission due to hot dust on the other side. Figure 3.1: A composite spectrum of type-I AGNs covering the range 0:340 m. The observations were obtained by several ground-based telescopes and Spitzer and were normalized to represent a typical intermediate-luminosity source. (Netzer, 2013) 3.3 The Dusty Torus Dust is the cornerstone of the uni

140. cation theory of active galactic nuclei (AGNs). Essentially, all types of AGNs are surrounded by an optically thick dust torus and are basically the same object but viewed from dierent lines of sight (Antonucci, 1993, Urry Padovani, 1995). The large diversity in the observational properties of AGNs (eg., optical emission-line widths and X-ray spectral slopes) is simply caused by the viewing-angle-dependent obscuration of the nucleus: those viewed face-on are un-obscured (allowing for a direct view of their nuclei) and recognized as Type I Seyferts, while those viewed edge-on are Type II Seyferts, with most of their central engine and broad line regions being hidden by the obscuring dust. 29

141. Chapter 3. The IR and Sub-mm Regime Apparently, key factors in understanding the structure and nature of AGNs are determining the geometry of the nuclear obscuring torus around the central engine and the obscuration (ie., extinction, a combination of absorption and scattering) properties of the circumnuclear dust. An accurate knowledge of the dust extinction properties is also required to correct for the dust obscuration in order to recover the intrinsic optical/UV spectrum of the nucleus from the observed spectrum and to probe the physical conditions of the dust-enshrouded gas close to the nucleus. The presence of an obscuring dust torus around the central engine was

142. rst indirectly indicated by the spectropolarimetric detection of broad permitted emission lines (characteristic of Type I Seyferts) scattered into our line of sight by free electrons located above or below the dust torus in a number of Type II Seyferts (Heisler et al., 1997, Tran, 2003) Direct evidence for the presence of a dust torus is provided by IR observations. The circumnuclear dust absorbs the AGN illumination and re-radiates the absorbed energy in the IR. The IR emission at wavelengths longward of 1 m accounts for at least 50% of the bolometric luminosity of Type II Seyferts. For Type I Seyferts, 10% of the bolometric luminosity is emitted in the IR. A near-IR bump (excess emission above the 2 10 m continuum), generally attributed to hot dust with temperatures around 1200-1500 K (near the sublimation temperatures of silicate and graphite grains), is seen in a few Type I Seyferts (Barvainis, 1987, Rodrguez-Ardila Mazzalay, 2006). Direct imaging at near- and mid-IR wavelengths has been performed for several AGNs and provides constraints on the size and structure of the circumnuclear dust torus (Elitzur, 2006). Spectroscopically, the 10 m silicate absorption feature and the 3.4 m aliphatic hydrocarbon absorption feature are widely seen in heavily obscured Type II Seyferts; in contrast, the 10 m silicate emission feature has recently been detected in a number of Type I Seyferts. To properly interpret the observed IR continuum emission and spectroscopy as well as the IR images of AGNs, it requires a good understanding of the absorption and emission properties of the circumnuclear dust. To this end, one needs to know the composition, size, and morphology of the dust - with this knowledge, one can use Mie theory (for spherical dust) to calculate the absorption and scattering cross sections of the dust from X-ray to far-IR wavelengths, and then calculate its UV/optical/near-IR obscuration as a function of wavelength, and derive the dust thermal equilibrium temperature (based on the energy balance between absorption and emission) as well as 30

143. Chapter 3. The IR and Sub-mm Regime its IR emission spectrum. This will allow us to correct for dust obscuration and constrain the circumnuclear structure through modeling the observed IR emission and images. The former is essential for interpreting the obscured UV/optical emission lines and probing the physical conditions of the central regions; the latter is critical to our understanding of the growth of the central SMBH. However, little is known about the dust in the circumnuclear torus of AGNs. Even our knowledge of the best-studied dust - the Milky Way interstellar dust - is very limited. Figure 3.2: A HST image of the gas and dust disk in the active galactic nucleus of NGC 4261. (Ngc4261 by Clh288 at en.wikipedia. Licensed under Public domain via Wikimedia Commons) 3.4 IR Spectra The value of spectral index () is (almost) constant in the IR region of the spectrum of an AGN, evident from Figure 3.3. The thermal IR bump is due to the emission from warm (T . 2000 K) dust grains. 31

144. Chapter 3. The IR and Sub-mm Regime Figure 3.3: A depiction of the typical features in the continuum observed for many AGNs. (Tengstrand et al., 2009) 3.4.1 The 1 m Minimum The existence of the IR bump longward of 1 m has led many authors to conclude that this emission must be thermal, as the required temperatures are in the right range (T . 2000 K) for hot dust in the nuclear regions. Sanders et al. (1988) have shown that a minimum in the SED at 1 m is a general feature of AGNs. The hottest dust has a temperature of 2000 K; at higher temperatures, dust grains sublimate. This upper bound of the temperature explains the constancy of the frequency where the NIR spectrum is the weakest, ie., at the Wien cut-o at a 2000 K blackbody. One can de

145. ne a `sublimation radius' as the minimum distance from the AGN at which grains of a given composition can exist. The dust grains closest to an AGN probably are graphite rather than silicate, as graphite has a higher sublimation temperature. The sublimation radius for graphite grains is r = 1:3L1=2 uv46T2:8 1500 pc; 32

146. Chapter 3. The IR and Sub-mm Regime where Luv46 is the central source UV luminosity in units of 1046 erg s1, and T1500 is the grain sublimation temperature in units of 1500 K (Barvainis, 1987). 3.4.2 IR Continuum Variability Clear evidence that the hot dust scenario for the origin of the IR continuum has some merit has been provided by the IR continuum variability characteristics. Unlike UV/optical variability with little if any time delay, the IR continuum shows the same variations as the UV/optical continuum, but with a signi

147. cant time delay. This is interpreted as a light-travel eect which occurs because of the separation between the UV/optical and IR continuum-emitting regions; whereas the UV/optical emission arises in a very compact region, the IR emission arises in dust that is far away from the central source. The variations occur as the emissivity of the dust changes in response to the UV/optical continuum that heats it. Within the sublimation radius, dust is destroyed. Farther out, however, it survives and is heated by the UV/optical radiation from the central source to approximately the equilibrium blackbody temperature. The IR continuum arises as this energy is re-radiated by the dust. In the FIR, the only AGNs that are found to vary are radio-loud sources. 3.4.3 The Submillimeter Break Observations of the FIR to sub-mm portion of AGN spectra have been made in a limited number of cases (Chini et al., 1989, Edelson Malkan, 1987, Hughes et al., 1993). These observations show that the sub-mm SED decreases rather sharply as one goes to longer wavelengths, so abruptly that in at least a few cases the spectral index longward of the sub-mm break must be less than the value of -2.5 expected in the case of a synchrotron self-absorbed spectrum (ie., F / v5=2). At these long wavelengths, a thermal spectrum can produce a cut-o this sharp because the emitting eciency of small grains is a sensitive function of frequency, Q / , typically with 2 (Draine Lee, 1984) so the emitted spectrum can have a very strong frequency dependence, F / 2+ . 33

148. Chapter 4 The Radio Regime 4.1 History The discovery of radio galaxies preceded the optical discovery of AGNs. It goes back to the late 1940s and the early 1950s. Many of these sources were later shown to have optical-UV spectra that are very similar to the various types of optically discovered AGNs. The main features of many such sources are single- or double-lobe structures with dimensions that can exceed those of the parent galaxy by a large factor and strong radio cores and/or radio jets in some sources that coincide in position with the nucleus of the optical galaxy. About 10 percent of all AGNs are core-dominated radio-loud sources. This provides an additional way to identify AGNs in deep radio surveys by correlating their radio and optical positions. Stars are extremely weak radio sources, and hence an optical point source that is also a strong radio source is likely to be a radio-loud AGN. The positional accuracy of optical and radio telescopes is one arcsec or better, and there is hardly any problem in verifying that the radio and optical emitters are one and the same source. Most of the early AGN samples were discovered in this way. A well-known example is the 3C radio sample, which includes some of the most powerful radio-loud, early-discovered AGNs such as 3C 48 and 3C 273. 34

149. Chapter 4. The Radio Regime 4.2 The Loudness of AGNs Like optically classi

150. ed AGNs, there are broad-line radio galaxies (BLRGs), the equivalent of the Type I sources; narrow-line radio galaxies (NLRGs), the spectroscopic equivalent of Type II AGNs; and even weak-line radio galaxies (WLRGs), the equivalent of LINERs. While most AGNs show some radio emission, there seems to be a clear dichotomy in this property. Hence, usually, the radio loudness parameter, R, is used to separate radio-loud from radio-quiet AGNs. R is a measure of the ratio of radio (5 GHz) to optical (B-band) monochromatic luminosity, R = L(5 Ghz) L(4400A ) = 1:5 105 L(5 Ghz) L(4400A ) ; where L(5 Ghz) and L(4400 A) represent the value of L at those energies. The dividing line between radio-loud and radio-quiet AGNs is usually set at R = 10. Statistics of a large number of AGNs show that about 10 percent of the sources are radio loud, with some indication that the ratio is decreasing with redshift. Much of the radio emission in radio-loud AGNs originates in a point-like radio core. The spectrum of such core-dominated radio sources suggests emission by a self-absorbed synchrotron source. Except for the self-absorption low frequency part, the spectrum is represented well by a single power law, F / R. Sources with R 0:5 are usually referred to as at-spectrum radio sources, and those with R 0:5 are steep-spectrum radio sources. There is a clear connection between the radio structure and the radio spectrum of such sources. Steep-spectrum radio sources show lobe-dominated radio morphology and are also less variables. Flat-spectrum sources have in general higher luminosity cores, larger amplitude variations, and weak or undetected lobes. This dichotomy is interpreted as a dependence on the viewing angle to the core. In steep-spectrum sources, one is looking away from the direction of the nuclear radio jet, and the radio emission is more or less isotropic. In at-spectrum sources, we are looking at a small angle into the core. The intensity is boosted due to the relativistic motion of the radio-emitting particles, and the variations are ampli

151. ed. In many cases, there is evidence for superluminal motion in such sources. 35

152. Chapter 4. The Radio Regime 4.3 The Fanaro-Riley Classi

153. cation It was

154. rst noticed by Fanaro Riley (1974) that the relative positions of regions of high and low surface brightness in the lobes of extragalactic radio sources are correlated with their radio luminosity. This conclusion was based on a set of 57 radio galaxies and quasars, from the complete 3CR catalogue, which were clearly resolved at 1.4 GHz or 5 GHz into two or more components. Fanaro and Riley divided this sample into two classes using the ratio RFR of the distance between the regions of highest surface brightness on opposite sides of the central galaxy or quasar, to the total extent of the source up to the lowest brightness contour in the map. Sources with RFR 0:5 were placed in Type I (FR-I) and sources with RFR 0:5 in Type II (FR-II). It was found that nearly all sources with luminosity L(178 Mhz) . 2 1025 h2 100 W Hz1 Sr1 were of Type I (FR-I) while the brighter sources were nearly all of Type II (FR-II). The luminosity boundary between them is not very sharp, and there is some overlap in the luminosities of sources classi

155. ed as FR-I or FR-II on the basis of their structures. For a spectral index of ' 1 the dividing luminosity at 5 GHz is L(5 Ghz) . 7 1023 h2 100 W Hz1 Sr1 At high frequencies, the luminosity overlap between the two classes can be as much as two orders of magnitude. Various properties of sources in the two classes are dierent, which is indicative of a direct link between luminosity and the way in which energy is transported from the central region and converted to radio emission in the outer parts. 4.3.1 Fanaro-Riley Class I (FR-I) Sources in this class have their low brightness regions further from the central galaxy or quasar than their high brightness regions (Figure 1.4 and Figure 4.1). The sources become fainter as one approaches the outer extremities of the lobes and the spectra here are the steepest, indicating that the radiating particles have aged the most. Jets are detected in 80% of FR-I galaxies. A jet can begin as one-sided close to the core, 36

156. Chapter 4. The Radio Regime but beyond a few kiloparsec it becomes two-sided and continuous, with an opening angle 8 that varies along its length. Along the jet the component of the magnetic

157. eld in the plane of the sky is at

158. rst parallel to the jet axis, but soon becomes aligned predominantly perpendicular to the axis. FR-I sources are associated with bright, large galaxies (D or cD) that have a atter light distribution than an average elliptical galaxy and are often located in rich clusters with extreme X-ray emitting gas (Owen Laing, 1989, Prestage Peacock, 1988) As the galaxy moves through the cluster the gas can sweep back and distort the radio structure through ram pressure, which explains why narrow-angle-tail or wide-angle-tail sources, say, appear to be derived from the FR-I class of objects. A typical FR-I galaxy is shown in Figure 4.1. This is the radio source 3C 449, which is optically identi

159. ed with a galaxy of type cDE4 at a redshift of 0.0181, so that 1 corresponds to 255 h1 100 pc. There are twin jets that are straight for 30 from the core, after which they deviate towards the west and terminate into diuse lobes. These jets and outer lobes are mirror symmetric about an axis through the core. The jets are generally smooth in appearance, but higher resolution observations show knots on a smooth ridge of emission, the southern jet being more knotty than the northern one. Within 10 of the nucleus, the surface brightness of the jets is much reduced. The jets widen at a non-uniform rate close to the core, with the greatest expansion occurring where the jets are faintest. Beyond 10 from the nucleus, the opening angle is constant at 7. The emission from the jets is highly polarized, the average polarization over the jets being 30%, and the projected magnetic

160. eld is perpendicular to the jet axis. 4.3.2 Fanaro-Riley Class II (FR-II) This class comprises luminous radio sources with hotspots in their lobes at distances from the center which are such that RFR 0:5. These sources are called edge-darkened, which was a terminology when the angular resolution and dynamic range used in observing the classical sources was not always good enough to reveal the hotspots as distinct structures. In keeping with the overall high luminosity of this type of source, the cores and jets in them are also brighter than those in FR-I galaxies in absolute terms; but relative to the lobes these features are much fainter in FR-II 37

161. Chapter 4. The Radio Regime Figure 4.1: VLA map of the FR-I galaxy 3C 449 at 1465 MHz, with angular resolution 4.8 3.4 arcsec2. The peak ux is 22.2 mJy per beam, with contours drawn at 5% intervals, beginning with the -5% contour. (Perley et al., 1979). galaxies. Jets are detected in 10% of luminous radio galaxies, but in nearly all quasars. The jets have small opening angles ( 4) and are knotty; the jet magnetic

162. eld is predominantly parallel to the jet axis except in the knots, where the perpendicular component is dominant. Figure 4.2 shows an example of an FR-II galaxy, which is a VLA map of the radio quasar 3C 47 made by Bridle et al. (1994). The most striking feature of the jets in the FR-II class is that they are often one-sided, as is clearly seen in Figure 4.2. Jet one-sidedness occurs at large (kpc) scales as well as in the milli-arcsecond jets which are found in compact cores through VLBI observations. The feature A in the jetted lobe is a hotspot, while feature H on the 38

163. Chapter 4. The Radio Regime unjetted side looks like one, but does not qualify for being a hotspot according to the criteria of Bridle et al. (1994). Figure 4.2: VLA map of the FR-II quasar 3C 47 made at 4.9 GHz with 1:45 1:13 arcsec2 resolution. G is the core, A the jetted hotspot. H does not meet the hotspot criteria of Bridle et al. (Bridle et al., 1994). FR-II sources are generally associated with galaxies that appear normal, except that they have nuclear and extended emission line regions. The galaxies are giant ellipticals, but not

164. rst-ranked cluster galaxies. The environment of FR-II sources does not show enhanced galaxy clustering over the environment of randomly chosen elliptical galaxies (Owen Laing, 1989, Prestage Peacock, 1988). Owing to the large dierences in the nature of the host galaxies and the environments of the FR-I and FR-II sources, it is possible that they are intrinsically dierent types of source not related to each other through an evolutionary sequence. 4.4 Radio Lobes and Jets Radio-loud sources usually consist of a radio core, one or two detectable jets, and two dominant radio lobes. The radio-quiet sources are less luminous at radio wavelengths by a factor of 103 to 104, consisting of a weak radio core and perhaps a feeble jet. The 39

165. Chapter 4. The Radio Regime increased level of activity in radio-loud AGNs is not con

166. ned to radio wavelengths, however; they also tend to be about three times brighter in X-rays than their radio-quiet counterparts. 4.4.1 The Generation of Jets The radio lobes are produced by jets if charged particles ejected from the central nucleus of the AGN at relativistic speeds. These particles are accelerated away from the nucleus in two opposite directions, powered by the energy of accretion and/or by the extraction of rotational kinetic energy from the SMBH via the Blandford-Znajek mechanism (Blandford Znajek, 1977). The jet must be electrically neutral overall, but it is not clear whether the ejected material consists of electrons and ions or an electron-positron plasma. The latter, being less massive, would be more easily accelerated. The disk's magnetic

167. eld is coupled (frozen in) to this ow of charged particles. The resulting magnetic torques may remove angular momentum from the disk, which would allow the accreting material to move inward through the disk. The incredible narrowness and straightness of some jets means that a collimating process must be at work very near the central engine powering the jet. A thick, hot accretion disk around the SMBH could provide natural collimation by funneling the out owing particles, as shown in Figure 4.3. Because the accreting material retains some angular momentum as it spirals inward through the disk, it will tend to pile up at the smallest orbit that is compatible with its angular momentum. Inside this centrifugal barrier, there may be a relatively empty cavity that can act as a nozzle, directing the accreting gases outward along the walls of the cavity. However, producing highly relativistic jets, as frequently observed, appear to be dicult to accomplish with this nozzle mechanism. Alternatively, magnetohydrodynamic (MHD) eects could play an important role in accelerating and collimating the relativistic ows. 4.4.2 The Formation of Radio Lobes As a jet travels outward, its energy primarily resides in the kinetic energy of the particles. However, the jet encounters resistance as it penetrates the interstellar 40

168. Chapter 4. The Radio Regime Figure 4.3: A Sketch of the electromagnetic out ows from the two sides of a rotating magnetized accretion disk owing to the unipolar dynamo action. (CYGAM (CYlindrical GAmma-ray Monitor), Russian Space Research Institute). medium within the host galaxy and the intergalactic medium beyond. As a result, the material at the head of the jet is slowed, and a shock front forms there. The accumulation and deceleration of particles at the shock front cause the directed energy of the jet to become disordered as the particles splash back to form a large lobe in which the energy may be shared equally by the kinetic and magnetic energy. The motion of the charged particles and magnetic

169. elds within the lobes of radio-loud objects contain an enormous amount of energy. For Cygnus A Figure 4.4, the energy of each lobe is estimated to be approximately 1053 to 1054 J, equivalent to energy liberated by 107 supernovae! 41

170. Chapter 4. The Radio Regime Figure 4.4: Contour images of the Cygnus A radio jet on various scales. (Carilli et al., 1996). 4.4.3 Accelerating the Charged Particles in the Jets The observations of jets are made possible by ineciencies in the transport of particles and energy out to the radio lobes. The spectra of the radio lobes and jets follow a power law, with a typical spectral index of ' 0:65. The presence of power-law spectra and a high degree of linear polarization strongly suggest that the energy emitted by the lobes and jets comes from synchrotron radiation. The loss of energy by synchrotron radiation is unavoidable, and the relativistic electrons in jets radiate away their energy after just 10,000 years or so. This implies that there is not nearly enough time for particles to travel out to the larger radio lobes. 42

171. Chapter 4. The Radio Regime This long travel time implies that there must be some mechanism for accelerating particles in the jets and radio lobes. As one possibility, shock waves may accelerate charged particles by magnetically squeezing them, re ecting them back and forth inside the shock. Radiation pressure may also play a role, but is alone not enough to generate the necessary acceleration. 4.4.4 Superluminal Velocities Although the standard model of jets and radio lobes requires a steady supply of charged particles moving at relativistic speeds, evidence for such high velocities is dicult to obtain. The absence of spectral lines in a power-law spectrum means that the relativistic velocity of the jet material cannot be measured directly but must be inferred from indirect evidence. The most compelling argument for relativistic speeds involves radio observations of material ejected from the cores of several AGNs, with so-called superluminal velocities. This eect is observed within about 100 pc of the AGN's center and probably continues farther out. Figure 4.5: The apparent superluminal motion of the M87 Jet. 43

172. Chapter 5 The Optical-UV Regime 5.1 History The

173. rst hint of the violent heritage of today's galaxies was found by Edward A. Fath (1908), who was observing the spectra of spiral nebulae. Although most showed an absorption-line spectrum produced by the combined light of the galaxy's stars, NGC 1068 displayed six bright emission lines. In 1926, Edwin Hubble recorded the emission lines of this and two other galaxies. Seventeen years later, Carl K. Seyfert reported that a small percentage of galaxies have very bright nuclei that are the source of broad emission lines produced by atoms in a wide range of ionization states. These nuclei were nearly stellar in appearance. 5.2 Spectrum Figure 3.3 is a rough schematic of the continuum observed for many types of AGNs. The most notable feature of this SED is its persistence over some 10 orders of magnitude in frequency. The wide spectrum is markedly dierent from the thermal (blackbody) spectrum of a star or the combined spectra of a galaxy of stars, and one can see that there are many non-thermal radiation processes that are going on at dierent stages in the AGNs. When AGNs were

174. rst studied, it was thought that their spectra were quite at. Accordingly, a power law of the form : F / was used to describe the 44

175. Chapter 5. The Optical-UV Regime monochromatic energy ux, F. The spectral index, , was believed to have a value of ' 1. The power received within any frequency interval between 1 and 2 is Linterval / 2 R 1 Fd = 2 R 1 F d = ln 10 2 R 1 Fd log10 ; so that equal areas under the graph of F vs. log10 correspond to equal amounts of energy. A value of ' 1 re ects the horizontal trend seen at the IR bump of

176. gure 3.3. The continuous spectra of AGNs are now known to be more complicated, involving a mix of thermal and non-thermal emission. However F / is still used to parameterize the continuum. typically has a value between 0.5 and 2 that usually increases with increasing frequency, so the curve of log10 L (or log10 F) vs. log10 in Figure 3.3 is generally concave downward. In fact, the value of is constant over only a limited range of frequencies, such as in the IR and visible regions of the spectrum. The shape and polarization of the visible-UV spectrum indicates that it can sometimes be decomposed into contributions from thermal sources (blackbody spectrum, low polarization) and non-thermal sources(power law spectrum, signi

177. cant polarization). The thermal component appears as the big blue bump in Figure 3.3, which can contain an appreciable amount of bolometric luminosity of the source. It is generally believed that the emission from the big blue bump is due to an optically thick accretion disk, although some believe that free-free emission may be responsible. 5.2.1 The Optical-UV Continuum and the Accretion Disk The best-understood disks are thin disks with or without X-ray emitting coronae. The majority of intermediate- and high-luminosity AGNs found in large surveys are thought to be powered by such disks. Thin-disk theory suggests that the SED of such systems contains a broad wavelength band where the spectral slope (L / ) is in the range 0-0.5. The classical thin-disk slope is = 1=3. X-ray emission from the hot corona and X-ray re ection from the surface of the disk are additional important characteristics of such systems. Much observational eort has been devoted to the measurement of and the characterization of the part of the continuum showing this 45

178. Chapter 5. The Optical-UV Regime Figure 5.1: A Composite Optical-UV Spectra of AGNs (Francis et al., 1991). Figure 5.2: A General View of the Optical-UV SED of AGNs. 46

179. Chapter 5. The Optical-UV Regime slope (the big blue bump). For MBH = 109M

180. and L=LEdd = 0:1, the peak of this emission is predicted to be around 1000 A. Even the most sophisticated calculated spectra are rather limited, and several models show spectra that dier considerably from the schematic 1=3 dependence and also, the infrared ( 1 m) part of the SED is dominated by non-disk emission, in particular, stellar emission by the host galaxy and thermal emission by warm dust, presumably in the central torus. In radio-loud AGNs, non-thermal emission can also contribute at this and even shorter wavelength bands. This obscures the part of the spectrum where the standard thin-disk theory predicts = 1=3. One must also consider the (yet hypothetical) possibility that additional processes, related perhaps to disk winds, are taking place and changing the observed SED. 5.3 Observations in the Optical-UV Region Optical images of luminous Type-I AGNs show clear signatures of point-like central sources with excess emission over the surrounding stellar background of their host galaxy. The non-stellar origin of these sources is determined by their SED shape and by the absence of strong stellar absorption lines. Type-II AGNs do not show such excess. The luminosity of the nuclear, non-stellar source relative to the host galaxy luminosity can vary by several orders of magnitude. In particular, many AGNs in the local universe are much fainter than their hosts, and the stellar emission can dominate their total light. For example, the V-band luminosity of a high-stellar-mass AGN host can approach 1044 erg s1, a luminosity that far exceeds the luminosity of many local Type-I AGNs. This must be taken into account when evaluating AGN spectra obtained with large-entrance-aperture instruments. The relative AGN luminosity increases with decreasing wavelength, and contamination by stellar light is not a major problem at UV wavelengths. The optical-UV spectra shown in Figure 5.4 and Figure 5.5 represent typical spectra of high-ionization luminous Type-I and Type-II AGNs. The added high-ionization is needed to distinguish such sources from low-ionization Type-I and Type-II sources. The Type-I spectrum is a composite composed of several thousand spectra of dierent redshift AGNs. This is done to illustrate the entire rest wavelength range of 900-7000 A using only ground-based observations. The data used to obtain this composite at 5000 A are based on spectra of lower luminosity, low-redshift 47

181. Chapter 5. The Optical-UV Regime Figure 5.3: Broadband spectral energy distributions (SEDs) for various types of AGNs. (Ho, 2008) Figure 5.4: The average optical-UV SED of several thousand high-luminosity Type I AGNs (Vanden Berk et al., 2001) 48

182. Chapter 5. The Optical-UV Regime objects, and the SED at those wavelengths is aected by host galaxy contamination. The Type-II spectrum covers a similar range, but this time, the spectrum is a combination of a ground-based optical spectrum with a space-borne (HST) UV spectrum. The striking dierences between the high-ionization Type-I and Type-II Figure 5.5: The spectrum of the low-luminosity, low-redshift type-II AGN NGC 5252. (Tsvetanov et al., 1996) spectra, which were the reason for the early classi

183. cation into Seyfert 1 and Seyfert 2 galaxies, are the shape and width of the strongest emission lines. Type-II AGNs show only narrow emission lines with typical full-width-at-half-maximum (FWHM) of 400 800 km s1. In Type-I spectra, all the permitted line pro

184. les, and a few semi forbidden line pro

185. les, indicate large gas velocities, up to 5000 10000 km s1 when interpreted as owing to Doppler motion. The line ratios and line widths of the 49

186. Chapter 5. The Optical-UV Regime Figure 5.6: Comparison of dierent broad-line pro

187. les in a typical Type-I AGN. (Netzer, 2013) forbidden lines in the spectra of Type-I sources are very similar to those observed in Type-II spectra and indicate that the basic physics in the narrow line-emitting region of both classes is the same. The broad emission lines can be used to map the gas kinematics very close to the central BH and to measure the BH mass. Study of the spectra of many thousand Type-I AGNs shows a considerable range in optical-UV continuum slope but little if any correlation between the slope and Lbol. Some of the observed dierences are attributed to a small amount of reddening in the host galaxy of the AGN or other sources of foreground dust. Although the spectra shown here clearly illustrate the large dierences in emission-line widths between Type-I and Type-II sources, observational limitations can make it dicult to detect weak broad emission lines. Slightly obscured or low-luminosity Type-I AGNs are occasionally classi

188. ed as Type-II based on their stellar-like continua and narrow emission lines. This can be the result of reddening of the broad wings of the H

189. line or a relatively strong stellar continuum, especially in large-aperture, low-spatial-resolution observations. Higher signal-to-noise (S/N), better-spatial-resolution observations of the same sources reveal, in some cases, very broad wings in one or more Balmer lines. The term broad emission lines, which is used to describe the permitted and semi-forbidden lines in Type-I AGNs, does not necessarily mean similar widths for all lines in all 50

190. Chapter 5. The Optical-UV Regime objects. The various broad emission lines show typically dierent widths, and in general, the width re ects the level of ionization of the gas, the source luminosity, and the mass of the central SMBH. Historically, it was found that broad emission lines in some Type-I AGNs are narrower than narrow emission lines in Type-II sources. A well-known example is the subgroup of narrow-line Seyfert 1 galaxies (NLS1s). This class of objects was historically de

191. ned as those Seyfert 1 galaxies with FWHM(H

192. ) 2000 km s1. In many such sources, FWHM(H

193. ) 1000 km s1, similar to the width of H

194. in many Type-II AGNs. Evidently, the distinction between Type-I and Type-II sources requires other criteria, such as the presence of a non-stellar continuum; strengths of emission lines typical of Type-I sources such as FeII emission lines; or the presence of a strong, unobscured X-ray continuum. There are also dierences in the shape and even the velocity of the same line in dierent objects. Some examples are shown in Figure 5.6. A similar remark should be made about the width of the narrow emission lines. For example, the width of the strong [O III] 5007 line can depend on the mass of the central SMBH (or, more accurately, the mass of the bulge). Thus [O III] 5007 lines with FWHM 1000 km s1 are commonly observed in high-redshift, large-MBH, large-Lbol AGNs. 5.4 Discovery by Optical-UV Properties As said, typical AGN SEDs are dierent in several ways from stellar SEDs, in which they cover a broader energy range and do not resemble a single-temperature blackbody. This dierence provides a simple and ecient way of discovering AGNs using broadband multicolor photometry. Several color combinations, based on three-band and

195. ve-band photometry, are useful in separating AGNs from stars by their color. Earlier methods were based on a UVB photometry in large areas of the sky. This three-band system is useful for discovering low-redshift sources but fails to detect many high-redshift objects because the spectrum gets eectively red and resembles the colors of nearby stars. In addition, even the low-redshift AGNs can be confused with the local population of hot white dwarfs. Some AGNs are intrinsically red, or reddened by dust, which results in colors that are not very dierent from those of stars. Moreover, intrinsically blue, high-redshift AGNs are eectively red due to the absorption of their short-wavelength radiation by intergalactic gas. Figure 5.7 illustrates this and 51

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