The Pennsylvania State University

The Graduate School

Department of Astronomy and Astrophysics

THE VIEW THROUGH THE WIND:

X-RAY OBSERVATIONS OF BROAD ABSORPTION LINE

QUASI-STELLAR OBJECTS

A Thesis in

Astronomy and Astrophysics

by

Sarah Connoran Gallagher

c 2002 Sarah Connoran Gallagher

Submitted in Partial Fulfillment of the Requirements for the Degree of

Doctor of Philosophy

May 2002 We approve the thesis of Sarah Connoran Gallagher.

Date of Signature

W. Nielsen Brandt Associate Professor of Astronomy and Astrophysics Thesis Adviser Chair of Committee

Jane C. Charlton Associate Professor of Astronomy and Astrophysics

Michael C. Eracleous Assistant Professor of Astronomy and Astrophysics

L. Samuel Finn Associate Professor of Physics

Gordon P. Garmire Evan Pugh Professor of Astronomy and Astrophysics

Steinn Sigurdsson Assistant Professor of Astronomy and Astrophysics

Peter I. M´esz´aros Professor of Astronomy and Astrophysics Head of the Department of Astronomy and Astrophysics iii Abstract

The 2–10 keV bandpasses and unprecedented sensitivity of modern X-ray obser- vatories have enabled new insights into the immediate environments of Broad Absorp- tion Line (BAL) QSOs. BAL QSOs, approximately 10% of the QSO population, exhibit deep, broad absorption lines from high ionization ultraviolet resonance transitions. These blueshifted absorption features are understood to arise along lines of sight which travel through radiatively driven winds with terminal velocities reaching 0.1–0.3c. These en- ergetic outflows are an important component of QSO environments; mass ejection is apparently fundamentally linked to the process of active mass accretion onto super- massive black holes. X-rays, generated in the innermost regions surrounding accreting black holes, travel through the nuclear environments to the observer. X-ray studies of BAL QSOs thus offer a privileged view through the wind. In this thesis, I present the results from X-ray surveys of BAL QSOs. This work includes both spectroscopic and exploratory X-ray observations designed to characterize the X-ray properties of specific objects as well as the population as a whole. Before this project, BAL QSOs were known to be X-ray weak relative to normal QSOs, however, the cause for this faintness had not been demonstrated. In general, I find that X-ray weak- ness in BAL QSOs results from intrinsic absorption of a typical QSO X-ray continuum. Fitting X-ray spectra with simple absorption models results in intrinsic column densities 22 2 ranging from (1–50) 10 cm− . Even at the spectral resolution currently available for × such observations, there is significant evidence for complexity in the absorber, perhaps arising from ionized gas with velocity structure that may partially cover the direct con- tinuum. From the X-ray continua above 5 keV, the X-ray to ultraviolet flux ratios ∼ are found to be consistent with normal QSOs, indicating that the spectral energy dis- tributions of BAL QSOs are not anomalous. For those BAL QSOs exhibiting BALs in low-ionization lines (e.g., Mg ii), the situation is qualitatively different. The direct line of sight to the primary X-ray continuum appears to be completely blocked by Compton- thick absorption. In this situation, only indirect X-rays, either scattered or reflected, are able to reach the observer. In the future, deeper observations of the brightest BAL QSOs as well as expanded exploratory surveys will offer additional insight into the state and location of the X-ray absorbers. Higher spectral resolution will clarify the physical state of the X-ray absorber as well as elucidate the relationship between the ultraviolet and X-ray absorbing gas. This information is essential for attaining the ultimate goal of constraining the mass outflow rate of luminous QSOs. iv Table of Contents

List of Tables ...... vii

List of Figures ...... viii

Acknowledgments ...... ix

Chapter 1. Introduction ...... 1 1.1 Active Galactic Nuclei ...... 1 1.1.1 The Standard Model ...... 1 1.1.2 Progress Towards Unification ...... 1 1.1.3 Radio-Loud/Radio-Quiet AGN ...... 2 1.2 Broad Absorption Line QSOs ...... 3 1.2.1 Continuum and Spectral Energy Distribution ...... 3 1.2.2 X-ray Properties (circa 1997) ...... 4 1.2.3 Broad Absorption Lines ...... 5 1.2.4 Polarization ...... 6 1.2.5 Variability ...... 6 1.2.6 Observational Subclasses ...... 7 1.2.7 Theoretical Models for BAL QSOs ...... 9 1.3 Motivation for this Project ...... 10 1.3.1 The Value of X-ray Studies ...... 11 1.4 Outline of the Thesis ...... 12 1.4.1 Chapter 2: Exploratory ASCA Observations of BAL QSOs . 12 1.4.2 Chapter 3: Heavy Absorption in Soft X-ray Weak AGN . . . 13 1.4.3 Chapter 4: Markarian 231: A Prototypical Low-Ionization BAL QSO? ...... 13 1.4.4 Chapter 5: X-ray Spectroscopy of QSOs with Broad Ultravi- olet Absorption ...... 14 1.4.5 Chapter 6: First Results from an Exploratory Chandra Survey of Large Bright Quasar Survey BAL QSOs ...... 14 1.4.6 Conclusions ...... 14

Chapter 2. Exploratory ASCA Observations of Broad Absorption Line Quasi-Stellar Objects ...... 16 2.1 Introduction ...... 16 2.2 Observations and Data Analysis ...... 19 2.2.1 Observation Details and Data Reduction ...... 19 2.2.2 Image Analyses and Count Rate Constraints ...... 19 2.2.3 Column Density Constraints ...... 20 2.3 Notes on Individual Observations ...... 25 2.4 Conclusions and Future Prospects ...... 27 v

Chapter 3. Heavy X-ray Absorption in Soft X-ray Weak Active Galactic Nuclei 29 3.1 Introduction ...... 29 3.2 Observations and Data Analysis ...... 32 3.2.1 PG 1011–040 ...... 34 3.2.2 PG 1535+547 (Mrk 486) ...... 37 3.2.3 PG 2112+059 ...... 40 3.3 Discussion and Conclusions ...... 43 3.3.1 PG 1011–040 ...... 43 3.3.2 PG 1535+547 (Mrk 486) ...... 43 3.3.3 PG 2112+059 ...... 44 3.3.4 X-ray Absorption as the Primary Cause of Soft X-ray Weakness 46 3.3.5 X-ray and UV Variability ...... 47 3.3.6 Soft X-ray Weak AGN and the X-ray Background ...... 47

Chapter 4. Markarian 231, a Prototypical Low-Ionization BAL QSO? ...... 50 4.1 Introduction ...... 50 4.2 Observations and Data Reduction ...... 52 4.3 Image Analysis ...... 53 4.3.1 Nucleus and Extended Emission ...... 53 4.3.2 Non-Nuclear Point Sources ...... 54 4.4 Variability Analysis ...... 55 4.5 Spectral Analysis ...... 56 4.5.1 Nucleus ...... 56 4.5.1.1 Spectrum Extraction ...... 56 4.5.1.2 Basic Spectral Modeling ...... 58 4.5.1.3 Reflection-Dominated Spectrum and Iron Line . . . 59 4.5.1.4 Alternate Model ...... 61 4.5.1.5 A Jet Contribution to the X-ray Flux? ...... 61 4.5.1.6 Spectral Checks ...... 62 4.5.2 Extended Galaxy Emission ...... 63 4.6 Discussion ...... 76 4.6.1 Mrk 231 as a Broad Absorption Line QSO ...... 76 4.6.1.1 Spectroscopic Classification ...... 76 4.6.1.2 X-ray Variability and Spectral Results ...... 77 4.6.2 Mrk 231 as a Starburst Galaxy ...... 80 4.7 Future Work ...... 80

Chapter 5. X-ray Spectroscopy of QSOs with Broad Ultraviolet Absorption Lines 83 5.1 Introduction ...... 83 5.2 Sample Data ...... 84 5.3 Results ...... 86 5.3.1 Normal Underlying QSO Continua ...... 86 5.3.2 Complex Absorption ...... 86 5.3.2.1 H 1413+117 ...... 87 5.4 Discussion ...... 88 vi

5.5 Future Work ...... 88

Chapter 6. First Results from an Exploratory Chandra Survey of Large Bright Quasar Survey Broad Absorption Line QSOs ...... 92 6.1 Introduction ...... 92 6.2 Observations and X-ray Data Analysis ...... 93 6.3 Testing for Significant Correlations ...... 96 6.4 Discussion ...... 100 6.5 Future Work ...... 101

Chapter 7. Conclusions ...... 106 7.1 Plans for the Future ...... 106

References ...... 108 vii List of Tables

2.1 BAL QSO General Information and Observing Log ...... 18 a 3 1 2.2 ASCA Detector Count Rates (10− photons s− ) ...... 21 2.3 Limits on Intrinsic Absorptiona ...... 22

3.1 General Information and Observing Log ...... 31 3.2 Observed ASCA Parameters ...... 35 3.3 Parameters of Model Fits to the ASCA Data ...... 38

4.1 Counts, Fluxes, and Luminosities ...... 57 4.2 Parameters for Model Fits to the Chandra Data for Mrk 231a ...... 74 4.2 Parameters for Model Fits to the Chandra Data for Mrk 231a ...... 75

5.1 Basic Properties of QSOs with Broad Ultraviolet Absorption Lines . . . 85

6.1 Basic Properties and Observation Log ...... 94 6.2 X-ray Properties ...... 97 6.3 Results from Bivariate Statistical Testsa ...... 98 viii List of Figures

1.1 Ultraviolet spectrum of PHL 5200, the first BAL QSO discovered . . . . 15 1.2 X-ray absorption by neutral gas of a QSO continuum ...... 15

2.1 Column density lower limits for PG 0043+039: αox ...... 24 2.2 Column density lower limits for PG 0043+039: αix ...... 25 3.1 Hubble Space Telescope ultraviolet spectrum of PG 2112+059 ...... 33 3.2 X-ray images of soft X-ray weak AGN ...... 34 3.3 ASCA spectra of PG 1011–040 ...... 36 3.4 ASCA spectra of PG 1535+547 ...... 39 3.5 Confidence contours: PG 1535+547 ...... 40 3.6 ASCA spectra of PG 2112+059 ...... 41 3.7 Confidence contours: PG 2112+059 ...... 42 3.8 Spectral energy distributions ...... 48 3.9 Rest-frame C iv absorption EW versus αox ...... 49 4.1 Chandra ACIS-S3 images of Mrk 231 ...... 64 4.2 Surface plots of X-ray images of Mrk 231 ...... 65 4.3 Radial surface-brightness profiles of the X-ray emission from Mrk 231 . 66 4.4 X-ray contours overlaid on optical images of Mrk 231 ...... 67 4.5 Hardness ratio map of Mrk 231 ...... 68 4.6 Image of putative off-nuclear X-ray source ...... 69 4.7 Light curves of X-ray emission from Mrk 231 ...... 70 4.8 Kolmogorov-Smirnov probability of variability as a function of energy . 71 4.9 ACIS-S3 spectra and models of the nuclear emission ...... 72 4.10 ACIS-S3 spectrum of the extended emission from Mrk 231 ...... 73 4.11 Rest-frame Hubble Space Telescope ultraviolet spectrum of Mrk 231 . . . 76 4.12 Diagram of a possible geometry for the nucleus of Mrk 231 ...... 81

5.1 Chandra ACIS-S3 spectra of QSOs with broad ultraviolet absorption . . 90 5.2 X-ray continua and spectral energy distributions of typical QSOs . . . . 91

6.1 Optical-X-ray Offsets ...... 95 6.2 Absorption line parameters versus αox ...... 103 6.3 Continuum properties versus αox ...... 104 6.4 Hardness ratio plots ...... 105 ix Acknowledgments

I am particularly grateful for four years of financial support from the Pennsylva- nia Space Grant Consortium as well as two years of funding from the Zaccheus Daniel Foundation for Astronomical Science. Both of these organizations allowed me to pursue valuable experiences related to interests in astronomy not directly associated with my thesis research. NASA’s Graduate Student Research Program was essential in enabling me to complete my program in five years and to present it at several meetings. Other grants contributed to specific projects, and I have thanked them in the appropriate chapters. Discovery science does not happen without such agencies.

The completion of this thesis would not have been possible without the support of many people who have contributed in very different ways to this journey through graduate school. I list them here in no particular order. I thank my committee as a whole; they have been reasonable and challenging in appropriate measures. I am grateful to my advisor Niel Brandt for the time and effort he has put towards my development as a professional astronomer. His high expectations have contributed significantly to my success to date. Jane Charlton encouraged my independence by providing me with the opportunity to work with the beautiful Hubble Space Telescope images of Stephan’s Quintet, and I value my continued connection to that field through her other ongoing projects. Mike Eracleous has been an unofficial advisor both on and off the soccer field, and I appreciate my good fortune in observing with a master. I was so fortunate to arrive at Penn State in time to work with both the Chandra X-ray Observatory and the Hobby-Eberly Telescope. There are not many institutions that give graduate students direct access to such data. I certainly appreciate first-hand the growing pains of new facilities as well as the excitement of first light. It always takes longer than anticipated, always. As I am sure Niel would agree, the same has been true of each of my papers. The Penn State X-ray group has provided a home port for me, and I thank Gordon Garmire for bringing part of the ACIS instrument team to Penn State and allowing me to participate. It is no exaggeration that Penn State is probably the best place around to be a student of X-ray astronomy — the fabulous data, expertise, and community are unparalleled. Leisa Townsley and Pat Broos, in addition to their rigor and skills in managing ACIS data of all sorts, are generous and loyal friends. I especially appreciate their sanity. The third desk in 405 Davey probably bears a dent from the many hours I have spent perching there. I am also grateful to George Chartas for including me as a collaborator on many projects. I owe all of my Italian vocabulary to Vlad Getman, who has made each day that I see him more cheerful. Joe Hill is an excellent example of a great woman with a head for hardware. It is inspiring to know someone who thinks CCDs are so cool. Franz Bauer and Gordon x

Richards generously contributed general advice on applying for jobs and discussions of interesting science. They are all fun to play with on the soccer field. The wisdom of several senior Penn State graduate students was invaluable. Laura Cawley made sure that I had all the information in hand to “do graduate school the right way”. Any strategic mistakes are therefore my own and cannot be attributed to a lack of information. Dave Andersen and Anna Jangren were excellent office mates and remain good friends, and I enjoyed our time together in 406 Davey. Mike Sipior has always been available to share computer advice, sardonic commentary, and excellent beer. I will forever feel a connection to Rajib Ganguly and Mike Weinstein, fellow travelers through classes and other horrors. Ann Hornschemeier has been a homework buddy, roommate, soccer friend, office mate, and general ally during my entire time in State College. I cannot imagine surviving comps and the other major trials of graduate school without her. I am indebted to my parents for developing my curiosity and for providing me with access to the means to pursue it. They never questioned the utility of an astrophysicist in the family, for which I am extremely grateful. I am also glad that my sister has recovered from the shock of being so closely related to a science geek. The level of encouragement and inspiration I receive from my husband Chris Smeenk can only be described as structural - absolutely essential, though perhaps not obvious from the exterior. This academic adventure has been a joint undertaking. xi

When I look down, I just miss all the good stuff; when I look up, I just trip over things. –Ani DiFranco 1

Chapter 1

Introduction

1.1 Active Galactic Nuclei

At the centers of some galaxies, there is a light source that outshines all of the stars in the galaxy. These so-called “active galactic nuclei” (hereafter AGN) exhibit broad- band spectra distinct from those of normal galaxies, whose emissions are the sum of their composite stars. The most luminous of these AGN (with M 23), quasi-stellar V ≤ − objects (QSOs or quasars), were first discovered as point-like sources at high Galactic latitude with unusually blue colors. QSOs are relatively rare, and the first found were at much greater distances than more common AGN. At such great distances, the light from their galaxies was completely overwhelmed by the luminous central sources.

1.1.1 The Standard Model AGN can generate the power of up to 1000 times an entire galaxy within a volume comparable to the size of our solar system. According to the standard picture, the accretion of material onto a supermassive (M = 106–109M ) black hole generates the incredible luminosity of these objects. As gaseous material spirals into the black hole, its gravitational energy is converted into optical and ultraviolet radiation via a dissipative accretion disk that heats the optically thick material to high temperatures (see chapters 3 and 4 of Peterson 1997 for more details). In the local Universe, normal galaxies contain supermassive black holes in their nuclei, though only 1 in 100 local galaxies exhibits luminous AGN activity (e.g., Peterson 1997, p. 4). In addition, the mass of each central black hole is correlated with the mass of the bulge of its host galaxy (e.g., Kormendy & Gebhardt 2001; Merritt & Ferrarese 2001). This connection provides strong evidence that the growth of supermassive black holes and their host galaxies are intimately connected. Investigating AGN, the manifestations of the accretion phase of supermassive black holes, is thus fundamental to understanding the evolution of galaxies in general.

1.1.2 Progress Towards Unification While the scenario outlined above provides a simple picture for the power source underlying AGN, the diversity of observed phenomena does not immediately appear to support this clean view. Some of the variety can be explained by an accident of view- ing angle, while other observations may require fundamental differences in the physical characteristics of different types of AGN. Khachikian & Weedman (1971) first defined two spectroscopically distinct types of active galaxies. The spectra of the first class, called Type 1, contained broad emission 2

3 1 lines (with a full width at half-maximum [FWHM] of v & 10 km s− ) of common permit- ted transitions, such as the Balmer lines, in addition to narrow emission lines (FWHM 2 1 of v 10 km s− ) of permitted and forbidden transitions such as [O iii] λλ4959, 5007. ∼ In contrast, Type 2 objects exhibited only the narrow emission lines. In unifying these two classes, one breakthrough observation was the detection of broad emission lines in the spectrum of polarized light from the Type 2 object NGC 1068 (Antonucci & Miller 1985). This indicated that this AGN was in fact a Type 1 object in which the region generating the broad emission lines was obscured. Only the broad emission line flux scattered into the line of sight could be seen. As the scattered flux is of much lower intensity than the direct emission from the rest of the nucleus, the scattered component could only be seen in polarized light. The obscuring medium invoked to explain the obstruction to the broad line region in Type 2 objects is a torus of molecular material at distances of several parsecs or more. This torus could be associated with other structures within the galaxy, for example a circumnuclear ring of star formation. Though low luminosity AGN are four times more likely to be Type 2 than Type 1, the Type 2 designation is rarer among the more luminous QSOs. This observational fact may result from selection biases against finding Type 2 QSOs; alternatively, more luminous AGN may not allow tori to survive for long (e.g., Antonucci 1993).

1.1.3 Radio-Loud/Radio-Quiet AGN AGN spectral energy distributions are distinguished from stars and galaxies by the presence of substantial power across many decades of frequency, from the infrared through the X-ray. A subset of “radio-loud” AGN also exhibit significantly more power in the radio regime relative to their optical and ultraviolet continua. Within this class of radio-loud AGN, there are Type 1 (or blazar, radio-loud quasar, and broad line radio galaxy) and Type 2 (or narrow line radio galaxy) subclasses as well. These have been successfully unified via orientation in a similar manner to the radio-quiet AGN, with the inclusion of an additional feature, the relativistic jet (see Urry & Padovani 1995 for a thorough review of this topic). These jets are often galaxy-scale or larger structures that contribute significantly to the bolometric power of an AGN, and they are the source of the radio emission which defines the radio-loud class. Formally, radio-loud AGN are defined ? ? as those with R & 10, where R is the ratio of flux densities at 5 GHz and 2500 A.˚ Radio-loud AGN compose approximately 10% of active galaxies. Efforts to elucidate the fundamental parameters driving the observed differences between radio-quiet and radio- loud AGN have not been successful as of yet. The distinction between the two classes may be related to environmental and/or evolutionary influences (e.g., Wills 1999). For example, Wilson & Colbert (1995) suggest that radio-loud QSOs may have maximally spinning black holes that can drive the formation of energetic jets. Other authors (e.g., Laor 2000) have found correlations between black hole mass and radio-loudness leading to claims that mass is the deciding factor, though these results are not universally accepted (Ho 2002). 3

1.2 Broad Absorption Line QSOs

As a result of evolution or the accident of orientation, specific populations of AGN offer privileged access to physical processes related to black hole growth. In particular, broad absorption line (BAL) QSOs enable study of the energetic outflows accompanying active black hole accretion. Approximately 10% of all radio-quiet QSOs exhibit broad ultraviolet absorption lines blueshifted from their corresponding emission lines. The P Cygni-type profiles of these lines indicate that our line of sight is probing energetic winds being expelled from the central engine with velocities up to 10–30% of the speed of light. As argued in the following sections, this outflow is believed to be a common element of the QSO environment. As such, it provides a potential mechanism for angular momentum loss from the accretion disk (e.g., de Kool 1997). Furthermore, the expulsion at high velocity of metal-rich gas may have important effects on star formation in the host galaxy as well as the evolution of accretion onto the black hole (Fabian 1999). For example, high velocity gas could conceivably shock the interstellar medium thus triggering star formation. Alternatively, the ejection of gas from the nucleus can halt star formation and black hole accretion by removing the fuel for these processes. On larger scales, these winds can also pollute the inter-galactic medium with metal-rich gas from galaxy bulges. Figure 1.1 illustrates the broad absorption line phenomenon with a sample spec- trum. As in this example, the broad absorption lines typically seen are those of the high- ionization resonance transitions of O vi λλ1032, 1038, N v λλ1239, 1243, Si iv λλ1394, 1403, and C iv λλ1548, 1551. However, observations with sufficient spectral coverage have de- tected more than two dozen broad absorption lines from almost 20 different ions (Arav 2001). The feature typically used to identify an AGN as a BAL QSO is broad C iv absorption. This ultraviolet diagnostic does not shift into the optical range readily ob- served from the ground (longward of 3700 A)˚ until z 1.4. As a result, the existence ∼ & of BAL QSOs at lower redshifts was not well documented before observations with the space-based International Ultraviolet Explorer (e.g., Turnshek & Grillmair 1986) identi- fied several.1 Turnshek & Grillmair (1986) also remarked upon the apparent deficit of BALs among lower luminosity AGN.

1.2.1 Continuum and Spectral Energy Distribution There are two extreme interpretations to the observation of a 10% BAL QSO frac- tion in color-selected QSO surveys. At one end, the BAL outflow is a common feature of all QSO nuclei that only covers 10% of the sky as seen from the continuum source. Al- ∼ ternatively, approximately one in 10 QSOs has a BAL wind with a covering fraction near unity. To address these possibilities, Weymann et al. (1991) systematically compared the rest-frame ultraviolet continuum and emission-line properties of BAL QSOs and a comparison sample of typical QSOs. They found that these features were “remarkably similar” between the two classes. In contrast, BAL outflows with very high covering fractions would be expected to result in differences in emission-line properties due to the

1The low signal-to-noise of some of these spectra did lead to some false identifications, e.g, of PG 1416 1291 (Green et al. 1997). − 4 additional contribution of scattered flux from the BAL region (e.g., Hamann et al. 1993). This result led to the conclusion that the BAL wind is likely to be a common feature of the environment of all QSOs. In comparing the spectral energy distributions of BAL and non-BAL QSOs of all types, Stocke et al. (1992) found a marked anti-correlation between radio-loudness and the presence of BALs. They approached the problem from both ends by search- ing for BALs among known radio-loud QSOs and by testing for radio-loudness among identified BAL QSOs. This apparently fundamental difference between radio-loud and radio-quiet QSO populations was taken to have profound implications for understanding the differences among the radio properties of luminous AGN. However, this apparently strong result has been challenged by the findings of the Faint Images of the Radio Sky at Twenty centimeters (FIRST) Bright Quasar Survey (FBQS; White 2000), which has dis- covered 29 radio-selected BAL QSOs. As discussed further in 1.2.6, several of these are § formally radio-loud (Becker et al. 1997; Becker et al. 2000). Subsequently, a previously well-known low redshift radio-loud QSO, PKS 1004+13, was found to have BALs (Wills et al. 1999). In fact, once corrected for incompleteness, the fraction of BAL QSOs among radio-detected Sloan Digital Sky Survey QSOs, 9%, is consistent with the fraction for ∼ those selected optically (Menou 2001). At this point, the exact nature of the relationship between radio and BAL properties is still unresolved.

1.2.2 X-ray Properties (circa 1997) From observational evidence and accretion physics, X-ray emission is understood to be a universal property of QSOs. Based on the rapid variability of the soft X-rays from AGN in conjunction with the standard black-hole paradigm, these photons are believed to be emitted on small scales, on the order of light minutes. This compact spatial extent is most naturally explained as encompassing the region immediately surrounding the central black hole. Energetically X-rays are significant, comprising between 2–20% of the radiant power of quasars. Though the exact theoretical underpinnings of this physical picture remain to be worked out in full, the primary AGN X-ray emission at energies up to 100 keV is believed to originate as ultraviolet and optical photons from the ∼ 7 9 accretion disk that illuminate a hot ( 10 − K) corona. The corona, or ionized plasma ∼ surrounding the black hole, reprocesses this radiation, perhaps via inverse Compton scattering. The result of this process is a non-thermal X-ray spectrum (e.g., Krolik 1999, pp. 202–216). This thesis will be primarily concerned with X-ray emission in the 0.5–10.0 keV band, a dynamic range of frequency as large as that from the near infrared to the ultraviolet. For luminous, radio-quiet QSOs, their emission in this regime can Γ usually be described empirically as a power law of the form FE = NEE− where NE is 2 1 1 the normalization at 1 keV in units of photons cm− s− keV− . The parameter Γ is called the photon index, and it is related to the spectral index, α: Γ = 1 + α. Though there is significant scatter in its value, Γ typically ranges from 1.7–2.3 (e.g., Reeves & Turner 2000; George et al. 2000). With the R¨ontgen Satellite (ROSAT ) X-ray observatory, BAL QSOs were first shown to be underluminous in X-rays relative to their ultraviolet flux (Kopko et al. 1994; Green et al. 1995). This observation could result from either of two situations: (1) 5

BAL QSOs are anomalous and do not generate X-rays in the manner of other radio-quiet AGN, or (2) BAL QSOs have intrinsic X-ray properties consistent with normal radio- quiet QSOs, but they are cloaked by absorbing gas which prevents soft X-rays from reaching the observer. Assuming the latter case is true, then the lack of X-ray detections can be used to constrain the amount of absorbing gas along the line of sight. Green & Mathur (1996) performed this analysis with a heterogeneous sample of 12 BAL QSOs with pointed ROSAT Position Sensitive Proportional Counter (PSPC) observations. Of this sample, only two were detected; one of these was subsequently shown not to contain BALs. The lack of detections implied intrinsic absorbing column densities of 22 2 N 2 10 cm− of neutral gas. In this case, N represents the number of hydrogen H& × H atoms in a gas with solar abundances along the line of sight through a tube with a cross-sectional area of 1 cm2. Though the hydrogen contributes only negligibly to the X- ray opacity, this parameter is commonly used to characterize the amount of intervening matter to the object in question. The lower limits from the ROSAT observations were at least an order of magnitude larger than contemporary measurements of NH from the ultraviolet absorption lines.

1.2.3 Broad Absorption Lines There is significant diversity in the nature of BALs. Some show strong structure suggesting distinct systems at several velocities, while others are extremely smooth. The velocity of BAL onset relative to the emission line peak varies significantly as does the velocity spread and terminal velocity of the outflow. The last can reach extremes 4 1 of 5 10 km s− in certain objects. The measured absorption equivalent widths ∼ × (EWa) range from 10–100 A,˚ but these values are not particularly illuminating given the complexity of the lines. In addition, with the difficulty of finding the continuum and accounting for the emission-line fluxes as well as the contribution from scattering, calculating EWa is not always straightforward. Weymann et al. (1991) defined two standardized parameters for characterizing the absorption lines. The first, the “BALnicity index” (BI), is a conservative measure of EWa motivated by the desire to exclude contributions from the emission lines as well as associated absorption.2 Specifically, the BI is measured by evaluating the following integral for the C iv absorption line:

3000 1 BI = [1 f(v)/0.9]Cdv (km s− ) (1.1) − − Z25,000 where f(v) is the normalized flux as a function of velocity displacement from the line center (Weymann et al. 1991, Appendix A). The value of C is set to zero until the quantity 1 in brackets has been continuously positive for 2000 km s− ; in this case, C equals 1. ≥ It is reset to zero if the value in square brackets becomes negative. While this appears to be an awkward working definition, the intention is to quantify the broad absorption,

2 1 Associated absorption is commonly defined as absorption within 3000 km s− of the emission-line redshift. It may result from gas clouds within the host galaxyrather than from the immediate vicinity of the black hole. 6 i.e., the continuous absorption — a dip of more than 10% below the continuum — with 1 widths exceeding 2000 km s− . These cautions should eliminate misclassifications due to intervening absorbers or a poorly defined continuum. The upper and lower limits are set to avoid confusion with the Si iv emission line and C iv absorption potentially associated with the host galaxy, respectively. By this definition, the maximum value for the BI is 1 20,000 km s− , and any QSO with BI = 0 is not considered a BAL QSO. Empirically, QSOs are observed to have BALs with onset velocity of absorption 1 3 1 ranging from 0 km s− to several 10 km s− . To quantify this property, the second ∼ standardized absorption-line measure defined is the “detachment index” (DI). This pa- rameter was considered useful for examining if detached absorption, that is, absorption with onset velocity significantly blueshifted from the C iv emission line, was related to other spectral properties. The procedure for calculating the DI is rather complex, but a brief description follows. As the onset of absorption is generally rather abrupt, measur- ing this value is straightforward. The goal is then to normalize the velocity of onset for each trough (measured from the systemic redshift as defined by the Mg ii emission line) with the width of the C iv emission line. This can be difficult as the blue wing of C iv (and sometimes the red wing as well) is often heavily absorbed. Further details are in Appendix B of Weymann et al. (1991).

1.2.4 Polarization In general, the distribution of continuum polarization fraction for BAL QSOs peaks at a higher value than that of normal QSOs (Schmidt & Hines 1999). Several pos- sible reasons have been put forth for this empirical result, and the most widely accepted is that there is a larger fraction of scattered flux in the typical ultraviolet BAL QSO spec- trum. This could result either from a specific geometry that would favor scattering into the line of sight (e.g., Goodrich & Miller 1995) or from gray attenuation of the primary continuum (Goodrich 1997). In the latter case, the reduced flux would cause BAL QSOs to be underrepresented in optically selected samples (Goodrich 1997; Schmidt & Hines 1999). More detailed spectropolarimetry has offered greater insight into the nature of the polarization. In the spectropolarimetry atlas of Ogle et al. (1999), detailed analysis revealed several general properties. The continuum polarization typically increases to the blue, consistent with dust scattering. Most significantly, the polarization increases in the BALs, indicating an excess of scattered relative to direct light. Furthermore, in most cases the emission lines and continua have different polarization position angles, and the emission lines are less polarized. Overall, Ogle et al. (1999) found the polarization results to be consistent with a unified picture for BAL and non-BAL QSOs whereby all QSOs have BAL regions.

1.2.5 Variability The largest campaign to monitor the absorption-line variability of BAL QSOs was undertaken by Barlow (1993). Over the three-year course of this investigation, approximately one-third of the 54 objects in that sample showed continuum variations of 0.1–0.3 mag on timescales of 0.25–1 year in the QSO rest frame. Of the smaller sample 7 of 23 with spectrophotometric monitoring, 65% showed variations of the normalized ∼ intensity in the absorption lines at some level. Notably, there was no evidence for changes in the velocity structure of any of the BALs. This latter result provides a significant theoretical constraint on the nature of the absorbing gas. The shape of the profile implies an accelerating flow, and the lack of velocity shifts is more consistent with a smooth wind rather than individual clouds. Though often changes in the normalized flux in the absorption line occurred in tandem with continuum variations, this was not always the case. According to Barlow (1993), this lack of one-to-one correspondence implies that the relationship between the continuum near the BALs and the ionizing continuum in the extreme ultraviolet regime is not simple. In addition, the BAL gas may be exposed to radiation that the observer does not see.

1.2.6 Observational Subclasses BAL QSOs exhibit a diversity of empirical properties that have led to the def- inition of several subclasses. In general, these groups are not mutually exclusive, and there is some debate about specific identifications. The more common ones are briefly described in this section. Low-Ionization BAL QSOs: The most well-documented subclass, the low-ionization BAL QSOs (hereafter LoBAL QSOs), exhibits BALs in the transitions of Mg ii λλ2796, 2803 and Al iii λλ1854, 1863 in addition to the more standard high-ionization lines. This group, approximately 15% of BAL QSOs in optically selected samples, was the notable exception to the finding that BAL and non-BAL QSOs have largely similar continuum and emission-line properties (Weymann et al. 1991). LoBAL QSOs tend to have red- dened ultraviolet and optical continua, indicative of dust extinction, as well as more pronounced optical Fe ii emission. Correcting these spectra for extinction like that seen in the Small Magellanic Clouds recovers a typical QSO continuum (Sprayberry & Foltz 1992). As a result of their red colors, LoBAL QSOs tend to be missed by standard optical selection techniques; however, they appear to be much more common in infrared-selected samples (e.g., the Infrared Astronomical Satellite AGN surveys of Low et al. 1989 and Boroson & Meyers 1992). For Mg ii gas to exist along the same line of sight as more highly ionized species, there are two reasonable possibilities in a monotonically accelerating gas: (1) the gas recombines along the outflow trajectory or (2) lower-ionization gas becomes more highly ionized as it is accelerated. In the first case, the Mg ii absorption will tend to be narrower and more blueshifted than the C iv absorption. In the latter case, the Mg ii will be at lower velocities. From observations of a small sample of LoBAL QSOs, Voit, Weymann & Korista (1993) found that the low-ionization line troughs exhibit similar structure and inhabit the low end of C iv absorption velocity parameter space. This suggests that the Mg ii absorbing gas is closer to the source of the radiation pressure, or that the gas is not accelerating monotonically. These authors favored the second option, as it is difficult to maintain low-ionization gas in the extreme radiation environment of an accreting supermassive black hole. However, other than for simplicity, there is no reason to assume that the outflow is spherically symmetric, in which case the reasoning 8 above does not necessarily hold. The issue of the geometry of the wind has been further addressed with dynamical modeling (see 1.2.7), but the relationship between high and § low-ionization BAL QSOs is still poorly understood. Radio-Loud BAL QSOs: As introduced in 1.2.1, radio-loud BAL QSOs have been § found in reasonable numbers by the FBQS, which is significantly more sensitive than previous radio surveys for QSOs. As a result, it has discovered many more QSOs with intermediate radio-to-optical flux ratios than were known in the past (White 2000). Cross-correlating Sloan Digital Sky Survey QSOs with FIRST sources has also iden- tified radio-loud BAL QSOs. In both surveys, the radio-loud BAL QSOs tend to lie in the intermediate regime of R?3 ranging from 10–500. In terms of radio luminos- 31.4 34.6 1 1 ity, LR, these objects span from 10 –10 erg s− Hz− . Those BAL QSOs with 32 1 1 LR & 10 erg s− Hz− would formally be considered radio luminous (Becker et al. 2000; Menou 2001). Currently, there are no known BALs found among the most radio-loud ? 3 QSOs with R & 10 (Becker et al. 2000; Menou 2001). Though the radio-loud subclass bears further study in its own right, there are no radio-loud BAL QSOs included in this thesis. “PHL 5200-like” BAL QSOs: With the discovery by Lynds (1967) of broad absorp- tion lines in the spectrum of PHL 5200, the study of BAL QSOs was born. However, with the prevalent irony of astronomical taxonomy, the “proto-typical” BAL QSO PHL 5200 is an atypical example of its class. In fact, Turnshek et al. (1988) used “PHL 5200-like” to describe a subset of BAL QSOs (currently with a membership of three) with narrow C iv emission lines and smooth C iv BALs indicating maximum outflow velocities of several 4 1 10 km s− . The C iv emission lines are unlikely to be inherently narrow, but the onset of absorption at low velocities makes them appear so. Two objects within this group, PHL 5200 and RS 23, also exhibit high continuum polarization of 5% (Hutsem´ekers & ∼ Lamy 2000; Lamy & Hutsem´ekers 2000). The third PHL 5200-like QSO, H 1413+117, has been identified as a LoBAL QSO (Hazard et al. 1984; Gallagher et al. 2002a), and therefore perhaps has less in common with the others. Though the utility of this distinct classification is debatable, it has been included here for completeness. Mini-BAL QSOs: After Weymann et al. (1991) standardized the definition of a BAL QSO with the introduction of the BALnicity index, higher quality data ousted several AGN that had been previously classified as BAL QSOs (e.g., PG 1411+442; Corbin & Boroson 1996). Typically, these sub-par objects have P Cygni-type broad 2 3 1 absorption lines with velocity widths of 10 − km s− and lower terminal velocities. ∼ Thus, they do not formally meet the BAL QSO criterion of Weymann et al. (1991), but they nonetheless show evidence for an energetic, outflowing wind along the line of sight. The main distinction between these “mini-BAL QSOs” and their bona-fide cousins is apparently their lower intrinsic luminosity. This group of AGN may provide continuity between the extreme winds seen in BAL QSOs and the associated absorbers in lower luminosity AGN, but this connection has not been thoroughly investigated. This classi- fication is largely informal, and as of this writing there exists no comprehensive reference for this group.

3R? = f /f . 5 GHz 2500 A˚ 9

1.2.7 Theoretical Models for BAL QSOs Theoretical models of the geometry and dynamics of an absorption-line outflow must incorporate the following observational constraints:

4 1 Material accelerated up to a few 10 km s− . • Coverage fractions of approximately 10–20% based on limits from emission-line • and polarization studies (e.g., Hamann et al. 1993; Ogle et al. 1999).

X-ray weakness (Kopko et al. 1994; Green & Mathur 1996; Gallagher et al. 1999; • Brandt et al. 2000).

Lack of variability in the velocity structure of the BALs on rest-frame timescales • of . 3/(1 + z) years (Barlow 1993). Location of the wind outside of or cospatial with the broad-line region based on • N v absorption of the Lyα emission line.

Observed ionization states of the absorbing gas, which appear to be roughly con- • stant over velocity space (e.g., Turnshek 1988).

The results from polarization studies provide important constraints on the geom- etry of the outflow region. The polarized continuum and absorption lines in the BAL wind indicate that the outflow is not spherically symmetric; the contribution from scat- tering implies asymmetry (e.g., Hines & Wills 1995). Furthermore, the observation of polarized flux in the BALs indicates that the scattered flux is not passing through the BAL outflow, i.e., the BAL wind does not fully cover the continuum source. The global coverage fraction is bounded by the lack of a significant contribution of scattered flux to the emission lines to be . 20% (Hamann et al. 1993). Assuming all QSOs have BAL regions sets the coverage fraction to 10%, the fraction of QSOs with BALs in optically ∼ selected samples. If the 10–20% covering fraction is not due to random sky coverage of the BAL wind, the two natural geometries for the outflow are polar or equatorial. The advantage of the equatorial outflow is that the accretion disk provides a natural source for the accelerated material. The problem of accelerating the BAL gas to the high velocities observed has led to several proposed solutions. The possible forces acting on the outflowing material in- clude gas pressure, magnetic pressure, and radiation pressure. Several pieces of evidence point to the last as the prime mover. First, the photon momentum from the luminosity absorbed in the BALs is comparable to simple estimates of the momentum in the out- flow (e.g., de Kool 1997). Perhaps most striking is the observation of line-locking, the separation of absorption systems by the velocity difference in the N v/Lyα lines (e.g., Korista et al. 1993) and the doublet spacing of C iv (Foltz et al. 1987). These observa- tions are difficult to explain if line pressure is not of paramount importance for driving the outflow. Earlier work on the outflow assumed that the gas is clumped in small clouds that require a confining medium such as magnetic pressure (de Kool & Begelman 1995) or gas pressure from a hot, diffuse medium. Taking a different tack, Murray et al. (1995) 10 developed a dynamical approach using a smooth accretion-disk wind as a source of both the broad emission and absorption lines. According to their picture, at radii starting at 1016 cm, radiation and gas pressure from the accretion disk lift material above the ∼ disk photosphere where it is struck by the force of ultraviolet through X-ray photons from the inner accretion disk and corona. The wind, initially co-rotating with the disk, is accelerated radially by the momentum of the line photons, and the material launched from the innermost radii reaches the highest terminal velocities. In order to launch the wind at such small distances from the high-energy continuum, Murray et al. (1995) introduced a rather ad hoc structure, the “hitchhiking gas,” to shield the outflow from becoming completely ionized by the soft X-ray continuum emission. The hitchhiking 23 24 2 gas has a column density of 10 − cm− , sufficient for absorbing soft X-ray emission. Primary objections to this model included the necessary fine tuning of the physical state of the hitchhiking gas and the small distances for both the broad emission line and broad absorption line material (e.g., Blandford 1995). Extending this model using hydrodynamical calculations, Proga et al. (2000) nat- urally produced a thick, highly ionized plasma that shields the fast outflow from soft X-rays. Their wind is also launched at similar radii to that of Murray et al. (1995), and the case for line-driven winds appears to be growing stronger. One potential difficulty re- mains: the observed lack of absorption-line velocity change. For winds launched at such small radii, instabilities are expected to result in variations on timescales of a few years. Though this has not been observed, the one major variability study by Barlow (1993) did not span enough time to rule out this model. Studies with baselines of a decade are more are required, particularly as most BAL QSOs are found at higher redshifts where the time dilation factors are non-negligible.

1.3 Motivation for this Project

Some fundamental questions remain regarding the nature of the energetic winds from luminous QSOs. From a physical point of view, one of the outstanding questions is the rate of mass ejection from the central engine. This parameter is of primary impor- tance for understanding the balance of accretion and outflow in luminous AGN. From an observational point of view, the relationship between absorption in the ultraviolet and X-ray remains ambiguous though one does not tend to appear without the other. The current difficulties with characterizing the absorber properties in both regimes ob- struct efforts at absorber unification. While X-ray studies offer better constraints on the column density than those in the ultraviolet, the dynamical information available from X-rays is limited. Constraining the specific nature of the X-ray absorber is beyond the signal-to-noise ratio and spectral resolution of existing X-ray data. Though ultraviolet spectroscopy allows high-resolution velocity analysis of the outflow, ultraviolet absorption lines typically exhibit saturation and partial covering. These effects make determining the specific physical parameters of the gas, e.g., column density, ionization state, and abundances, extremely difficult. In fact, analyses of ab- sorption lines over the last decade suggested that the abundances of the outflowing wind were 10–100 times the solar abundances (e.g., Hamann 1997). This conclusion was later found to result from non-black saturation: scattered flux often “fills in” absorption line 11 troughs (e.g., Goodrich & Miller 1995). Furthermore, at a given velocity the continuum source can be partially covered by saturated absorption (e.g., Arav et al. 1999). These phenomena make standard curve-of-growth analysis essentially useless for measuring col- umn densities. Thus, ultraviolet absorption lines will not easily enable measurements of the mass outflow rate in luminous AGN.

1.3.1 The Value of X-ray Studies As X-ray photons travel from the vicinity of the black hole, matter they traverse imprints the X-ray spectrum with signatures of its physical state. X-ray spectra thus have the potential to offer excellent probes of the immediate environment of the black hole via absorption studies. Ideally, X-ray observations would answer the following questions: (1) Are the underlying X-ray continua and spectral energy distributions of BAL QSOs the same as normal QSOs? (2) What is the physical state (such as the column density, ionization parameter, covering fraction, and velocity structure) of the absorbing gas? High sensitivity coupled with broad bandpasses serve to address the first question in this thesis, while the second will require significantly longer observations at the higher spectral resolutions of transmission or reflection gratings to be completely resolved. In the soft X-ray regime, absorption is dominated by bound-free transitions at the ionization energies of the innermost electrons of abundant metals. X-rays are not as sensitive to the exact nature of the gas as ultraviolet photons; gas in all states, from dust to molecules to partially ionized plasma, will leave a mark on the X-ray spectrum. For neutral gas, the power-law continuum will have an exponential cut off up to an energy 23 2 set by the column density of the gas. For example, a column density of 10 cm− will absorb X-rays up to approximately 3 keV. The effect of different column densities of neutral gas on a typical AGN X-ray spectrum is shown in Figure 1.2. If the gas is partially ionized, the opacity of the metals with the lowest atomic numbers such as carbon and nitrogen becomes negligible as they lose all of their electrons. Therefore, the X-ray spectrum of an AGN with an ionized absorber will have some flux at the lowest energies. In addition, the primary absorption edges will be from abundant elements in higher ionization states; O vii(739 eV) and O viii(871 eV) are commonly seen. In some circumstances, the absorbed spectrum is complicated by partial coverage of the continuum source by the absorber. In this case, the observed spectrum is the sum of two (or more) spectra, one with a clear line of sight to the continuum source. This type of spectrum can be observed either when the continuum source has a larger angular extent than the absorber or when X-ray photons are scattered into the line of sight around the absorber. In both situations, the observer gains an unobscured line of sight to the continuum. A spectrum with a partially covering absorber will also have additional soft flux contributed by the unobscured spectrum. With high signal-to-noise ratio data, an ionized absorber can be distinguished from partial coverage by the shape of the low energy spectrum. This straightforward picture of X-ray absorption will be complicated by significant velocity dispersion in the absorbing gas. In this case, the opacity from bound-bound transitions can no longer be ignored, and measurements of the column density that neglect this additional opacity may become unreliable. 12

The X-ray observatories utilized for the work in this thesis include ROSAT (1990– 1999; Trump¨ er 1982), the Advanced Satellite for Cosmology and Astrophysics (ASCA, 1993–2001; Tanaka et al. 1994), and the Chandra X-ray Observatory (Chandra, 1999– present; Weisskopf et al. 2000). The two primary instrument types were proportional counters (ROSAT and ASCA) and charge coupled devices (CCDs; ASCA and Chandra). Of these two types of devices, CCDs offer the advantage of higher spectral resolution (E/∆E = 50–100). Both proportional counters and CCDs can provide coverage of both the soft (0.5–2 keV) and hard (2–10 keV) X-ray bands. The softer band is more useful for studying the structure of absorption while the data from the hard band can better constrain the shape of the underlying continuum as well as probe more extreme column 23 2 densities (& 10 cm− ), particularly of BAL QSOs at higher redshifts.

1.4 Outline of the Thesis

In general, the research strategy of this thesis follows two distinct tacks. Along the first, shallow exposures of samples of several objects are analyzed in an attempt to characterize the X-ray properties of the larger BAL QSO population. Secondly, deeper X-ray observations of a few highly absorbed QSOs were undertaken to examine the specific nature of the X-ray spectra of individual objects. These two methodologies are combined in order to obtain data on the most objects given limited observatory resources while still pursuing detailed studies of a few objects. Brief descriptions of the content of each of the chapters follows. These subsequent chapters have been written as stand- alone documents, and as such the introductory material may be repetitive. Most of these chapters have been previously published; this will be noted in footnotes where appropriate.

1.4.1 Chapter 2: Exploratory ASCA Observations of BAL QSOs In this chapter, the constraints on the X-ray emission of BAL QSOs set with ROSAT are extended significantly using the hard-band sensitivity of ASCA for a sample of eight BAL QSOs. This was the first moderate-sized sample of sensitive BAL QSO observations above 2 keV, and the BAL QSOs in this sample are among the optically brightest known with B = 14.5–18.5. Despite the ability of 2–10 keV X-rays to pene- trate large column densities of gas, BAL QSOs are shown to be extremely weak sources above 2 keV, and we are only able to add two new 2–10 keV detections (0226–104 and IRAS 07598+6508) to those previously reported. By comparison with non-BAL QSOs of similar optical continuum brightness, we derive the column densities needed to sup- press the expected X-ray fluxes of these BAL QSOs. In several cases column densities 23 2 5 10 cm− are required for a neutral absorber with solar abundances. These & × were the largest X-ray column densities yet inferred for BAL QSOs, and they exceeded ROSAT lower limits by approximately an order of magnitude. Optical brightness does not appear to be a good predictor of 2–10 keV bright- ness for BAL QSOs, but these data do suggest that the BAL QSOs with high optical continuum polarizations may be the X-ray brighter members of the class. We discuss the implications of our results for future observations with Chandra and XMM-Newton. 13

If the objects in this sample are representative of the BAL QSO population, precision X-ray spectroscopy of most BAL QSOs will unfortunately prove difficult in the near future.

1.4.2 Chapter 3: Heavy Absorption in Soft X-ray Weak AGN With the complete Boroson & Green (1992) sample of Palomar-Green AGN (Schmidt & Green 1983) with z < 0.5, Brandt et al. (2000) clearly demonstrated from ROSAT data that BAL QSOs are not the only QSOs to exhibit weakness in 0.5–2.0 keV X-rays, and most of those that do have ultraviolet absorption of some sort. The ROSAT observations were unable to address the following important question: are these AGN in- trinsically X-ray weak or are they just highly absorbed? To investigate this issue further, we present ASCA observations of three of these soft X-ray weak AGN: PG 1011–040, PG 1535+547 (Mrk 486), and PG 2112+059. In general, our ASCA observations support the intrinsic absorption scenario for explaining soft X-ray weakness; both PG 1535+547 22 23 2 and PG 2112+059 show significant column densities (N 10 –10 cm− ) of absorb- H ≈ ing gas. Interestingly, PG 1011–040 shows no spectral evidence for X-ray absorption. The weak X-ray emission may result from very strong absorption of a partially covered source, or this AGN may be intrinsically X-ray weak. PG 2112+059 is a Broad Ab- sorption Line (BAL) QSO, and we find it to have the highest X-ray flux known of this class. It shows a typical power-law X-ray continuum above 3 keV; this is the first direct evidence that BAL QSOs indeed have normal X-ray continua underlying their intrinsic absorption. Finally, marked variability between the ROSAT and ASCA observations of PG 1535+547 and PG 2112+059 suggests that the soft X-ray weak designation may be transient, and multi-epoch 0.1–10.0 keV X-ray observations are required to constrain variability of the absorber and continuum.

1.4.3 Chapter 4: Markarian 231: A Prototypical Low-Ionization BAL QSO? In a Chandra snapshot survey of BAL QSOs, Green et al. (2001) demonstrated that LoBAL QSOs are significantly weaker in X-rays than normal BAL QSOs. Given this constraint, very few LoBAL QSOs will be accessible for even moderate-quality spec- troscopy with the current generation of X-ray observatories. To examine one of this class of extremely X-ray faint objects, Chandra observed Markarian 231, a LoBAL QSO and ultraluminous infrared galaxy, for 40 ks with ACIS-S3. These data reveal new informa- tion on both the starburst and QSO components of the X-ray emission of the nearest LoBAL QSO known. The bulk of the X-ray luminosity is emitted from an unresolved nuclear point source, and the spectrum is remarkably hard with the majority of the flux emitted above 2 keV. Most notably, significant nuclear variability (a decrease of 45% in approximately 6 hours) at energies above 2 keV indicates that Chandra has ∼ probed within light hours of the central black hole. Though we concur with Maloney & Reynolds (2000) that the direct continuum is not observed, this variability coupled with the 188 eV upper limit on the equivalent width of the Fe Kα emission line argues against the reflection-dominated model put forth by these authors based on their ASCA data. Instead, we favor a model in which a small, Compton-thick absorber blocks the direct X-rays, and only indirect X-rays from multiple lines of sight can reach the observer. 14

1.4.4 Chapter 5: X-ray Spectroscopy of QSOs with Broad Ultraviolet Ab- sorption With the broad bandpasses of ASCA and Chandra, we are just beginning to ac- cumulate X-ray observations with enough counts for spectral analysis at CCD resolution of QSOs with broad ultraviolet absorption lines. From a sample of eight QSOs (includ- ing four BAL QSOs and three mini-BAL QSOs) with 0.5–10.0 keV spectra with more than 200 counts, general patterns are emerging. Their power-law X-ray continua are typical of normal QSOs with Γ 2.0, and the signatures of a significant column density 23 2 ≈ [N (0.1–4) 10 cm− ] of intrinsic, absorbing gas are clear. Correcting the X-ray spec- H ≈ × tra for intrinsic absorption recovers a normal ultraviolet-to-X-ray flux ratio, indicating that the spectral energy distributions of this population are not inherently anomalous. In addition, a large fraction of our sample shows significant evidence for complexity in the absorption, i.e., a simple neutral absorber does not provide an acceptable fit to the data. In contrast, LoBAL QSOs apparently suffer from Compton-thick absorption com- pletely obscuring the direct continuum in the 2–10 keV X-ray band, complicating any measurement of their intrinsic X-ray spectral shapes.

1.4.5 Chapter 6: First Results from an Exploratory Chandra Survey of Large Bright Quasar Survey BAL QSOs In an effort to update the work of Chapter 2, we are in the process of obtaining a complete sample of exploratory X-ray observations with Chandra. In this chapter, we present the first results from this survey of Large Bright Quasar Survey BAL QSOs. This sample will eventually include all of the BAL QSOs from the LBQS with z > 1.35, the redshift at which the C iv emission line enters the LBQS bandpass. This survey is composed of short (5–7 ks) ACIS-S3 observations and was designed to significantly constrain X-ray fluxes and hardness ratios. This sample avoids the selection biases of ear- lier, heterogeneous surveys while providing X-ray data of similar quality. Of 15 observed BAL QSOs, 11 are detected. In general, our results are consistent with absorption as the primary cause of X-ray weakness for the majority of BAL QSOs. At present, neither the BALnicity index nor the detachment index shows evidence for significant correla- tions with the X-ray properties of the sample. Once again, those BAL QSOs with low- ionization absorption lines are found to be consistently X-ray weaker than BAL QSOs with only high ionization lines.

1.4.6 Conclusions In this final chapter, I briefly summarize the main results from this thesis. I also outline plans for future observations to further investigate the connection between ultraviolet and X-ray absorption and the nature of the BAL wind. 15

Fig. 1.1 A rest-frame ultraviolet spectrum of the first BAL QSO discovered, PHL 5200, with common broad emission and absorption lines and blends labeled. A low-ionization BAL QSO would also show blueshifted Mg ii absorption. These data, originally taken for the survey by Weymann et al. (1991), were kindly provided by Kirk Korista.

Fig. 1.2 The effect of different neutral column densities of absorbing gas on a typical 24 2 power-law QSO X-ray spectrum. Column densities of 10 cm− will block X-rays up to 10 keV. Cosmological redshifting of the spectrum can enable detection of harder ∼ 24 2 X-rays by current instruments. Note that for NH& 10 cm− , Compton scattering must also be taken into account. 16

Chapter 2

Exploratory ASCA Observations of Broad Absorption Line Quasi-Stellar Objects

2.1 Introduction

While Broad Absorption Line Quasi-Stellar Objects (BAL QSOs) allow us to view substantial and energetically important gas outflows that are probably present in most QSOs (e.g., Weymann 1997), their study in the X-ray regime has not yet matured. BAL QSOs are known to be much weaker in the soft X-ray band (< 2 keV) than QSOs that lack BALs (e.g., Kopko et al. 1994; Green et al. 1995; Green & Mathur 1996, hereafter GM96), and only a few BAL QSOs have reliable X-ray detections. Their low soft-X-ray fluxes may arise as a result of photoelectric X-ray absorption by the same outflowing matter that makes the BALs in the ultraviolet. Hard X-ray (> 2 keV) observations of BAL QSOs can, in principle, test the absorption hypothesis as well as constrain the properties, in particular the column densities, of BAL QSO outflows. Column density constraints for BAL QSOs are difficult to obtain from ultraviolet data due to BAL saturation (e.g., Arav 1997; Hamann 1998), and the limited X-ray data now available suggest that previous column density estimates from ultraviolet lines were too small by a factor of 100 or more, assuming the ultraviolet and X-ray absorbers are the same. ∼ Because the photoelectric X-ray absorption cross section is a strongly decreasing function of energy, BAL QSOs that are weak in the soft X-ray band (e.g., for ROSAT ) could be much brighter at higher X-ray energies (e.g., for ASCA).2 In the hard X-ray band, only the BAL QSOs PHL 5200 (z = 1.98, B 18.5) and Mrk 231 (z = 0.042, B 14.5) have ≈ ≈ been studied in detail. Mathur, Elvis & Singh (1995) argue for the presence of heavy 23 2 X-ray absorption with N 1 10 cm− in PHL 5200, although the observed flux is H ∼ × low making reliable spectral analysis difficult (see 3 for further discussion). The ASCA § 22 23 2 data for Mrk 231 also suggest a large intrinsic column density ( 2 10 –10 cm− ; ∼ × Iwasawa 1999; Turner 1999). Since at least 10% and perhaps up to 30–50% of all QSOs have BAL gas along ∼ our line of sight (e.g., Goodrich 1997; Krolik & Voit 1998), our lack of knowledge about the X-ray properties of BAL QSOs represents a serious deficiency. In an attempt to remedy this situation, we performed ASCA observations of several of the BAL QSOs

1The work in this chapter has been previously published as Gallagher, S. C., Brandt, W. N., Sambruna, R. M., Mathur, S. & Yamasaki, N. (1999) Exploratory ASCA observations of broad absorption line quasi-stellar objects. Astrophysical Journal 519:549. 2The ROSAT band is from 0.1–2.4 keV and peaks from 0.9–1.4 keV, while the ASCA band is from 0.6–9.5 keV and has high sensitivity in the 3–7 keV range. To good approximation, the 8 photoelectric X-ray absorption cross section is proportional to E − 3 . Thus a 6 keV X-ray is 120 times more penetrating than a 1 keV X-ray. ∼ 17 that seemed most likely to be X-ray bright above 2 keV. Given the unknown hard X-ray properties (e.g., fluxes and spectral shapes) of BAL QSOs, we adopted a conservative strategy of making moderate-depth ‘exploratory’ observations of a fairly large number of objects. In this manner, we could learn about the 2–10 keV properties of as many BAL QSOs as possible without being too heavily invested in the uncertain results from any one object. BAL QSOs found to be reasonably X-ray bright based on exploratory observations could then be followed up with deeper X-ray spectroscopic observations. We chose as our targets some of the optically brightest (in the B and R bands) BAL QSOs, since there is a general correlation between optical brightness and intrinsic X-ray brightness for QSOs (e.g., Zamorani et al. 1981). In this paper we present ASCA results for five of the optically brightest BAL QSOs known: PG 0043+039, 0043+008, 0226–104, PG 1700+518 and LBQS 2111–4335. In the B band, these objects are all brighter than the ASCA-detected BAL QSO PHL 5200. In addition, we include the ASCA results from archived observations of IRAS 07598+6508, Mrk 231 and PHL 5200. Redshifts, B and R magnitudes, and Galactic column densities for the BAL QSOs in our sample are given in Table 2.1. Many of our BAL QSOs are comparably bright in the optical band to radio-quiet (RQ) QSOs which ASCA has studied in significant detail (e.g., PG 1634+706 and PG 1718+481, Nandra et al. 1995; IRAS 13349+2438, Brandt et al. 1997), and our sample is on average significantly brighter at B than the high-redshift RQ QSOs studied with ASCA by Vignali et al. (1999). Aside from being some of the optically brightest BAL QSOs in the sky, the objects in our sample span a range of other properties (e.g., absorption line strength, absorption line shape, redshift, optical continuum polarization, infrared luminosity). While they do not form a rigorously defined complete sample, they do appear to comprise a reasonably representative subsample of the BAL QSO population as a whole.

As mentioned above, if BAL QSOs suffer from heavy X-ray absorption then they will be much brighter above 2 keV than at lower energies. For example, a hypothetical 23 2 BAL QSO at z = 2 with an absorption column density of 5 10 cm− is expected to × have 20 times more X-ray flux in the 2–10 keV band than in the 0.1–2.0 keV band ∼ (observed-frame with a photon index of Γ = 2.0 for the underlying power law). The 2–10 keV sensitivity of ASCA thus makes it a superior tool to ROSAT for studying the X-ray properties of BAL QSOs with X-ray column densities in the wide range from 22 2 1 24 2 (1–500) 10 cm− (column densities substantially larger than σ− = 1.50 10 cm− × T × are optically thick to Thomson scattering and are thus impenetrable to X-rays with 2 energies below mec ). Furthermore, even moderate-length ASCA non-detections can set important constraints on the 2–10 keV fluxes and internal column densities of BAL QSOs. These are important for planning Chandra/XMM-Newton spectroscopic obser- vations, and in some cases ASCA non-detections can raise the current ROSAT lower limits on BAL QSO X-ray column densities by roughly an order of magnitude (from 22 2 23 2 5 10 cm− to 5 10 cm− ). If the outflowing BAL material causes the inferred ∼ × ∼ × X-ray absorption, this increases the implied BAL QSO mass outflow rates and kinetic luminosities. 18 , t R the e Sirola Cleary 4 for (1995), (1991); curren (ks) , & GIS 28.3 20.0 41.7 20.6 22.6 16.4 23.9 23.0 E 2111–4335 al. age the Johnson v w-ionization et Exp. lo Sa er, Heiles w ev & LBQS 13 References d w (1997); sho de Ho for (ks)/ CCD CCD CCD CCD CCD CCD CCD CCD to and kman t. Mo c Neff photographic noted. Neugebauer wn 2 & Lo Log Exp. 9 as ectrum CCD 22.7/2 15.7/2 37.2/1 19.2/1 21.5/1 10.2/2 23.0/1 24.4/1 kno presen (1996), SIS e sp is hings b (1983), unication). al. y catalogue; 98 defined 98 96 94 94 96 98 98 Hutc Observing et erscripts: ma 7 is optical y comm Jul Jul Oct Jun Dec Dec Apr Mar Green sup Obs. Date and BALQSO & 09 17 29 21 21 05 23 24 ate The (1997), Murph this c de. priv ) 12 H absorption 2 al. magnitude (1992)—APM hmidt umerical i N − mo i that n i R Sc et (9) (9) (10) (11) (12) (10) (13) (11) 1 cm 1998, Information Al therein). 3.0 2.5 2.0 2.6 4.5 1.0 4.2 4.8 Irwin 20 (1992), with (PH) The t & (10 Galactic al. urnshek indicates Morris broad de. T General et erscripts: 6 Heigh BALs. (2) (4) (4) (3) (7) (3) (3) (4) mo (S. references that t indicated sup name b for QSO ain Stark /15.6 /17.7 /15.2 /13.1 /13.2 /14.4 /16.0 /17.4 McMahon /R are Pulse F and 3 (1992), 11 B (1) (3) (5) (6) (7) (1) (8) (5) in in target BAL al. suggests 17.3 18.4 15.9 18.0 14.5 16.7 15.4 14.5 umerical 1992 n et the corrected (1995) fluxes, classification done done e densities (1989), e 2.1. ers k c een ere z (1991) es ere with after b uum w 2.256 2.146 0.384 1.981 0.042 1.708 0.292 0.149 w Mey Sto able 5 al. Chaffee tin T Wilk column not & r, definitiv et i & e con (L) theses & a v ations H ations indicated w ha oltz (L) for Gunn F a from paren Morris allo are kman Boroson observ observ c in in ett, (L) arget Lo not Name T SIS ‘L’ GIS (see 2111–4335 ted (1998); 07598+6508 Hew do 5200 231 8 0043+039 1700+518 Magnitudes All An References All al. computed a b c d e Elvis, R data et presen 10 BALs magnitudes (1979) and Mrk PG PHL 0226–104 IRAS 0043+008 PG LBQS 19

2.2 Observations and Data Analysis

2.2.1 Observation Details and Data Reduction In Table 2.1 we list the relevant observation dates, exposure times and instrument modes for our objects. The data resulting from these observations were reduced and analyzed with ftools and xselect following the general procedures described in Brandt et al. (1997). We have used Revision 2 data (Pier 1997) and adopted the standard Revision 2 screening criteria.

2.2.2 Image Analyses and Count Rate Constraints We used xselect to create full (0.6–9.5 keV for SIS and 0.9–9.5 keV for GIS) and hard (2–9.5 keV) band images for each of the four ASCA detectors. We also created summed SIS0+SIS1 and GIS2+GIS3 images in the full and hard bands. Image analysis was performed using ximage (Giommi, Angelini & White 1997). We first attempted to independently check the ASCA astrometry using serendipitous X-ray sources within the fields of view. We correlated ASCA serendipitous source positions with objects in coincident ROSAT /Einstein images as well as the ned/simbad databases. We were able to independently verify the astrometry for the observations of all BAL QSOs but PG 0043+039, Mrk 231 and PHL 5200. We comment that the ASCA astrometry is generally quite reliable (see Gotthelf 1996), and our independent checking is only done for confirmation. We then searched all the images described above for any significant X-ray sources that were positionally coincident with the precise optical positions of our BAL QSOs. In order to determine the observed SIS count rate or upper limit for a given BAL QSO, the 0.6–9.5 keV counts were extracted from a circular target region with a 3.20 radius centered on the optical BAL QSO position. This provided counts for the target region, Nt. An annular or circular source-free background region was chosen and similarly extracted to give background counts. The background counts were normalized to the area of the target region to obtain Nb which was then subtracted from Nt. If the difference was 3 < 3σ where σ = √Nb, the observation was considered a non-detection. In this case an upper limit on the count rate was taken to be 3√Nt/te where te is the exposure time listed in Table 2.1. For observations with 3σ detections, the count rate was calculated ≥ as (N N )/t . A similar procedure was followed for the GIS detectors. However, for t − b e the GIS the extracted energy range was from 0.9–9.5 keV, and the radius of the target region was 50. The results of these analyses are presented in Table 2.2, and we comment on special cases in 3. Only four of these optically bright BAL QSOs were detected with § high statistical significance: 0226–104, IRAS 07598+6508, Mrk 231 and PHL 5200. We have chosen our target region sizes based upon 7.4 of The ASCA Data Reduction Guide, § and we have found that varying the target region sizes within the recommended ranges does not materially change our basic results. When calculating Nb we have investigated at least two background regions for each SIS/GIS image, and we find that our results

3The areas of the background and source regions where generally similar, and so this measure of the error in the background region is a reasonable approximation. 20 also do not materially depend upon our choice of background region. When our SIS observations were made using 2 CCD mode (see Table 2.1), we took all background photons from the chip that contained the target. We also give observed frame 2–10 keV fluxes or upper limits for our BAL QSOs in Table 2.2. These were computed for SIS0 with pimms (Mukai 1997) using a power law 23 2 with Γ = 2. The intrinsic column density was taken to be 1 10 cm− or the value from × column three of Table 2.3 (whichever is larger).4 Note that cosmological redshifting of the X-ray spectrum can in some cases greatly reduce the effect of the intrinsic absorption. This allows us to set more sensitive upper limits on the 2–10 keV fluxes for some high redshift objects due to the energy dependence of the ASCA spectral response.

2.2.3 Column Density Constraints Since X-ray spectra cannot be modeled for most of the BAL QSOs in our sample, we have followed the method adopted by GM96 to place lower limits on their intrinsic column densities. Briefly, we assumed the underlying optical-to-X-ray continuum shape of a typical RQ QSO and added an intrinsic absorbing column until the predicted count rate matched our observed count rate or upper limit. This type of analysis relies upon the plausible but unproven assumption that our targets have typical RQ QSO X-ray continua viewed through an intrinsic absorbing column of gas. Comparisons of BAL QSO and non-BAL QSO optical/ultraviolet emission lines and continua show that these two types of QSOs are remarkably similar, consistent with the view that these are not two inherently different classes of objects (e.g., Weymann et al. 1991). We also take the gas in the absorbing column to have solar abundances (Anders & Grevesse 1989), and this appears to be consistent with the currently available data (see Arav 1997 and Hamann 1998). Further justification for the assumptions that underlie this general method may be found in GM96. One additional issue not discussed by GM96 is that resonant absorption lines may be a significant source of X-ray opacity due to the large velocity dispersions of BAL QSO outflows (compare with 3 of Kriss et al. 1996). We have neglected this § effect for consistency with all previous work and because there is presently no proof that the BAL gas also causes the inferred X-ray absorption. A detailed treatment of resonant absorption line opacity in the X-ray band will be complex and dependent upon the details of the velocity dispersion in the BAL outflow, but this opacity is expected to modify the inferred column density by less than a factor of 2 (J. Krolik, private ∼ communication). Theoretical calculations examining the importance of this effect would be a valuable addition to the literature. The optical-to-X-ray spectral index, αox, is the slope of a nominal power law connecting the rest-frame flux density at 2500 A˚ to that at 2 keV (QSOs with large values of αox are those that are X-ray faint). Typical RQ QSOs are observed to have α = 1.51 0.10 (this value is from 4.3 of Laor et al. 1997 for an optical flux density at ox  § 2500 A;˚ A. Laor, private communication). The 2 keV flux density can be predicted given

4In practice, pimms is only able to work with column densities at z = 0. Therefore, to include intrinsic column densities at z > 0 in pimms we used xspec (Arnaud 1996) to find the column density at z = 0 that produces equivalent absorption. Galactic absorption was also taken into account in this process. 21

a 3 1 Table 2.2. ASCA Detector Count Rates (10− photons s− )

b Target Name SIS0 SIS1 GIS2 GIS3 Observed F2 10 − PG 0043+039 < 1.6 < 1.6 < 1.8 < 1.9 < 2 0043+008 < 2.2 < 2.1 < 2.3 < 4.4c < 1 0226–104 2.5 < 2.2 2.2 5.1 1 c c IRAS 07598+6508 < 2.3 < 3.3 < 1.7 4.2 . 4 Mrk 231 7.2 6.4 6.0 7.1 8 PG 1700+518 < 1.9 < 3.2c < 2.3 < 2.4 < 3 LBQS 2111–4335 < 4.0c < 2.1 < 2.3 < 2.5 < 2 PHL 5200 3.8 6.3 < 2.9 6.3 1

aSIS/GIS count rates are for 0.6–9.5 keV/0.9–9.5 keV, and count rate upper limits are for 3σ unless otherwise noted. b 13 2 1 F2 10 is the observed 2–10 keV flux in units of 10− ergs cm− s− computed− using a power law with Γ = 2. For the intrinsic column density 23 2 we use either 1 10 cm− or the value from column three of Table 2.3 × (whichever is larger). The Galactic column density is also included (see 2.2). § cPhoton excess above 3σ in the target region, but no evidence of a point source at the optical position of the target. See 3 for discussion. § 22

Table 2.3. Limits on Intrinsic Absorptiona

Target Assumed X-ray SIS0/GIS3 Intrinsic 22 2 b Name Photon Index NH/(10 cm− )

PG 0043+039 1.7 > 47/> 40 2.0 > 33/> 26 0043+008 1.7 > 1.8/ · · · 2.0 / · · · · · · 0226–104 1.7 / · · · · · · 2.0 / · · · · · · IRAS 07598+6508 1.7 > 72/50 2.0 > 56/36 Mrk 231c 1.8 4 PG 1700+518 1.7 > 59/> 53 2.0 > 44/> 36 LBQS 2111–4335 1.7 > 2.6/> 3.9 2.0 / · · · · · · PHL 5200 1.7 / · · · · · · 2.0 / · · · · · ·

a For αox = 1.6, unless otherwise noted. bAn ellipsis indicates that the intrinsic column density was con- 2 sistent with 0 cm− . cFor details on the photon index and column density constraints for Mrk 231 see 3, Iwasawa (1999), and Turner (1999). § 23 the rest-frame 2500 A˚ flux density and an expected value for αox. Flux densities at 2500 A˚ for PG 0043+039 and PG 1700+518 were obtained by interpolating the continuum flux density data of Neugebauer et al. (1987) with a power law. For our non-PG BAL QSOs, we derived the flux density at 2500 A˚ from the observed B magnitude (see Table 2.1) using the flux zero point of Marshall et al. (1983) and extrapolating along an assumed optical continuum power law with slope αo = 0.5. The Galactic extinction corrections used by GM96 and a K-correction for the power law were also included. An αox value of 1.6 was used to predict an underlying rest-frame flux density at 2 keV for each BAL QSO in our sample. We chose this value of αox for consistency with GM96 and note that it is reasonably conservative (also see the discussion in 4.1 § of GM96). We assumed that the underlying X-ray continuum shape was a power law with photon index Γ, and we then calculated the expected power-law normalization at 1 keV in the observed frame. We entered this model into xspec (Arnaud 1996) allowing for the presence of absorption by a column of neutral gas intrinsic to the BAL QSO. Galactic absorption was also included with the column densities given in Table 2.1. Using the simulation routine fakeit with the spectral response matrices for the ASCA SIS0 and GIS3 detectors [described in 10.4.1 of the AO-7 ASCA Technical Description § (AN 98-OSS-03)], a predicted count rate was generated. The intrinsic column density was then increased until the predicted count rate from the simulation no longer exceeded the measured count rate or upper limit. In this manner, an upper limit on the count rate yielded a lower limit on the intrinsic column density. We have focused on the SIS0 and GIS3 detectors because for the standard pointing position our targets lie closer to the optical axes of these telescopes. Combining like detectors (SIS0/SIS1 or GIS2/GIS3) would not substantially improve the column density constraints due to the fact that the vignetting for ASCA is highly dependent upon off-axis angle and is thus noticeably worse for SIS1 and GIS2. The results from this analysis are presented in Table 2.3. Since X-ray spectral shapes vary among the RQ QSO population, we have performed these calculations with both Γ = 1.7 and Γ = 2.0 in order to cover the typical range of photon index values (see Reeves et al. 1997). Using PG 0043+039 as an example, we have also illustrated the dependence of the inferred column density upon the intrinsic αox value in Figure 2.1. Recent evidence indicates that some, if not all, BAL QSOs have significantly attenuated optical/ultraviolet continua compared to non-BAL RQ QSOs (e.g., Goodrich 1997; see this paper for the precise definition of ‘attenuation’). In this case, an underlying flux density at 2 keV that is predicted from the attenuated flux density at 2500 A˚ will be artificially low. While this issue is difficult to address with precision at present, correcting for it would tend to strengthen our above results. Even larger intrinsic column densities would be required to suppress the larger predicted X-ray fluxes. For two of our BAL QSOs with near-infrared continuum flux densities in the literature, PG 0043+039 and PG 1700+518, we have examined this matter by attempting to predict the underlying 2 keV flux density from the 1.69 µm flux density (again using the data of Neugebauer et al. 1987). Laor et al. (1994) argue that the 1.69 µm flux density is a reasonably good predictor of the 0.3 keV flux density, and the Laor et al. (1997) data show that it is also a reasonably good predictor of the 2 keV flux density. We have computed αix, the spectral slope of a nominal power law between 1.69 µm and 2 keV, for the 20 QSOs from 24

Fig. 2.1 Column density lower limits for PG 0043+039 derived following the method described in the text. We show the inferred column density lower limit as a function of the intrinsic (i.e., corrected for X-ray absorption) αox. The square data points are for an X-ray photon index of Γ = 2.0 and the circular dots are for an X-ray photon index of Γ = 1.7. The dotted lines show the constraints from SIS0, and the solid lines show the constraints from GIS3. The open triangle at an intrinsic αox = 1.6 illustrates the typical BAL QSO column density lower limit found by GM96 based on ROSAT data (although GM96 did not present ROSAT data for PG 0043+039). The two vertical dashed lines show the typical range of αox for RQ QSOs (see the text). The numbers along the right hand side of the plot show the Thomson optical depth of the corresponding column density. Note that our inferred column densities are within a factor of 3 of being ∼ optically thick to Thomson scattering. 25

Laor et al. (1997) without intrinsic absorption. We find that α = 1.29 0.09, and that ix  αix has a comparable dispersion to αox for the same set of QSOs. We have repeated our column density estimates for PG 0043+039 and PG 1700+518 using αix in place of αox, and we find that the inferred column density lower limits are within a factor of two of those presented in Table 2.3. We show our results for PG 0043+039 in Figure 2.2, and 23 2 for PG 1700+518 we find an intrinsic column density > 8.5 10 cm− when α = 1.29 × ix and Γ = 2.0.

Fig. 2.2 Column density lower limits for PG 0043+039 derived following the method described in the text. We show the inferred column density lower limit as a function of the intrinsic (i.e., corrected for X-ray absorption) αix. The symbols and line styles are as for Figure 2.1, and the typical range of αix is discussed in the text.

2.3 Notes on Individual Observations

PG 0043+039: The optical and ultraviolet properties of this BAL QSO were recently studied in detail by Turnshek et al. (1994), and the ASCA analysis for it was straightfor- ward. Our large inferred X-ray column density combined with the intrinsic color excess of E(B V ) 0.11 may suggest that the absorber is dust poor (compare with 3.1.1 of − ≈ § Turnshek et al. 1994). 0043+008 (UM 275): Zamorani et al. (1981) claimed that this BAL QSO was detected in a 1.6 ks Einstein observation, although our analysis of the Einstein data made us skeptical of this claim. Wilkes et al. (1994) also give only an Einstein upper limit. 0043+008 is undetected in our much deeper (23 ks) ASCA observation. Based on our analysis, it appears that the claimed Einstein detection was actually an unrelated source lying 2.10 from the precise optical position. In order to obtain the tightest possible 26 constraints on the SIS and GIS count rates for 0043+008, we have excluded this source from the target regions before calculating the count rate upper limits. This source does not cause serious confusion, but there was a statistical photon excess in the target region for GIS3 that was not consistent with a point source or the optical position of 0043+008. We suspect this excess is due to imperfect removal of all the photons from the Einstein source. 0226–104: Despite the fact that this BAL QSO has been intensively studied (e.g., Korista et al. 1992), the coordinates stated for it in the literature are inconsistent with h m s each other. We have used the coordinates α = 02 28 39 .2, δ = 10◦11010.000 2000 2000 − (T. Barlow, private communication). 0226–104 was detected in all but the SIS1 detector as a point source (the SIS1 non-detection can be understood as due to the larger off-axis angle of 0226–104 in this detector). We excluded counts from two nearby sources before calculating the count rates. IRAS 07598+6508: This BAL QSO was marginally detected by ROSAT with a count 3 1 rate of 1.8 10− count s− , but it was impossible to determine whether the ob- ≈ × served X-ray emission was associated with accretion activity or a circumnuclear star- burst (Lawrence et al. 1997). There was an X-ray source coincident with the optical position of IRAS 07598+6508 in the GIS3 detector, and the target region counts were 8.6σ above background for the full band image (the source was also seen in the hard band image). There were photon excesses in the target regions for SIS0 and SIS1, but the lack of distinct point sources at the optical position led us to conservatively treat these as upper limits. The sole detection in the GIS3 detector might be understood as due to the better high-energy response of the GIS detectors and the smaller off-axis angle of GIS3 compared to GIS2. Our lower limit on the absorption column density is 20 times higher than that ∼ of GM96, and this suggests that the dust-to-gas ratio in the absorbing material may be extremely small (compare with 5.2 of GM96). § Mrk 231: This object has the lowest redshift in our sample and has sometimes been 46 1 classified as a Seyfert 1, but we feel its large luminosity (LBol > 10 erg s− ) and strong absorption lines justify its inclusion (e.g., it is also included in the sample of Boroson & Meyers 1992). Mrk 231 has recently been studied with ASCA by Iwasawa (1999) and Turner (1999). Our spectral analysis is consistent with the analyses in these papers and thus we do not detail it here. In Table 2.3 we give the column density discussed in 4.1.2 § of Iwasawa (1999). However, even the fitted ASCA column density may be substantially smaller than the true column density to the black hole region due to possible electron scattering and partial covering effects in this extremely complicated system (see 4 for § further discussion). PG 1700+518: There were no ASCA sources close to the optical position for this BAL QSO, and we were able to place tight upper limits on the count rate for all but the SIS1 detector. A statistical photon excess in the SIS1 target region was not from any obvious point source and thus we do not consider it a detection. A 40 ks RXTE observation of PG 1700+518 has also been performed with principal investigator P. Green. Given that our ASCA observation is & 15 times more sensitive than the RXTE observation, we expect an RXTE non-detection. 27

LBQS 2111–4335: There was clearly no point source at the optical position of the BAL QSO. However, a ninth magnitude A star (HD202042) was coincident with an X- ray source detected 2.20 from LBQS 2111–4335 (the A star may have a late-type binary companion that creates most of the X-ray emission). We excluded the counts from this source before calculating the count rate upper limits. Even after excluding the source from the extraction region, photon counts above the noise level were still evident in the SIS0 detector. However, we do not consider this to be a detection. PHL 5200: The ASCA data for this object were first presented by Mathur et al. (1995). This BAL QSO was clearly detected in all but the GIS2 detector. Due to other sources in the field, we find that the results of spectral analysis are highly sensitive to the details of the background subtraction. Given this sensitivity, our analyses suggest that any 23 2 column density from 0–5 10 cm− is consistent with the ASCA data (even at 68% × confidence). Thus, while the ASCA data certainly do not rule out the presence of a large column density absorber, they do not convincingly show one to be present either.

2.4 Conclusions and Future Prospects

We have presented the first moderate-sized sample of sensitive BAL QSO obser- vations in the 2–10 keV band. This band is potentially an important one for BAL QSO spectroscopy due to the ability of > 2 keV X-rays to penetrate large column densities, and the properties of our sample are likely to be more representative of the behavior of the BAL QSO population as a whole than those derived from the two single-object ASCA studies to date (see 1). While we detect a significantly larger fraction of our § objects than GM96, we find that in general even the optically brightest BAL QSOs known are extremely weak in the 2–10 keV band. Assuming that our BAL QSOs have typical RQ QSO X-ray continua and are weak due to intrinsic X-ray absorption, we find 23 2 column densities 5 10 cm− in several cases. These are the largest X-ray column & × densities yet inferred for BAL QSOs, being about an order of magnitude larger than the lower limits of GM96. If the same outflowing gas makes both the X-ray and ultraviolet absorption, our improved column density limits also raise the implied mass outflow rate and kinetic luminosity by about an order of magnitude. Alternatively, it is possible that BAL QSOs are intrinsically X-ray weak, perhaps due to an orientation dependence of the X-ray continuum flux. Krolik & Voit (1998) have discussed the possibility of optical/ultraviolet continuum anisotropy in BAL QSOs (due to accretion disk limb-darkening which dims objects viewed equatorially), but they suggest that anisotropy of the X-ray emission is likely to be weaker than at lower energies. If the power-law X-ray emission of RQ QSOs indeed originates in an optically thin accretion disk corona near a Kerr black hole, the intrinsic equatorial emission may in fact be enhanced by relativistic effects (e.g., Cunningham 1975). An important corollary of our work is that optical brightness (in the B or R bands) is not a good predictor of X-ray brightness for BAL QSOs. Two of the four BAL QSO X-ray detections in our sample (0226–104 and PHL 5200) are actually among our optically faintest objects. It appears that the prototype of the class, PHL 5200 (Lynds 1967), has an unusually high X-ray flux for a BAL QSO given its optical flux, and thus its X-ray properties cannot be taken as representative of the BAL QSO population. In 28 fact, its X-ray brightness is typical of non-BAL QSOs of its optical magnitude, and we do not find the evidence for a large intrinsic column density in this object to be compelling (see 3). Finally, all four of our detected BAL QSOs have notably high § optical continuum polarizations (2.3–5%; Schmidt, Hines & Smith 1997; 7 of our 8 BAL QSOs have optical continuum polarization data in the literature), and it is worth investigating if high-polarization BAL QSOs tend to be the X-ray-brighter members of the class. We refrain from making a formal claim about this issue at present due to our limited sample size, but another example of this phenomenon may be IRAS 13349+2438 (see 3.1.3 of Brandt et al. 1997). High optical continuum polarization and X-ray flux § could occur together if the direct lines of sight into the X-ray nuclei of BAL QSOs were 25 2 usually blocked by extremely thick material (say 10 cm− ). In this case, one could ∼ only hope to detect X-rays when there is a substantial amount of electron scattering in the nuclear environment by a ‘mirror’ of moderate Thomson depth. Then, any measured X-ray column density for a highly polarized BAL QSO (e.g., Mrk 231 and PHL 5200) might not reflect the true column density to the black hole region but rather the column density along the electron-scattered path (compare with 5 of Goodrich 1997). We have § also looked for other common properties of our four X-ray detected BAL QSOs and none is apparent. For example, they are not all ‘PHL 5200-like’ BAL QSOs (see Turnshek et al. 1988). Aside from the physical constraints placed upon the X-ray column densities in BAL QSOs, our results also have important practical implications for future X-ray spec- troscopy of these objects. Consider the representative case of PG 1700+518. The ASCA 3 1 SIS and GIS count rates for this BAL QSO are < 3.2 10− count s− , and its 2– 13 × 2 1 10 keV observed flux is constrained to be < 3 10− erg cm− s− . Using pimms, × 3 1 this translates into an observed Chandra ACIS-I count rate of < 9 10− count s− . × While this BAL QSO may be detectable with Chandra, high-quality spectroscopy of it will be difficult and perhaps impossible (assuming no strong variability). Even in the most optimistic case, where the flux is just below our upper limit, it will take 110 ks ∼ to obtain the 1000 counts needed for moderate-quality spectroscopy.5 A high-quality ∼ 6 ( 10 000 count) spectrum would require an immodest & 10 s of Chandra time. For the ∼ 2 1 XMM-Newton EPIC-PN, the count rate will be < 2.6 10− count s− . Spectroscopy × may be more tractable in this case, but it will still prove inordinately expensive if the flux is a factor of a few times smaller than our ASCA upper limit. We also comment that some of our implied column density lower limits are close to becoming optically thick to Thomson scattering. If the X-ray column densities in many BAL QSOs are optically thick to Thomson scattering, this will make it difficult to improve the above situation by observing at even higher energies.

Acknowledgments: We thank E. Feigelson, E. Gotthelf, J. Krolik, K. Mukai, D. Schnei- der, and an anonymous referee for helpful discussions. This paper is based upon work supported by NASA grant NAG5-4826 (SCG), the NASA LTSA program (WNB), and NASA contract NAS-38252 (RMS).

5The 1000 count and 10 000 count criteria are those given in 1.5 of the AXAF Proposers’ Guide, and we find that these criteria work well in practice. Similar§ criteria are adopted in figures 29–33 of the XMM Users’ Handbook. 29

Chapter 3

Heavy X-ray Absorption in Soft X-ray Weak Active Galactic Nuclei

3.1 Introduction

The canonical, blue X-ray loud AGN have received much attention though there are other numerically significant AGN populations (e.g., Elvis 1992). Here we study one of the ‘minority’ populations: UV-excess AGN (U B 0.44) with weak X-ray − ≤ − emission. X-ray weak AGN were first noted in Einstein data (e.g., Elvis & Fabbiano 1984; Avni & Tananbaum 1986). ROSAT extended these studies and identified a significant population of type 1 AGN notably faint in the soft X-ray band (0.1–2.0 keV) relative to their optical fluxes (e.g., Laor et al. 1997). Although not as strong as the radio- loud/radio-quiet dichotomy (e.g., Kellermann et al. 1989), soft X-ray weakness is quite a strong effect. The Laor et al. (1997) soft X-ray weak (SXW) AGN lie at least an order of magnitude below the mean of the X-ray to optical flux ratio for the Palomar Green (PG) AGN sample (Schmidt & Green 1983). Possible explanations for this distinct trait are intrinsic X-ray absorption, inherently different spectral energy distributions, and variability of the X-ray and/or optical fluxes. Only significant optical variability has been disallowed by optical monitoring (e.g., Giveon et al. 1999). SXW AGN have yet to be systematically investigated with X-ray spectroscopy in order to determine the cause for their faintness. Brandt, Laor & Wills (2000; hereafter BLW) have identified and studied the SXW AGN population in the Boroson & Green (1992; hereafter BG92) sample of all PG QSOs with z < 0.5. Specifically, they took SXW AGN to be those with the optical to X- ray spectral index, α 2.2 In soft X-rays, these sources are 10–30 times fainter ox ≤ − relative to their optical fluxes than most AGN. This effect is much more extreme than any possible luminosity dependence of αox (e.g., Green et al. 1995). For example, even the most luminous objects in Green et al. (1995) have α 1.7. ox ≥ −

1The work in this chapter has been previously published as Gallagher, S. C., Brandt, W. N., Laor, A., Elvis, M., Mathur, S., Wills, B. J. & Iyomoto, N. (2001) Heavy X-ray absorption in soft X-ray weak active galactic nuclei. Astrophysical Journal 546:795. 2 The parameter αox is the slope of a power law defined by the flux densities at rest-frame ˚ 3000 A and 2 keV. A large negative value of αox indicates weak X-ray emission. The mean value for radio-quiet AGN is 1.48 (e.g., Laor et al. 1997) with a typical range from 1.7 to 1.3 ≈ − − − (e.g., BLW). Other authors prefer α0 , the slope of a power law between the flux densities at ox rest-frame 2500 A˚ and 2 keV. For a slope of α between 2500 and 3000 A,˚ α0 =1.03α 0.03α u ox ox− u (BLW). Values of α0 from the literature have been converted to α (assuming α = 0.5) to ox ox u − facilitate direct comparisons; the original values are indicated in parentheses. 30

BLW demonstrated a strong correlation between αox and the equivalent width (EW) of intrinsic C iv absorption, suggestive of a physical link between soft X-ray weakness and UV absorption. Intrinsic X-ray absorption does depress soft X-ray flux; however, the underlying X-ray power-law continuum will recover in the harder 2–10 keV 23 2 band if N 5 10 cm− . The absorption hypothesis is thus testable with spectro- H . × scopic X-ray observations which makes such sources ideal targets for ASCA. ASCA offers access to penetrating X-rays and thus the ability to probe much larger column densities of absorbing gas than ROSAT . If soft X-ray weakness does indeed arise from intrinsic absorption, this is a significant step towards understanding the link between X-ray and UV absorbers. Precisely relating the X-ray and UV absorbers has proved difficult in practice although there is highly suggestive evidence for some connection (e.g., Mathur et al. 1998, BLW, and references therein). We have undertaken a project to observe the 10 SXW AGN studied by BLW over a broad X-ray band to test the absorption interpretation for a significant number of objects. In addition to identifying absorption, X-ray spectral analysis can constrain the nature of the absorbing gas, e.g., column density, ionization state, covering factor, and velocity distribution. These quantities are important for determining the total mass outflow rates from AGN. We proposed ASCA observations of the seven SXW AGN from BLW that had not yet been observed, and here we report on the three we were able to observe before the end of the normal ASCA operations phase. The SXW AGN in this study are more extreme in soft X-ray weakness than warm absorber AGN, and they cover a large range in UV absorption up to the strength of Broad Absorption Lines (BALs). BAL QSOs lie at the extreme end of the BLW correlation of C iv absorption EW with αox, but those examined thus far (e.g., PG 0043+008 and PG 1700+518; Gallagher et al. 1999) have not offered enough X-ray photons for significant spectroscopic work. Therefore, X-ray analyses of BAL QSOs have always assumed an underlying X-ray power law the same as that observed in normal QSOs. We aim to investigate by direct spectroscopic observations whether the range of absorption properties represents a continuum in the X-rays. If this is the case, it would support the assumption that the extreme X-ray weakness observed in BAL QSOs is due to heavy intrinsic absorption. Basic properties of the observed SXW AGN and the observations are listed in Table 3.1. Following, we present background information on each of the objects. 31 y 30 27 30 rates v ation t Murph Oct Jun No Date coun (1) Observ 1999 1999 1999 PSPC Exp. T erscripts: (ks) d sup OSA R 33.9/38.9 31.9/37.5 29.7/36.7 Time Log SIS/GIS de. and ) 2 mo umerical − n t c 1 1 2 1987) H cm N Observing 6.26 1.35 3.78 with 20 al. et (10 and b x 11 45 00 . . . o 2 2 2 indicated α − − − are 1-CCD/pulse-heigh (Neugebauer Information 2 ˚ A) in . ( iv W. a 1 19 4.6 C < BL (1995). EW fluxes densities General age 3 2 7 from v . . . erformed V p uum Sa 27 22 22 M − − − 3.1. EW tin column & ere i w con H able z ˚ A T kman 0.457 0.038 0.058 c ations Lo absorption 3000 B Galactic (2) 15.5 15.3 15.5 iv observ C for from and Name SIS/GIS (1996) 1011–040 2112+059 1535+547 W). Estimated All Rest-frame References arget al. a b c d T PG PG PG et (BL 32

1 PG 1011–040: The small Hβ Full Width at Half Maximum (FWHM) of 1440 km s− for this SXW AGN makes it a Narrow-Line Seyfert 1; its type 1 nature is supported by the presence of Fe ii emission (BG92). PG 1011–040 does not show any evidence for C iv or Lyα absorption lines in analyses of available IUE spectra (BLW) and is potentially useful for comparison with the other targets for exploring whether UV and X-ray absorption are always related. PG 1535+547 (Mrk 486, 1 Zw 1535+55): This source has strong optical Fe ii emission first noted by de Veny & Lynds (1969), and it is the X-ray weakest, detected SXW AGN in the BLW sample. It has a Narrow-Line Seyfert 1 spectrum with Hβ 1 FWHM=1480 km s− and is notable as the PG AGN with the highest optical continuum polarization (P =2.5%; Berriman et al. 1990). The structure in the polarization and C iv absorption has been studied in detail by Smith et al. (1997). PG 1535+547 was the only radio-quiet AGN observed by Neugebauer & Matthews (1998) to exhibit significant near- infrared variability from 1980–1998. PG 2112+059: This SXW AGN is one of the most luminous low-redshift (z < 0.5) PG QSOs with M = 27.3, and it has the second largest C iv absorption EW, 19 A,˚ of V − all the AGN in the BLW sample. HST spectra revealed broad, shallow C iv absorption 1 making this a BAL QSO with a balnicity index of 2980 km s− (Jannuzi et al. 1998; ≈ BLW; and see Figure 3.1), well above the lower limit that defines the class (Weymann et al. 1991). There are no spectra in the literature that cover the Mg ii region, and so it is not known whether PG 2112+059 has broad Mg ii absorption lines (see Boroson & Meyers 1992 for further discussion of Mg ii BAL QSOs) . This source was detected by ROSAT which is unusual for a BAL QSO (Kopko et al. 1994; Green & Mathur 1996). Since few BAL QSOs are even detected in X-rays, observations of this source provide a valuable test of the hypothesis that BAL QSOs have the same underlying X-ray continua as normal QSOs.

3.2 Observations and Data Analysis

In Table 3.1, we list the relevant observation dates, exposure times, and instru- ment modes for our targets, and SIS images for each source are displayed in Figure 3.2. For PG 1535+547 and PG 2112+059, we were able to use serendipitous X-ray sources in archival ROSAT PSPC data to confirm the ASCA pointing. PG 1011–040 was only observed off-axis with ROSAT ; thus the ASCA pointing could not be independently con- firmed. However, ASCA pointings are generally quite reliable, and the detected source position was coincident with the optical position within the expected error (Gotthelf 1996). The spectra resulting from these observations were extracted with xselect, a program in the ftools package, following the general procedures described in Brandt et al. (1997b). We have used Revision 2 data (Pier 1997) and adopted the standard Revision 2 screening criteria. In Table 3.2, the ASCA count rates or count rate up- per limits for each detector are listed. Count rates were determined by subtracting the counts in a source-free background region (normalized to the area of the source region) from those in the source region, and then dividing by the exposure time. The count rates in Table 3.2 and the spectra below are filtered in energy from 0.6–9.5/0.9–9.5 keV in the SIS/GIS detectors in order to eliminate poorly calibrated channels. Fluxes in the 33

Fig. 3.1 HST UV spectrum of PG 2112+059. The data from 1617–3277 A˚ (observed frame) have been previously published (Jannuzi et al. 1998), but the shorter wavelength data are presented here for the first time. Observed wavelength ranges, observing modes, and observation dates are as follows: 1327–1613 A,˚ GHRS G140H grating, 1995 July 31; 1598–2311 A,˚ FOS G190L grating, 1992 September 19; and 2223–3277 A,˚ FOS G270H, 1992 September 19. The data have been dereddened using the Galactic NH given in 21 2 Table 3.1 and E(B V ) = N /(5.0 10 cm− ) (Diplas & Savage 1994; Predehl & − H × Schmitt 1995). The GHRS and FOS data are in excellent agreement in the region of wavelength overlap even though the flux-density calibration may be uncertain for these small slit observations. The combined data as presented here show more clearly than previously published spectra the shallow and very broad absorption lines against the QSO continuum interpolated between, e.g., 1600 and 2300 A.˚ There may be up to four equally spaced line-locked absorption systems identifiable in the O vi λλ1031,1037 and C iv λλ1548,1550 doublets and possibly some identified in Lyα and N v λλ1238,1242 (see Foltz et al. 1987 for an excellent example of line-locking in a BAL QSO). 34

Fig. 3.2 Full-band (0.6–9.5 keV) co-added SIS0 and SIS1 images for each of the SXW AGN in this study. Images have been adaptively smoothed at the 3σ level using the algorithm of Ebeling, White & Rangarajan (2000). The white cross marks the precise optical position of each source. The white circle is centered on the X-ray source with a radius of 4000 representing the standard ASCA pointing error.

0.5–2.0 keV and 2.0–10.0 keV bands are calculated from the best-fitting models with the SIS0 normalization (see Table 3.2). Unless otherwise noted, we used the rest-frame 3000 A˚ continuum flux densities from Neugebauer et al. (1987) and the rest-frame 2 keV flux densities from the best-fitting SIS0 model to calculate αox. Galactic absorption is fixed in all of the model fits, and errors stated are always for 90% confidence conserva- tively taking all model parameters to be of interest other than absolute normalization. 1 1 1 We adopt H0 = 70 km s− Mpc− and q0 = 2 throughout.

3.2.1 PG 1011–040 PG 1011–040 was detected in all instruments except GIS2 at greater than 3σ above the background. Low signal-to-noise ratio spectra were successfully extracted from SIS0 and GIS3 for analysis with the X-ray spectral analysis tool xspec (Arnaud 1996). There were not enough counts in the SIS1 detector to warrant spectral analysis. The background spectrum for SIS0 was extracted from a region on the CCD free from point source contamination, while that for GIS3 was extracted from blank-sky background files as detailed in Ikebe et al. (1995) in order to mitigate vignetting effects. The spectra were first fit with a power law with fixed Galactic absorption. The fit was statistically acceptable, and the resulting photon index, Γ 2 (see Table 3.3), is ≈ consistent with those typically observed in radio-quiet type 1 AGN, Γ = 1.5–2.5 (Brandt et al. 1997a; Reeves & Turner 2000). However, PG 1011–040 has an extremely low 41.8 1 2–10 keV X-ray luminosity, L2 10, of 10 erg s− . In order to determine if X-ray absorption is the cause of the w−eakness of PG 1011–040 in soft X-rays, we examined the spectrum for any indications of absorption. If the intrinsic neutral column density 23 2 were . 10 cm− , then the X-ray spectrum would start to recover to the level of the incident power law by 3–5 keV. Such a recovery is not evident in the spectrum, and ∼ thus we do not see evidence for absorption of this magnitude (see Figure 3.3). Given the 35 t c all x 03 75 02 . . . o SIS0 2 1 2 α − − − with represen ) 1 eV 9 6 3 − . . . k s 0 0 0 est-fitting 2 Errors b    − 1 5 1 . . . cm the 4 7 1 text. 2.0–10.0 normalization erg and the del 13 2 3 eV − in . . 13 12 . . k 0 0 mo 0 ed (10 1987) +0 −   b SIS/GIS. 4 7 59 the . . . al. 2 0 0 for in Flux 0.5–2.0 et describ arameters eV P 8 k 8 as . . 7 0 0 . range 0   dels  0 ASCA . ) 8 1 mo . − ed 7/14 7/9 s . (Neugebauer . 0 7/2 0 . ct SIS0 1  GIS2/GIS3 confidence  3 6 7 < − . 0.6–9.5/0.9–9.5 . Observ fluxes 9 6 90% (10 a 8 6 . . uum 3.2. rates; 0 0 the est-fitting 1 t Rate b tin    t en 3 7 able . . con the T coun giv ˚ A 1/16 8/6 7/4 Coun . . for 0 0  SIS0/SIS1   are fluxes 21 3000 3 7 . . fixed. 6 5 del from fluxes mo in Name kground-subtracted parameters tegrated 2112+059 1011–040 1535+547 range In Bac Calculated arget del. a b c T PG PG PG mo other the 36

Fig. 3.3 ASCA SIS0 and GIS3 observed-frame spectra of PG 1011–040 fit with a power- law model. Filled circles are the data points for the SIS0 detector while plain crosses are for the GIS3 detector. The ordinate for the lower panel, labeled χ, shows the fit residuals in terms of σ with error bars of size one.

21 2 current data, we constrain the intrinsic column density to be N 5 10 cm− , though H ≤ × the poor photon statistics prevent detailed modeling of the data. Another possibility is that the observed power law may not be the direct nuclear source, but only an indirect scattered component. If this is the case, then the direct component can reasonably be modeled as an additional power law with Γ = 2.0 and the expected 1 keV normalization for an AGN of this optical brightness assuming an intrinsic αox of 1.5. A large neutral 24 2 − column density of N > 1.0 10 cm− is then required to extinguish the primary H × continuum throughout the ASCA bandpass. If the observed X-ray emission is only scattered light, one might expect a narrow iron Kα emission line; given the poor quality of the data, the EW for a line in the range from 6.4–6.97 keV can only be constrained to be < 6.6 keV. A nearby source 2.05 away in the SIS1 and GIS detectors complicated the mea- surement of count rates. We searched the Palomar Optical Sky Survey for an optical counterpart to this nearby source, but none was apparent. The confusing source was quite faint in SIS1, and the lack of its detection in SIS0 can be understood as resulting from the different position of each detector with respect to the optical axis of the tele- scope. For GIS2, the count rate upper limit for PG 1011–040 in Table 3.3 is 3√Ns/t, where Ns is the number of counts in the source cell and t is the exposure time. Counts from the nearby source mentioned above were excluded. We have investigated potential variability between the 350 s ROSAT All-Sky Survey observation and our ASCA exposure. Although the constraints are not tight, we do not find highly significant evidence for variability. 37

3.2.2 PG 1535+547 (Mrk 486) PG 1535+547 was clearly detected in all four detectors with enough total counts, 940, for spectral analysis. The source spectra were extracted using circular source ≈ cells with radii 2.07/4.01 for the SIS/GIS to maximize the signal-to-noise ratio, and the background regions were extracted as described in 3.2.1. The following spectral analysis § was also done using different background regions on the CCDs for the SIS as well as source-free background regions on the GIS. The results were consistent and thus are not sensitive to the choice of background region. For the first fit to the data in all four detectors, a power law with fixed Galactic +0.17 absorption was used. The resulting photon index, Γ = 0.45 0.22, is abnormally flat for a typical type 1 AGN. In addition, the residuals showed a −systematic excess below 1 keV indicating that the fit was poor. A χ2 of 134.7 for 119 degrees of freedom was obtained indicating that this model can be rejected with 84% confidence according to the P (χ2 ν) probability distribution as defined in 11.1 of Bevington & Robinson (1992; see | § Table 3.3). Next, we fit a power law with fixed Galactic absorption also allowing for the +0.25 presence of intrinsic absorption. The fitted photon index was again low, Γ = 0.44 0.32, with negligible intrinsic absorption and χ2=134.7 for 118 degrees of freedom; this mo− del can be rejected with 86% confidence. A plot of the residuals again showed an excess below 1 keV. Forcing the photon index to lie in the typical range for radio-quiet AGN, Γ =1.5–2.5, only worsened the fit. Finally, we fit a model adding intrinsic neutral absorption partially covering the incident power law. The choice of this model was motivated by the presence of low energy photons which would be completely absorbed by total coverage of the emission source by neutral gas. An effective partial covering situation could result from two lines-of-sight to the source of the X-ray emission. In this scenario, the first component would be direct emission absorbed by gas along the line-of- sight, and the second component would be unabsorbed emission scattered by electrons towards the observer. Since scattering results in an increased fraction of polarized flux, the high observed optical polarization of PG 1535+547 suggests that scattering is indeed occurring in the nucleus. For the partial covering model, the resulting best-fitting model +0.83 23 2 parameters were an intrinsic neutral absorption column density of (1.23 0.54) 10 cm− , +7 +0.92 − 2 × a covering fraction of 91 26%, and Γ = 2.02 0.95 (see Figure 3.4). The χ was 117.5 for 117 degrees of freedom,−a significant impro−vement over all of the previous fits at the 99% confidence level according to the F -test. The residuals showed no obvious trends, and the partial covering model explains the additional soft photons evident in the simple absorption models while also providing a physically reasonable photon index. Confidence contour plots for the two partial covering parameters are displayed in Figure 3.5. From 42.1 1 this best-fitting model to the ASCA data, we obtain L2 10=10 erg s− (absorption- 42.4 1 − corrected L2 10=10 erg s− ). For completeness, we also tried an ionized absorber model, but the− fit left clear systematic residuals. The partial-covering absorption model was preferable, with ∆χ2 = 17 for the same number of degrees of freedom. − 38

Table 3.3. Parameters of Model Fits to the ASCA Data

Modela Γ Parameter Values χ2/ν P (χ2 ν)b | PG 1011–040 1. Power law +0.44 1.93 0.39 ...... 53.6/47 0.24 − PG 1535+547 1. Power law +0.17 0.45 0.22 ...... 134.7/119 0.16 2. Power law−with intrinsic partial-covering absorption +0.92 22 2 +8.3 2.02 0.95 NH (10 cm− ) 12.3 5.4 117.5/117 0.47 − +0− .07 Coverage Fraction 0.91 0.26 PG 2112+059 − 1. Power law +0.08 1.44 0.09 ...... 182.8/151 0.04 2. Power law,− E > 3 keV rest frame (2 keV observed frame) +0.23 1.94 0.21 ...... 69.0/76 0.70 3. Power law−with intrinsic, neutral absorption +0.26 22 2 +0.51 1.97 0.23 NH (10 cm− ) 1.07 0.37 142.2/150 0.66 4. Power law−with intrinsic partial-covering absorption − +0.40 22 2 +1.4 1.98 0.27 NH (10 cm− ) 1.0 0.49 142.0/149 0.65 − −+0.03 Coverage Fraction 0.97 0.26 5. Power law with an ionized absorption edge − +0.28 +0.09 1.88 0.24 Rest-frame Edge Energy (keV) 0.66 0.15 142.7/149 0.63 − +−>10 Maximum Optical Depth, τ 9.7 7.9 −

aAll models have fixed Galactic absorption as given in Table 1. Models with intrinsic absorption have z fixed and assume neutral gas with solar abundances unless otherwise noted. bThe probability, if the given model were correct, that this value of χ2/ν or greater would be obtained where ν is the number of degrees of freedom. 39

Fig. 3.4 ASCA SIS and GIS observed-frame spectra with the best-fitting model for PG 1535+547 (see 3.2.2). Filled circles are the data points for the SIS detectors while § plain crosses are for the GIS detectors. The ordinate for the lower panel, labeled χ, shows the fit residuals in terms of σ with error bars of size one.

According to the ASCA data, the measured α = 2.03 (see Table 3.2) is greater ox − by ∆αox=0.42 than that calculated in BLW from the 2678 s ROSAT PSPC observation performed on 1993 February 8. This dramatic change corresponds to a factor of 13.5 ≈ increase in flux for a fixed X-ray spectral shape. However, a closer look at the PSPC data indicates that PG 1535+547 was not actually detected in the 0.5–2.0 keV band but only in the 0.1–0.5 keV band. In the 0.5–2.0 keV band, there were 5 photons in the 1.02 radius source cell with 2 expected from the background. Using the method of Kraft, Burrows & Nousek (1991), the upper limit on the number of counts from the target is 8.54 at the 95% confidence level. This corresponds to an observed count rate 3 1 of < 3.2 10− ct s− . To estimate α , BLW assumed a photon index of Γ = 3.1 × ox based on the Hβ FWHM (see 3.3 of Laor et al. 1997 for discussion) and then found § the normalization that gave the proper count rate for the entire 0.1–2.5 keV band. If instead the 0.5–2.0 keV count rate upper limit is used to estimate an upper limit on the 2 keV flux density, then α is measured to be < 2.17. The lower limit on the flux ox − increase for a fixed spectral shape is then a factor of 2.4. Though the limited photon statistics in the ROSAT observation preclude spectral analysis, we can still compare the two observations in order to determine if the increase in αox is greater than expected from measurement and systematic errors. We took the best-fitting ASCA SIS0 spectral model and convolved it with the ROSAT PSPC response using xspec. From this, we +1.1 3 1 predicted a PSPC count rate of (4.9 ) 10− ct s− in the 0.5–2.0 keV band (errors 0.9 × are for 90% confidence), clearly abov−e the measured upper limit. Though the absolute calibrations of the two satellites have a systematic offset, measured fluxes in the 0.5– 2.0 keV band are generally consistent within errors (Iwasawa et al. 1999); this cannot 40

Fig. 3.5 Confidence contours for the two partial-covering parameters in the spectral analysis of PG 1535+547. The contours are for 68%, 90%, and 99% confidence. In the plot on the left, the photon index, Γ, is allowed to vary freely. On the right, Γ is limited to the range 1.5–2.5 to keep it within the bounds observed for Seyfert 1 galaxies. Since intrinsic absorption can masquerade as a flattening of the power-law slope, a lower limit on Γ narrows the range of acceptable covering fractions. account for the increase in brightness. The source has varied significantly in the six years between the two observations.

3.2.3 PG 2112+059 In all four ASCA instruments, PG 2112+059 was clearly detected with enough counts, a total of 1950, for spectral analysis. Source spectra were extracted as de- ≈ scribed above using circular source cells with radii 2.06/3.08 arcmin for the SIS/GIS. The GIS source cell was chosen slightly smaller than in 3.2.2 in order to avoid including § photons from a faint source located 4.01 from the target. Background spectra were ≈ also extracted as described in 3.2.1, and we have verified that our results below are not § sensitive to the details of the background subtraction. For the first fit to the data in all four detectors, a power law with fixed Galactic absorption was used. A χ2 of 182.8 for 151 degrees of freedom was obtained indicating that this model can be rejected with 96% confidence. The resulting photon index, Γ = +0.08 1.44 0.09, is rather flat for a typical QSO. In addition, the residuals showed a systematic deficit− below 1 keV which also indicated that the fit was not good in the soft band and suggested the presence of intrinsic absorption. In order to model only the continuum, we excluded the data below 3 keV in the rest frame (2 keV in the observed frame). Above 3 keV, the absorption of incident photons by neutral or partially ionized gas should be much less than in the softer band, thus allowing the underlying power-law continuum to be observed and fit directly. These data were then fit with a power law plus Galactic 41

Fig. 3.6 ASCA SIS and GIS observed-frame spectra of PG 2112+059 fit with a power-law model above 2 keV which is then extrapolated back to lower energies. Filled circles are the data points for the SIS detectors while plain crosses are for the GIS detectors. The ordinate for the lower panel, labeled χ, shows the fit residuals in terms of σ with error bars of size one. Note the low-energy residuals suggestive of a warm absorber or partial covering model.

+0.23 absorption. This resulted in a steeper photon index of Γ = 1.94 0.21, more characteristic of radio-quiet QSOs, and χ2=69.0 for 76 degrees of freedom. The− data below 3 keV were then included, and the model extrapolated to the lowest ASCA energy bins (0.87 keV rest frame). The resulting plot, shown in Figure 3.6, is suggestive of the possible nature of the intrinsic absorber. The spectra appear to recover towards the power-law continuum in the lowest energy bins, hinting at either partially ionized absorbing gas or partial covering of the continuum. In contrast, a neutral, cold absorber would depress almost all flux below 3 keV, the energy where the continuum recovers. We fit several models ≈ to try to characterize the absorbing gas. The details of each fit are listed in Table 3.3. 22 2 Briefly, intrinsic absorption by neutral gas with column density 1.1 10 cm− was × preferable at the > 99% confidence level to a straight power law with only Galactic absorption according to the F -test (∆χ2 = 40.6). In addition, adding neutral intrinsic − +0.26 absorption resulted in a photon index of Γ = 1.97 0.23 which is consistent with Γ for the fit to the data above 3 keV in the rest frame (see− Figure 3.7). Fitting the data with a warm absorber or a partial covering model also yielded acceptable fits to the data; however, the fits were not statistically significant improvements over that with a neutral absorber (see Table 3.3). An iron Kα line is not detected; we can set an upper limit on the EW of a narrow line from 6.4–6.97 keV of < 210 eV. 43.7 PG 2112+059 is a luminous X-ray source with L1 2 10 and L2 10 44.6 1 − ≈ − ≈ 10 erg s− , and so any possible starburst X-ray contribution would negligibly influ- ence our spectral fitting. Though absorption of one form or another is certainly required, 42

Fig. 3.7 Confidence contours in the spectral analysis of the BAL QSO PG 2112+059 for the two parameters of interest in the absorbed power-law model (see model 3 in Table 3.3). The contours are for 68%, 90% and 99% confidence. Note that even within the 99% confidence contours Γ is within the range for normal QSOs. these data do not have a high enough signal-to-noise ratio to distinguish between var- ious absorption models. Regardless of the nature of the absorber, the photon index is consistently 2.0, completely typical for a radio-quiet QSO. We have investigated the ≈ impact of the changing SIS calibration on the results of our spectral-fitting (T. Yaqoob 1999, private communication) and find that our general conclusions are not substantively affected. The archival ROSAT PSPC data from a 21.1 ks observation on 1991 November 15 cannot be simultaneously fit with the ASCA data without allowing the power-law normalization to drop by a factor of > 7. If we convolve the ASCA SIS0 model with the ROSAT PSPC response, then the predicted 0.5–2.0 keV count rate for the ROSAT +0.13 2 1 observation is (1.83 0.16) 10− ct s− . This value is almost a factor of four larger than − × 3 1 the observed PSPC count rate in the same band of (4.8 0.7) 10− ct s− . The ROSAT  × data have a low signal-to-noise ratio, but spectral analysis allows for a measurement of the 2 keV flux density which is relatively insensitive to the specific model. Recalculating αox more robustly from spectral fitting of the data (rather than assuming a model and using the PSPC count rate as in BLW), one obtains α = 2.07, similar to the BLW value ox − of 2.11. The measured α = 1.75 from the best-fitting ASCA data is significantly − ox − higher than any αox consistent with the ROSAT data; ∆αox of 0.32 between the two observations corresponds to an increase in the rest-frame 2 keV flux density by a factor of 7.2. The discrepancy between the relative increases in the 2 keV flux density and the ≈ observed versus predicted PSPC count rates can be reconciled if the X-ray spectral shape has varied, perhaps due to changes in the absorbing material along the line-of-sight. For example, a change in ionization state of the material along the line-of-sight could strongly affect the rest-frame 2 keV flux density but cause smaller changes to the observed frame 43

0.5–2.0 keV integrated flux. Clearly, PG 2112+059 has brightened substantially in the soft X-ray band in the decade since the ROSAT observation. In fact, PG 2112+059 is no longer SXW as defined by BLW.

3.3 Discussion and Conclusions

3.3.1 PG 1011–040 From our rudimentary spectral analysis, PG 1011–040 appears to have a typical power-law X-ray continuum with no evidence for absorption beyond the Galactic column density. Forcing a normal underlying AGN X-ray spectrum in addition to the observed 24 2 component requires an intrinsic neutral column density of NH > 10 cm− to extinguish it to the top of the ASCA bandpass; such a large column density would be extraordinary for a type 1 AGN. Detecting a narrow iron Kα emission line with a large EW would provide evidence for such a situation, but we are unable to set meaningful constraints with the current data. Given that the exact relationship between UV and X-ray absorbers is still unclear, such heavy absorption can perhaps be reconciled with the lack of detected C iv absorption if the gas is dense and cold or very close to the nucleus. Dense, cold gas can effectively absorb X-rays without creating strong absorption in the high-ionization C iv line as can gas lying interior to the UV emission region. Alternatively, this may be an intriguing example of an AGN with intrinsically weak X-ray emission. However, the normal He ii λ4686 emission (EW= 10.2 A;˚ BG92) suggests that the intrinsic X-ray emission cannot be extremely abnormal (Korista, Ferland & Baldwin 1997; BLW). A depressed ionizing continuum in the extreme UV (EUV) would not produce significant He ii emission, and it is difficult to reconcile a normal EUV continuum with drastically reduced X-ray emission. Regardless, PG 1011–040 merits further study in both the X- ray and UV regions of the spectrum in order to more tightly constrain the apparent lack of absorption in both bands (see Figures 3.8 and 3.9). In addition, UV emission line studies of this object are important for further investigating the intrinsically weak X-ray continuum scenario. The ASCA data provide a measurement of αox consistent with that derived from the ROSAT PSPC count rate, and so this source remains firmly in the SXW category. From archival Einstein data, α = 1.89 (α0 = 1.93) was measured (Tananbaum ox − ox − et al. 1986), indicating that this source has been consistently weak in soft X-rays for at least two decades. IUE data from 1993 and 1995 show no evidence for UV variability.

3.3.2 PG 1535+547 (Mrk 486) PG 1535+547 shows clear signs of intrinsic absorption by a large column density of 23 2 gas (N 1.2 10 cm− ) with only partial covering of the power-law continuum. The H ≈ × partial covering model is particularly interesting with regard to the optical polarization properties of this AGN since it implies possible electron scattering of X-rays around intrinsic absorbing clouds. This result is consistent with the known high polarization of the optical continuum flux (P =2–8%; Smith et al. 1997) and closely resembles the X-ray results for the SXW AGN PG 1411+442 (Brinkmann et al. 1999). PG 1411+442 23 2 has a comparable internal column density, N =1.3 10 cm− , and covering fraction, H × 44

+0.02 fcov=0.97 0.09. From the broad-band spectral energy distribution of PG 1535+547 (see Figure 3.8),− the reddened continuum in the optical and UV is clearly seen, as well as the recovery of the X-ray continuum to within the range expected given the 3000 A˚ flux density. The amount of reddening in the optical and UV, though significant, is still much less than the A 76 expected along the direct line-of-sight if the absorbing material V ≈ has roughly a Galactic dust-to-gas ratio (Burstein & Heiles 1978). Though dust may be responsible for the optical polarization, it is an inefficient scattering medium for X-rays; electrons are the most likely candidates for X-ray scatterers. The increase in X-ray brightness of PG 1535+547 in the six years between the ROSAT and ASCA observations is notable. Determining whether the continuum flux has increased or the absorption has decreased is not possible given the limited ROSAT data. During the same epoch, the near-infrared flux of PG 1535+547 in 4 bands from 1.27–3.7 µm also varied significantly, increasing by 0.5 mag from 1993 to 1996 and ≈ then dropping by 0.5 mag between 1996 and 1998 (Neugebauer & Matthews 1999). ≈ Our analysis of archival UV data from IUE indicates that on 1982 November 17 and 1984 May 7, PG 1535+547 was 30% brighter than in earlier IUE spectra (1982 June 6 and September 26). Later HST spectra (1992 September 19) are consistent with the earliest IUE data.

3.3.3 PG 2112+059 With this ASCA observation, PG 2112+059 is found to have the largest X-ray flux of any BAL QSO known. For comparison, the next brightest ASCA-observed BAL QSO, PHL 5200, had broad-band count rates 3 times lower in all detectors (Mathur et al. ≈ 1995; Gallagher et al. 1999). Our result of a typical QSO X-ray continuum with absorp- tion is robust. Regardless of the nature of the absorbing gas, neutral, partially ionized or partially covering, the photon index is 2.0. This is the first direct spectral evidence ≈ for a typical QSO X-ray continuum in a BAL QSO and supports the assumption that they are indeed normal QSOs cloaked by absorbing gas (e.g., Green & Mathur 1996; Gallagher et al. 1999). Our best-fitting model to the ASCA data finds an intrinsic, 22 2 neutral column density of 10 cm− . The data are not of sufficient quality to deter- ≈ mine the ionization state of the gas; ionized gas would require a larger column density for the same amount of continuum absorption. An additional issue complicating the determination of the line-of-sight column density is the unknown velocity dispersion of the gas. Currently available models for X-ray absorption by neutral or partially ionized gas do not usually account for any velocity dispersion within the absorber. Thus they are severely limited for the study of BAL QSOs which may have velocity dispersions as 4 1 large as 10 km s− if the X-ray and UV absorbing gas have similar velocity structure. ∼ At these high dispersions, bound-bound X-ray absorption becomes significant and may dominate over bound-free absorption (see 2.2 of Netzer 1996; D. Chelouche & H. Netzer § 2000, private communication). The measured neutral column density for PG 2112+059 is more than an or- der of magnitude lower than the intrinsic column densities implied for the BAL QSOs PG 0946+301 (Mathur 2000) and PG 1700+518 (Gallagher et al. 1999), and it may be that this lower column density is related to the relatively shallow BAL troughs of this 45 object (see Figure 3.1). Future observations of a larger sample of BAL QSOs can test if the depths in BAL troughs are correlated with the observed X-ray column density, but care must be taken to evaluate the amount of scattering which can partially ‘fill in’ BALs. The measured column density also appears to be insufficient for the ‘hitchhiking’ gas postulated by Murray et al. (1995). In their model, thick, highly ionized gas with 23 2 N 10 cm− absorbs soft X-rays and shields the BAL wind from becoming com- H ≈ pletely stripped of electrons. The X-ray absorbing gas is stalled at the base of the flow and is distinct from the gas causing the broad UV absorption lines. Though our data suggest that this is not the case for PG 2112+059, higher spectral resolution observations are required to constrain the velocity of the X-ray absorbing gas. From Einstein data, Tananbaum et al. (1986) measured an upper limit of αox < 1.78 (α0 < 1.82) which is weaker than the current value, but consistent with the − ox − ROSAT observations. Since then, the variability of α from 2.07 to 1.75 over the ox − − course of a decade makes further UV studies as well as higher signal-to-noise ratio X- ray observations of this source important. In particular, the demonstrated variability of PG 2112+059 in the X-rays can be compared with the UV BALs. The probable X- ray spectral variability hints at changes in absorption structure, either of the ionization parameter or the column density. BAL QSOs generally have little variability in their UV absorption line profiles (e.g., Barlow 1993). However, if the UV and X-ray absorbing gas are closely related and the increase in X-ray flux is a result of changing absorption structure, then the dramatic change in αox should occur in concert with significant changes in the UV absorption lines. The gravitationally lensed QSO PG 1115+080 is an example of a mini-BAL QSO which appears to show variable X-ray absorption (Chartas 2000b) as well as clear changes in UV absorption (Michalitsianos, Oliversen & Nichols 1996). The archival HST FOS data of PG 2112+059 were taken in 1992, approximately one year after the 1991 ROSAT observation and seven years before this ASCA observation. PG 2112+059 is scheduled for a possible Cycle 9 HST observation as part of a UV absorption snapshot survey. From the spectral energy distribution of PG 2112+059 (see Figure 3.8), it is apparent that the HST UV spectra are not continuous with the Neugebauer et al. (1987) UV photometric data points which we used to calculate α in Table 3.2. In the 13 ox ≈ years between the 1980 N87 measurements and the HST FOS observation, it is likely that the UV/optical flux has diminished. Optical variability is not unreasonable as between 1992 and 1996, part of the time between the ROSAT and ASCA observations, PG 2112+059 brightened by 0.18 mag in the V band (Rabbette et al. 1998). Additional photometric data up to the present epoch show continued variability of 0.07 mag in  V (Garcia-Rissmann et al., in preparation). Analysis of IUE and HST data from 1986 to 1995 shows no evidence for UV continuum variability. If instead of using N87 values the HST continuum data are extrapolated to rest-frame 3000 A,˚ the ROSAT and ASCA values for α are 1.94 and 1.62, respectively. Our ASCA absorption-corrected α is ox − − ox then 1.55, within 1σ of the mean of the Laor et al. (1997) radio-quiet AGN distribution − of αox. This is further evidence supporting the picture of BAL QSOs as AGN with typical underlying X-ray continua. Since the strongest signatures of X-ray absorption are edges at fairly low energies (lesssim2 keV), the characteristics of the X-ray absorbing gas are most readily studied 46 in low-redshift objects where rest-frame photons below 2 keV can still be detected ≈ 43.7 1 with available X-ray instruments. With its high X-ray luminosity, L1 2 10 erg s− 44.2 1 44.6 − 1≈ (absorption-corrected L1 2 10 erg s− ) and L2 10 10 erg s− , PG 2112+059 − ≈ − ≈ is unlikely to suffer significant contamination by a nuclear starburst as the most X-ray luminous starburst galaxy known is more than 200 times fainter (Moran et al. 1999). Two other relatively X-ray bright BAL QSOs, CSO 755 (z = 2.88) and PHL 5200 (z = 1.98), are at much higher redshifts, and thus PG 2112+059 offers the most promise for detailed absorption studies with the new generation of X-ray observatories. With higher spectral resolution and greater sensitivity at low energies, parameters such as the ionization state and column density of the absorber can be probed more thoroughly. Ideally, grating observations of this source have the potential to elucidate the dynamical structure of the X-ray absorbing gas.

3.3.4 X-ray Absorption as the Primary Cause of Soft X-ray Weakness The number of SXW AGN in BLW with published ASCA spectral constraints is now six out of ten: PG 0043+008, PG 1011–040, PG 1411+442, PG 1535+547, PG 1700+518, and PG 2112+059. As the number of SXW AGN observed in the 2– 10 keV band grows, the general identification of soft X-ray weakness with intrinsic X-ray absorption becomes more plausible. PG 1411+442 had been shown previously to have 23 2 substantial X-ray absorption with a neutral column density of N 1.3 10 cm− H ≈ × (Brinkmann et al. 1999). With two of our sample, PG 1535+547 and PG 2112+059, we have also examined objects with more X-ray absorption than is typically seen in type 1 AGN but still with enough photons for direct spectroscopic analysis. The BAL QSOs PG 0043+008 and PG 1700+518 were not detected by ASCA which suggests that they 23 2 suffer absorption by a column density larger than 5 10 cm− (Gallagher et al. ≈ × 1999). PG 1011–040 is the only object to show no evidence for strong UV or X-ray absorption. With the exception of perhaps PG 1011–040, the SXW AGN which have been spectroscopically studied in the hard band show no evidence for unusual intrinsic spectral energy distributions. In addition, the occurrence of UV and X-ray absorption is now one- to-one; we have found evidence for X-ray absorption in those sources with detected UV absorption. These observations lend direct support to absorption as the primary cause for soft X-ray weakness as put forward by BLW. After correcting the observed X-ray spectra of PG 1411+442, PG 1535+547, and PG 2112+059 for intrinsic absorption, one obtains photon indices and αox values consistent with typical type 1 AGN. This serves to emphasize the universality of X-ray emission in QSOs, and the value of αox as a parameter for identifying QSOs with interesting UV and X-ray absorption properties. The three AGN in this study and PG 1411+442 are ideal targets for higher throughput X-ray spectroscopy and additional UV observations. The most compelling issues to address remain the nature of the relationship between the X-ray and UV ab- sorbers and the dynamics of the X-ray absorbing gas. Since the bulk of the gas in the inner regions may be most readily observable in X-rays, it is essential to measure the velocity of this gas to determine the mass-outflow rates of AGN. 47

3.3.5 X-ray and UV Variability

The αox variability of two SXW AGN, PG 1535+547 and PG 2112+059, suggests that other sources in the intermediate range ( 2.0 < α < 1.7) may also vary and − ox − perhaps even join the ranks of SXW AGN at times. One possible example of such interesting variability is PG 0844+349 which is generally not SXW, but in one out of five X-ray observations had an intermediate α of 1.88 (α0 = 1.92; Wang et al. ox − ox − 2000). Additional 0.1–10.0 keV X-ray observations of the sources known to vary in αox are essential for determining whether the absorbing gas or the underlying continuum it- self is changing. PG 2112+059 appears to be an example of a SXW AGN whose spectral shape has changed over time most likely as a result of changing absorption structure (see 3.2.3). Simultaneous or near-simultaneous UV observations are also required to § elucidate the precise relationship of the UV and X-ray absorbers and calculate αox con- sistently. Currently, we have direct evidence for variability in X-ray flux which translates into horizontal motion in the BLW C iv absorption EW versus αox plot (see Figure 3.9). As of yet, we do not have corresponding multi-epoch C iv absorption EW measure- ments and cannot determine if the plot is as dynamic vertically. Though monitoring of UV BALs by Barlow (1993) showed little variability, the matter requires further in- vestigation. If the variability is a result of bulk motions of absorbing clouds across the line-of-sight, then the AGN currently towards the absorbed end of the correlation may show more dramatic changes over time as their clouds continue to move.

3.3.6 Soft X-ray Weak AGN and the X-ray Background All three of the SXW AGN studied above, and indeed all of the SXW AGN of BLW, were originally selected by their blue optical color, and BLW found no evidence for a systematic difference between the optical continuum slopes of SXW and non-SXW AGN. Thus, given our X-ray spectral fitting results, it appears that a non-negligible fraction of blue type 1 AGN suffer significant X-ray absorption and have hard X-ray spectra qualitatively similar to that of the X-ray background. Such objects need to be remembered when constructing AGN synthesis models for the X-ray background, and indeed some blue QSOs with hard X-ray spectra have been found in X-ray background surveys (e.g., Fiore et al. 1999; Brandt 2000).

Acknowledgments: We thank F. Hamann for thoughtful and constructive comments. We acknowledge the support of a NASA Graduate Student Researchers Program Grant and the Pennsylvania Space Grant Consortium (SCG), NASA LTSA grant NAG5-8107 and NASA grant NAG5-7256 (WNB), and NASA LTSA grant NAG5-3431 (BJW). This research has made use of data obtained through the NASA Extragalactic Database (NED) and the High-Energy Astrophysics Science Archive Research Center (HEASARC) Online Service, provided by NASA’s Goddard Space Flight Center. 48

Fig. 3.8 Spectral energy distributions for each SXW AGN. Solid squares are infrared and optical continuum data from Neugebauer et al. (1987; N87), and thin solid curves are UV spectra from IUE or HST as indicated. ROSAT data points are marked with stars for PG 1011–040 and PG 1535+547, and the ROSAT best-fitting model for PG 2112+059 is indicated with a dotted line. ASCA data are for the best-fitting SIS0 model; the thick gray line is the same model corrected for intrinsic absorption for PG 1535+547 and PG 2112+059. The vertical gray bar indicates the spread of expected 2 keV fluxes from the 3000 A˚ flux densities of N87 based on the typical range of α , 1.7 <α < 1.3; ox − ox − the dashed vertical line for PG 2112+059 is derived from the HST data. All data have been corrected for Galactic reddening and absorption as described in the caption to Figure 3.1. See 3.3 for further discussion. § 49

Fig. 3.9 Plot of C iv absorption EW (in the rest frame) versus αox for the BG92 AGN (adapted from BLW). Solid dots are radio-quiet AGN and half-filled triangles are radio- loud AGN. SXW AGN (α 2) have been labeled with the right ascension parts ox ≤ − of their names. Large, open circles are SXW AGN with published 2–10 keV spectral studies. BAL QSOs are designated with boxes after their names. PG 0844+349 is also labeled with the hatched circles representing the measured αox values from four X-ray observations (Wang et al. 2000) other than the ROSAT PSPC one used by BLW. Note the remarkable X-ray variability of PG 0844+349, PG 1535+547, and PG 2112+059 as indicated with the dotted lines; PG 2112+059 is no longer ‘soft X-ray weak’ according to the definition of BLW. 50

Chapter 4

Markarian 231, a Prototypical Low-Ionization BAL QSO?

4.1 Introduction

The extraordinary galaxy Markarian 231 was discovered in 1969 as part of a survey searching for galaxies with strong ultraviolet continua (Markarian 1969). With the first spectrum, it was identified as unique among the Markarian sample because of deep, blue-shifted Na D absorption evocative of broad absorption line quasi-stellar objects (BAL QSOs; Arakelian et al. 1971). Soon after, Adams & Weedman (1972) also identified broad absorption lines from Ca ii H and K and He i at the same velocity as the Na D system. In addition, they first described the long tidal tails and disturbed morphology of the galaxy. Since those early observations, Mrk 231 has maintained its reputation as an ex- ceptional object and continues to be a favorite target in all wavelength regimes. Years before the Infrared Astronomical Satellite (IRAS) was launched, Mrk 231 was known to have an infrared luminosity on a par with QSOs ( 4 1012 L ; Rieke & Low 1972, ∼ × 1975). Weedman (1973) initially proposed that this energy could be non-thermal ul- traviolet emission reradiated thermally by dust. Even after IRAS expanded the known population of ultraluminous infrared galaxies, Mrk 231 remains one of the most luminous overall and the most powerful known at z < 0.1 (e.g., Surace et al. 1998). An association of Mrk 231 with BAL QSOs was suggested from the start (Arake- lian et al. 1971), but confirmation of this status required ultraviolet spectroscopy. The first published International Ultraviolet Explorer (IUE) spectra showed no obvious C iv or Mg ii emission or absorption, but the signal-to-noise ratio of the data was low (Hutch- ings & Neff 1987). Hubble Space Telescope (HST ) ultraviolet spectra confirmed the BAL QSO identification of Mrk 231 with the discovery of broad Mg ii and Fe ii absorp- tion (Smith et al. 1995). Broad absorption lines are found in 10% of optically selected ∼ QSOs, but Mg ii BAL QSOs only comprise 10% of that population. Other spectral ∼ characteristics that Mrk 231 shares with low-ionization BAL QSOs include reddened optical and ultraviolet continua (e.g., Boksenberg et al. 1977; Hutchings & Neff 1987), strong optical Fe ii emission (e.g., Adams & Weedman 1972; Boksenberg et al. 1977), and weak [O iii] emission (e.g., Boroson & Meyers 1992). Though the emission of many ultraluminous infrared galaxies appears to be dom- inated by energetic starbursts, Mrk 231 has been repeatedly identified as an exception

1The work in this chapter is currently in press: Gallagher, S. C., Brandt, W. N., Chartas, G., Garmire, G. P. & Sambruna, R. M. (2002) X-raying the ultraluminous infrared galaxy and broad absorption line QSO, Markarian 231 with Chandra. Astrophysical Journal . 51

(e.g., Goldader et al. 1995; Krabbe et al. 1997; Rigopoulou et al. 1999), and many pieces of evidence point toward an accreting black hole as the major power source behind the enormous infrared luminosity (though see Downes & Solomon 1998 for an alternate view). 1 The optical spectrum exhibits broad Hα (FWHM = 2870 km s− ; e.g., Boroson & Mey- ers 1992) and asymmetric Hβ emission. The bright nuclear source is unresolved at radio, infrared, and optical wavelengths (e.g., Ulvestad, Wrobel, & Carilli 1999; Lai et al. 1998; Surace & Sanders 1999), and a parsec-scale radio jet with low apparent speeds has been resolved with Very Long Baseline Array (VLBA) observations (Ulvestad et al. 1999a). Variability on time scales of years has been observed in the radio emission (Ulvestad et al. 1999b) and in the Na D and He i absorption lines (Boroson et al. 1991; Kollatschny et al. 1992; Forster et al. 1995). The IRAS color, F25/F60 (the ratio of the flux density at 25 µm to that at 60 µm), is characterized as “warm” (> 0.25), F25/F60 = 0.29 (e.g., Low et al. 1988). This implies a non-thermal continuum or warm dust, more likely to result from the hard ionizing continuum of an active nucleus. In addition, the ratio of 1 infrared luminosity to mass in H2, LIR/MH2 = 225 L M − , is difficult to explain with- out a powerful QSO contributing the bulk of the infrared luminosity (Sanders, Scoville & Soifer 1991).2 In HST images, the point-like nucleus has colors inconsistent with a starburst, and M 21.6 with no correction for reddening (Surace et al. 1998). From B ∼ − the shape of the optical and ultraviolet continua and infrared line ratios, the reddening is estimated to be A 2 mag (Boksenberg et al. 1977; Cutri et al. 1984). Correcting V ∼ for reddening then places Mrk 231 at QSO luminosities (M 24) that can account B . − for the infrared energy (e.g., Surace et al. 1998). Though the primary power behind the incredible far-infrared luminosity of Mrk 231 is almost certainly an active nucleus, the galaxy is also undergoing an energetic starburst. The B R colors of optical knots are indicative of active, early-type star formation − (Hutchings & Neff 1987), and HST images of the nucleus clearly resolved blue, star- forming knots (Surace et al. 1998). Most dramatically, in the inner kiloparsec, a nuclear 1 ring of active star formation with a rate estimated to be > 100 M yr− has been stud- ied with Very Large Array and VLBA observations (Carilli et al. 1998; Taylor et al. 1999). Furthermore, Sanders et al. (1987) observed CO emission suggestive of many spi- ral galaxies worth of concentrated molecular gas that has since been mapped with high resolution by Bryant & Scoville (1996). The optical morphology of Mrk 231 is irregular, and the asymmetries in the structure of the host galaxy are likely the result of a merger. Tidal tails extending more than 3000 to the North and South of the galaxy clearly visible in deep images further support the merger scenario. The combination of starburst and luminous QSO make Mrk 231 a classic example of the transition from starburst to QSO in the paradigm outlined by Sanders et al. (1988). X-ray studies of ultraluminous infrared galaxies are essential since they provide a direct probe not only of the environment of an active nucleus but also of the end-products of stellar evolution. Mrk 231 is of additional interest as a BAL QSO. BAL QSOs are

2For reference, the total Sanders et al. (1991) sample has a median value of L /M IR H2 1 ≈ − 26 L M , and NGC 4418, the galaxy with the next highest value after Mrk 231, has LIR/MH = 2 1 89 L M − .

52 notoriously weak X-ray sources, and few X-ray spectra of them exist in the literature. Since the HEAO-1 and Einstein era, Mrk 231 has been identified as anomalously X-ray weak given the implied power of its active nucleus (Eales & Arnaud 1988; Rieke 1988). Mrk 231 was first detected in X-rays with the ROSAT Position Sensitive Proportional Counter (PSPC; Rigopoulou et al. 1996), and extended emission was observed with the High Resolution Imager (HRI; Turner 1999). ASCA observed this source twice, and spectral analysis indicated both starburst and non-thermal emission in the X-rays (Iwa- sawa 1999; Turner 1999). With the most recent 100 ks ASCA observation, Maloney & Reynolds (2000) argued that the hard X-ray spectrum is dominated by X-rays “re- flected” from neutral or nearly neutral material, and they modeled an Fe Kα emission line with EW 290 eV. In this model, the obscuration to the nucleus is so severe that no ≈ direct X-rays below 10 keV escape along the line of sight. This implies an absorption ∼24 2 column density & 10 cm− and suggests an intrinsic X-ray luminosity more typical of a luminous QSO, although the precise value is highly uncertain. In this paper, we present a Chandra ACIS-S3 observation of Mrk 231. Previous X-ray spectral analysis of this target has been confused by the overlap of the starburst and nuclear components due to the large point spread functions (PSFs) of earlier X- ray missions. With the excellent 100 spatial resolution of Chandra, the nuclear and disk components can be spatially separated for spectral analysis. At z = 0.042, 100 is 1 1 1 approximately 0.8 kpc for H0 = 75 km s− Mpc− and q0 = 2 . The low Galactic column 20 2 density, N = (1.03 0.10) 10 cm− (Elvis, Lockman & Wilkes 1989), also makes H  × Mrk 231 an excellent target for studying the soft X-rays of an active starburst galaxy.

4.2 Observations and Data Reduction

Mrk 231 was observed with the back-illuminated Chandra ACIS-S3 detector (G. P. Garmire et al., in preparation) for 39.8 ks on 2000 October 19. The data were pro- cessed with the standard Chandra X-ray Center (CXC) pipeline from which the level 1 events file was used. The data were filtered on good time intervals resulting in a “live” exposure of 39.3 ks. The standard pipeline processing introduces a 0.005 software ran- domization into the event positions, which degrades the spatial resolution by 12% ≈ (Chartas et al. 2002). This randomization is done to avoid aliasing effects noticeable in observations with exposure times less than 2 ks. As spatial resolution was crucial for ∼ our analysis and our observation was significantly longer, we removed this randomiza- tion by recalculating the event positions in sky coordinates using the CIAO 2.0 tool, acis process events. The data were filtered to keep “good” ASCA grades (grades 0, 2, 3, 4, and 6). Filtering on “status” removed 30 X-ray events from the nucleus because ∼ X-rays were inappropriately tagged as cosmic ray afterglows. Therefore, the values in the status column for cosmic ray afterglows were ignored for the final filter on “good” status. In order to evaluate the reliability of our subsequent analysis, we checked the data for evidence of pile-up. Pile-up occurs when two or more photons arrive at a pixel during a single, 3.2 s frame and alters the spectrum and PSF. In the raw and processed data, there was no indication of a read-out streak (due to X-rays from a bright source which arrived as the CCD was being read out), which would indicate a count rate sufficient 53 to result in significantly piled-up data. However, given the observed count rate of the 1 nucleus (0.039 ct s− ), this source is within the regime where pile-up at the 5% level ∼ is a possibility. At this low level, our statistical errors are likely to be greater than any effects of pile-up. Throughout this paper, wherever the analysis could be sensitive to the effects of pile-up, we carefully test our results. The complete analysis of the possible effects of pile-up, which requires detailed spectral information, is deferred to 4.5.1.6. § The light curve of the entire CCD was examined for evidence of variability. The background did flare by a factor of 2–3 for 2.2 ks near the end of the observation. For ∼ spectral analysis of diffuse emission and image analysis, this time interval was removed leaving a total of 800 counts in this region. The count rate of the nuclear source in ∼ Mrk 231 was high enough, and the source cell region small enough, that the increase in background was not significant, and so the higher background times were kept for the nuclear light curve and spectral analysis. The nuclear source cell contained a total of 1450 counts; over the course of the flare, less than 2 additional background counts ∼ over the entire ACIS-S3 band would be expected in a circular region with a radius of 200.

4.3 Image Analysis

4.3.1 Nucleus and Extended Emission Mrk 231 is one of the 212 extragalactic sources that define the International Celestial Reference Frame, and thus coordinates accurate to . 1 mas are available from Very Large Baseline Interferometry (Ma et al. 1998). The centroid of the nuclear X-ray emission from Mrk 231 is within 100 of this position; this is within the typical absolute astrometric errors of Chandra. In order to examine qualitatively the spatial structure of the X-ray emission, the events were filtered into three broad energy bands: full (0.35–8.0 keV), soft (0.35– 2.0 keV), and hard (2.0–8.0 keV). The lower bound of 0.35 keV was chosen to avoid a strong feature in the instrumental background, and above 8 keV the effective area of ∼ the High-Resolution Mirror Assembly (HRMA) drops while the particle background rises sharply. Examining the spatial distribution in these bands gives some indication of the emission mechanisms in different physical regions. Generally, thermal, starburst emission is soft with kT . 2 keV while non-thermal, nuclear emission has a harder component with emission up to and beyond the highest energies of the ACIS-S3 bandpass. To bring out subtle features, the X-ray images were each adaptively smoothed at the 2.5σ level using the algorithm of Ebeling, White & Rangarajan (2000). The detection of soft, extended emission with the ROSAT HRI (Turner 1999) is confirmed, but the superior resolution of Chandra enables additional structure within the galaxy to be discerned. Generally, the extended emission is asymmetric, and the nuclear source is point-like and off-center with respect to the extended emission (see Figure 4.1). The galaxy X-ray emission is on the same scale as the optical emission though the tidal tails are not apparent in X-rays. From a comparison of the soft and hard images (Figures 4.1b and 4.1c), the galaxy emission can be seen to be primarily soft (E < 2.0 keV), consistent with thermal, starburst emission. In the hard band where non-thermal emission would be expected to dominate, the emission is much more compact, and the nucleus becomes 54 even more apparent. At all energies, the nucleus contributes significantly to the X-ray flux; this is most clearly seen in the surface plots of the adaptively smoothed soft-band and hard-band images (Figure 4.2). Radial surface-brightness profiles were also calculated in order to examine quan- titatively the intensity distribution; the X-ray centroid of the nuclear emission was used as the origin. Out to radii of 3000, the galaxy emission in the soft band is above ∼ the background. However, outside of 1800, the hard-band emission is consistent with ∼ the background level (see Figure 4.3a). All of the extended hard-band emission can be attributed to the wings of the PSF of the central source (see 4.5.1.6 for the complete § analysis). The spatial structure of the X-ray emission is not obviously correlated with optical or radio structures. In particular, the bright knots of starburst emission are not apparent in the X-rays (Figure 4.4). However, the X-ray emission enhancement to the Northeast (see the X-ray contours in Figure 4.4) of the nucleus is coincident with the contours of the G B (green) color from Hutchings & Neff (1987). These authors speculated that − the green emission is likely from nebular [O iii]; the structure is reminiscent of ionization cones seen in Seyfert 2 galaxies. To compare different X-ray emission mechanisms with the multi-wavelength data, we used the adaptive binning code, AdaptiveBin (Sanders & Fabian 2001), to create a hardness-ratio map. Such a map indicates the value of the hardness ratio, which is defined as the ratio of 2.0–8.0 keV to 0.35–2.0 keV counts, at any given location. The algorithm of Sanders & Fabian (2001) calculates the hardness ratio in each pixel, and then it combines adjacent pixels until the errors (propagated from Poisson statistics) on the ratio have dropped below a given fractional tolerance. Given our small numbers of counts, we chose a moderate tolerance of 0.40; i.e., the hardness- ratio error is less than 40% in each combination of pixels. The background-subtracted, hardness-ratio map overlaid with full-band X-ray contours is displayed in Figure 4.5. The nuclear region spans a noticeably higher range of hardness ratios, HR 0.6–1.8, nuc ≈ than the extended galaxy (HR 0.1–0.4) and shows obvious small-scale structure. gal ≈ With the larger number of counts in this region, finer structure can be distinguished as the pixels are smaller. An inner region directly to the West of the nucleus also shows hard X-ray emission (HR = 1.25 0.4) which corresponds to the reddest B R contours  − from Hutchings & Neff (1987). This relationship suggests that the hardness is likely the result of intrinsic absorption that would also redden the optical color.

4.3.2 Non-Nuclear Point Sources The CIAO 2.0 tool, wavdetect, which implements a wavelet source-detection al- gorithm, was used to search the full, soft, and hard-band images. Several point sources within 30 of the X-ray centroid of Mrk 231 were detected, but none (other than the nucleus) was found within the optical extent of the galaxy. However, visual inspection of the data suggested the presence of a point source within 300 of the nucleus (see Fig- ure 4.6). Based on our experience with the algorithm, wavdetect is not very sensitive to faint point sources on surface-brightness gradients such as those found near bright point sources. Therefore, the lack of a wavdetect detection of this putative point source is not unreasonable. 55

Within a 100-radius source cell of the possible point source, 23 full-band counts were registered where 10 are expected from the local background. This background ∼ was calculated as the mean number of counts per pixel in an annulus with an inner radius of 100 and an outer radius of 200 normalized to the area of the source cell. There is only a 0.03% chance of this occurrence due to Poisson fluctuations. For maximum sensitivity, the entire observation was used for this measurement; with such a small source cell the effect of the background flare should be negligible. The errors in the source and background measurements were determined using the Poisson values for 90% confidence tabulated in Gehrels (1986); they were then added in quadrature to calculate the error +6.6 in the background-subtracted source counts: 13 4.9 counts for a full-band count rate +1.7 4 1 − of (3.3 1.3) 10− ct s− . Assuming this source has a power-law slope typical of an − × +1.1 15 2 1 X-ray binary, Γ = 1.7, the 0.35–8.0 keV flux, F , equals (2.2 ) 10− erg cm− s− , X 0.8 × though this value is not sensitive to the assumed model. At the− redshift of Mrk 231, this +3.4 39 1 corresponds to a luminosity, L = (7.3 ) 10 erg s− . X 2.8 × If this is a single point source radiating− isotropically, then the Eddington limit at this luminosity suggests a massive compact object, M = 35–90 M . Such a source is consistent with observations of other starburst galaxies where several point sources with comparable luminosities and implied masses have been detected (e.g., Makishima et al. 2000). Alternatively, an off-nuclear point source with this flux could be a luminous supernova remnant (e.g., Schlegel 1995) such as was found in the Circinus Galaxy (Bauer et al. 2001). Given the low numbers of counts, this could also be a spatially extended area of enhanced X-ray emission such as a superbubble.

4.4 Variability Analysis

Nuclear light curves were examined for variability, and the full-band count rate did apparently show a subtle decrease over the course of the observation. To exam- ine this possible variability quantitatively, the full-band event arrival times from a 200- radius source cell were compared with a constant count rate using the non-parametric Kolmogorov-Smirnov (K-S) test. The K-S statistic of 0.036 gave a probability of 0.042 for the data being consistent with a constant count rate. In order to investigate possible variability further, the data were divided into soft and hard bands and tested separately with the K-S test. This strategy was motivated by the different emission mechanisms contributing in different bands; the soft-band emission may have a significant starburst component while the hard-band emission should be attributed to the active nucleus. The soft-band light curve was consistent with constant flux, and a K-S statistic of 0.033 in- dicated a probability of 0.353 that the data were consistent with the null hypothesis. In contrast, the hard-band light curve yielded a K-S statistic of 0.069 giving a probability of only 0.002 that the data were consistent with a constant count rate; this result is highly significant. The binned light curves are presented in Figure 4.7, and the bin size was chosen to obtain at least 50 counts in each bin. The hard-band bin with the lowest flux contains 55 10% of the counts in the bin with the highest flux corresponding to a  flux decrease of 45 14% over 20 ks.  To investigate further the energy dependence of the variability, we applied a K-S test to the photon arrival times in different energy bands, starting with 0.35–8.0 keV. At 56 each step, the lower limit to the energy band was increased, and the K-S test reapplied. The results of this analysis are presented in Figure 4.8; the energy band showing the most 4 highly significant variability (with a K-S probability of 4.2 10− ) was 2.75–8.0 keV. This × energy-dependent variability can be understood either as energy bands with softer lower limits suffering from dilution by a soft, constant flux component, or as a lack of soft flux from the changing component. For example, the starburst emission, which contributes primarily to the 0.5–2.0 keV emission even within the nuclear source cell (see 4.5 for § further discussion), is generated on physical scales of parsecs to kiloparsecs. Because of the larger scale of this emission, the flux from the starburst would not be expected to vary on time scales as short as the exposure time of this observation. In contrast, the small-scale, accretion-generated X-rays dominate the power of the nuclear emission at energies & 2.0 keV. Variability on the time scale of this observation is thus much more likely to arise from this physically compact region. In the hardest bands, the number of counts decreases rapidly as a result of the smaller effective area of the telescope, thus reducing the sensitivity of the K-S test to variability. In order to confirm that low-level pile-up was not causing false variability, the K-S test was also applied to an annular region with an inner radius of 0.005 and outer radius of 300. Since possible pile-up at the 5% level will only affect the innermost pixels of a ∼ point source, an annular source region is much less susceptible to pile-up, though it also contains less than half the counts in each band. As with the circular source-cell region, the hard band showed significant variability with a K-S probability of 0.013 of a constant count rate. The soft-band light curve was consistent with a constant count rate with a K-S probability of 0.892.

4.5 Spectral Analysis

The 100 spatial resolution of Chandra allowed us to perform spectroscopy on phys- ically distinct regions. Throughout the spectral analysis, Galactic absorption was fixed 20 2 at 1.03 10 cm− (see 4.1). All errors are for 90% confidence with all fitted param- × § eters taken to be of interest other than absolute normalization, unless otherwise noted. The counts, fluxes, and luminosities of different regions within the galaxy are listed in Table 4.1, and the spectral fitting results described within this section are summarized in Table 4.2.

4.5.1 Nucleus 4.5.1.1 Spectrum Extraction Analysis of ground-calibration data from the ACIS back-illuminated (BI) CCDs (which include S3) showed well before launch that these devices exhibit significant charge transfer inefficiency (CTI). CTI occurs during both parallel and serial transfers, but it is mild enough that the energy resolution is not position-dependent. However, the gain, quantum efficiency, and grade distribution change subtly across the device. In an attempt to mitigate the effects of CTI, Townsley et al. (2001b) have developed a software tool for correcting CTI through a forward-modeling algorithm incorporating a 57 d ) from 8 1 − source − · · 2 · s · · · · · · /L the / / / 2 erg 2 3 . . − of 12 5 . 6 7 40 . 17.1/153 1 2 0 L (10 determined = t. 0 areas q are ) onen 1 the − and s c to 16 , 8 · · · 2 1 · · · +13 − − comp − · · · 2 − Errors c t 45 / / / 3 3 6 4 3 0 / cm /F ...... 1 2 2 0 0 1 . . Mp 4.2. 1 − +0 +0 +0 − − − 5 erg − +1 1 8 1 5 . . . . 0 4 − . 2 1 3 s 14 F normalized 5 able dominan − T km een Luminosities (10 the b see 75 e of b v = and ts 0 ha dels; H mo Coun 4.5, 4 14 6 15 § Fluxes, 042,     . in 0 ts, 36 36 = 204 240 normalization eV) z kground k est-fitting Coun b defined Bac using the 4.1. absolute ts (0.5–8.0 regions, the 38 from 23 18 29 able 4.5. T  Coun    on § calculated in 510 323 833 1447 kground are limits Source calculated bac defined are the are a in Region confidence ts Region luminosities fluxes Emission Region 90% Coun The Regions The Inner Outer a b c d Diffuse Nucleus the regions. 58 detailed physical CCD model (Townsley et al. 2001a). This model has been tuned using an extensive library of on-orbit calibration data and tested on astrophysical sources. CTI-corrected data require redistribution matrix files (rmfs) and quantum efficiency uniformity (qeu) files distinct from those distributed by the CXC; these have been made publicly available.3 CTI correcting ACIS-S3 data allows a single rmf to be used across the whole CCD thus making analysis of extended sources more tractable. In addition, recovering the initial event grades makes grade filtering uniform across the CCD, and accounting for the charge lost due to transfers improves the energy resolution of the X-ray spectrum. After correcting the data for CTI, the spectrum of the nuclear region was extracted in pulse-invariant (PI) format from a 200-radius circular cell. The background spectrum was extracted from an annulus centered on the nucleus with an inner radius of 2.005 and an outer radius of 400. The qeu file for the CTI-corrected data was used to generate an ancillary response file (arf) using the CIAO 2.1 tool mkarf which combines the effective area of the mirrors with the qeu and quantum efficiency file. Using the FTOOL grppha, the data were grouped into bins with at least 15 counts for χ2 fitting. The data were fit with the X-ray spectral analysis tool XSPEC (Arnaud 1996). As the quantum efficiency calibration below 0.5 keV is uncertain, the counts below this value were ignored. The data above 8.0 keV were also ignored for the reasons described in 4.3.1. § 4.5.1.2 Basic Spectral Modeling The data were first fit with a power-law model. Though the fit was statistically acceptable (χ2/ν = 88.2/81), the residuals showed systematic effects. In particular, there were positive residuals below 1.0 keV and above 5 keV. Also, the best-fitting photon index, Γ = 0.46 0.09, was unusually hard; the typical photon index for a radio-quiet  QSO is in the range, Γ = 1.7–2.3 (e.g., Reeves & Turner 2000; George et al. 2000). As Mrk 231 hosts a starburst ring within 100 of its nucleus (Bryant & Scoville 1996; Carilli et al. 1998; Taylor et al. 1999), a plausible additional component to add is a Raymond- Smith thermal plasma (Raymond & Smith 1977). The plasma abundance was initially fixed at solar, a reasonable value for an energetic starburst hosting a luminous QSO. The temperature was constrained to lie within the bounds found for local starbursts by Ptak et al. (1999), kT = 0.2–3.0 keV. Compared to the simple power law, this model was preferred at the > 95% confidence level according to the F -test (∆χ2 = 7.8 − for two additional fit parameters). For the thermal plasma, the temperature, kT = +>1.9 1.1 0.2 keV, although poorly constrained, is consistent with that previously found by − +0.11 Iwasawa (1999). However, the photon index remained atypically flat, Γ = 0.31 0.20. Allowing the abundances to vary did not substantially improve the fit (∆χ2 = 1.1),− and +0.63 the best-fitting abundances were poorly constrained and low, Z = 0.32 0.32 Z . As the model fitting was not very sensitive to this parameter and the best-fitting− values appear to be unreliable, the abundances were fixed to the solar values for the remainder of the analysis.

3http://www.astro.psu.edu/users/townsley/cti/ 59

Intrinsic absorption can often masquerade as a flattening of the power-law photon index, and so intrinsic absorption was added to the power-law component. This did not improve the fit (χ2/ν = 80.5/78), and the best-fitting intrinsic column density was negligible. The photon index remained at an unusually flat value, Γ 0.3. To push ≈ this physical picture further, the photon index was constrained to lie in a reasonable range for radio-quiet QSOs, Γ = 1.5–2.5. The model fit for Γ “pegged” at the lower bound of 1.5 with χ2/ν = 121.2/78. Two additional absorption models were also tested, partial-covering and ionized absorption (Done et al. 1992; Magdziarz & Zdziarski 1995), with the photon index constrained as for the previous example. In both cases, the fits were statistically worse than the simple power law plus thermal plasma, χ2/ν = 92.5/77 and χ2/ν = 103.5/77, respectively, and the photon index value pegged at Γ = 1.5. In addition, the models underfit the flux at energies & 5 keV.

4.5.1.3 Reflection-Dominated Spectrum and Iron Line Following the results of Maloney & Reynolds (2000), we next tried fitting a ther- mal plasma plus a power law reflected from neutral material. Reflected spectra are commonly seen in Seyfert 1 galaxies in addition to the primary power-law component; the latter contributes the majority of the flux in the Chandra band. The reflected com- ponent is physically understood to arise from the direct X-ray continuum illuminating a Compton-thick accretion disk or molecular torus. In Compton-thick Seyfert 2 galaxies, which are viewed along a line of sight through the putative torus, the primary continuum is completely absorbed. Therefore, reflection dominates the emission. The XSPEC reflection model, pexrav (Magdziarz & Zdziarski 1995), was used to fit the observed data. Input parameters to this model include the inclination (i) and abundances of the reflecting medium, as well as the photon index and cutoff energy (Ec) of the illuminating power-law spectrum. The photon index was allowed to vary, but the other parameters were set to the default values in order to reduce the number of free parameters: cos(i) = 0.45, Z = Z , and Ec = 100 keV. The model was statistically acceptable, χ2/ν = 83.7/78. The best-fitting value and errors for the reflection scal- +>10 000 ing factor, R = 5737 2637 , suggest that reflection is dominating the spectrum with essentially no direct con− tinuum reaching the observer. Setting R to produce a pure re- flection spectrum gives consistent values for the photon index and the Raymond-Smith temperature. +0.28 The best-fitting photon index, Γ = 2.52 0.30, though closer to the typical val- ues for radio-quiet QSOs than those given by the− fits to a simple power-law model, is unusually high. Based on the ASCA correlation for radio-quiet QSOs between Γ and the FWHM , a photon index of 2.1 for Mrk 231 would be more typical (Reeves Hβ ∼ & Turner 2000). In addition to the unusually steep photon index, the lack of strong Fe Kα line emission accompanying a reflected continuum is also striking. Adding an unresolved Fe Kα line with E = 6.4 keV and σ = 0.01 keV to the best-fitting model only decreased χ2 by 0.2. Maloney & Reynolds (2000) measured the Fe Kα equivalent +190 width to be EW = 290 170 eV; however, we do not detect an Fe Kα line at 6.4 keV with high statistical significance.− Our 90% confidence upper limit on the equivalent width (EW < 188 eV), within the errors of the Maloney & Reynolds (2000) value, is much 60 lower than would be expected with a reflection-dominated model (EW 1–2 keV; e.g., ∼ Matt et al. 1996). Our detection of significant variability above 2 keV is inconsistent with the pro- posed model of Maloney & Reynolds (2000) for the nucleus where the reflecting medium is the torus, and any scattering medium is on even larger size scales (see their Figure 3). Both of these structures, capable of redirecting X-rays from the nucleus into the line of sight, are on physical scales too large for the observed short-term variability. In any case the reflected component would be unlikely to have energy-dependent variability. No ab- sorption of the reflected component is required; therefore, the reflected component would be unlikely to only exhibit variability above 2 keV as it dominates the flux below 2 keV as well. Since the reflected component is an unlikely source of the observed variability, adding a third component, in addition to the Raymond-Smith plasma and the reflection- dominated continuum, is thus justified. For the spectrum of this third component, we chose an absorbed power law; adding absorption is motivated by the requirement that this component not dominate the flux at the lowest energies for consistency with the observed hard-band variability. For this fit, the photon indices for both the power-law component and the incident spectrum on the reflecting medium were fixed at Γ = 2.1 according to the correlation of Reeves & Turner (2000). Given the quality of the data, this was necessary in order to obtain meaningful constraints on the other model parameters. The free parameters of interest were then the Raymond-Smith plasma temperature, the intrinsic absorption column density of the power-law continuum, and the reflection scaling factor of the reflection component. This fit gave a value of χ2/ν = 69.0/77, a significant improvement at the > 95% confidence level over all of the previous models. The Raymond-Smith +1.1 temperature, kT = 1.1 0.1 keV, was consistent throughout all of the model fits; the bump in the spectrum at− 1 keV corresponds to strong Fe L shell emission expected ∼ +427 from a plasma of this temperature. The reflection scaling factor, R = 697 223, remained large. This factor is very sensitive to the data below 1 keV; a higher normalization− for the Raymond-Smith plasma component or an additional starburst component at lower temperature would push the reflection scaling factor to higher values corresponding to a completely reflection-dominated spectrum. The value for the best-fitting column density +1.3 22 2 of the power-law component, NH = (2.1 0.9) 10 cm− , effectively eliminates flux from − × the power-law component at energies . 2 keV. In summary, we find that a three-component model provides a statistically ac- ceptable fit to the observed nuclear spectrum. This complex model is comprised of a Raymond-Smith plasma with kT 1.1 keV, an unabsorbed reflection-dominated con- ≈ +2.5 22 2 tinuum, plus an absorbed power law with N = (2.9 ) 10 cm− . Though this H 1.0 × composite model provided a good match to the observ−ed spectrum (see Figure 4.9a), it has three significant physical problems. First, the strength of the absorbed power law in this model is not sufficient to dilute the EW of the strong Fe Kα line expected from the reflected component to the observed upper limit, EW < 188 eV. In this model, the absorbed power law contributes 20% of the observed flux at 6.4 keV. Manually ∼ increasing the normalization of the absorbed power-law component to provide the nec- essary dilution near 6.4 keV results in a continuum spectrum inconsistent with these data. Secondly, the absorbed power law does not provide enough flux above 2 keV to 61 explain the observed amplitude of hard-band variability (see Figure 4.9b). Finally, the variability time scale implies a physical size of 1015 cm for the emitting medium, much ∼ smaller than any plausible reflecting medium such as the BAL wind or a molecular torus.

4.5.1.4 Alternate Model Mrk 231 exhibits complex frequency-dependent optical and ultraviolet polariza- tion that has been attributed to scattering by dusty clouds (e.g., Goodrich & Miller 1994). The nuclear continuum polarization rises with decreasing wavelength and may approach values as high as 20% in the near ultraviolet (Smith et al. 1995). The absorption and emission-line polarization is highly structured, and Smith et al. (1995) used three sepa- rate scattering media to explain the behavior of the polarization position angle. Given this evidence that scattering is an important mechanism contributing to the optical and ultraviolet emission, we considered that a similar situation may be present in the X-rays. Following this reasoning, we examined the possibility that the flat X-ray spectrum is the sum of multiple scattered power-law components with different absorption column den- sities along each line of sight. This combination could reproduce the unusually flat X-ray spectrum without requiring a large contribution from a reflection-dominated continuum. Assuming Γ = 2.1 as in 4.5.1.3, a minimum of three separate power laws, with 21.1§ 22.5 23.8 2 column densities of NH= 10 , 10 , and 10 cm− , are required for an acceptable fit to the data (∆χ2/ν = 66.0/75). The best-fitting temperature for the Raymond-Smith plasma component is consistent to within errors with the previous models. In this sce- nario, the second power law dominates the flux between 1.5–5 keV (see Figure 4.9d), ∼ the energy regime contributing most significantly to the observed variability. This model has a slightly lower value of reduced χ2 compared to the reflection-dominated plus ab- sorbed power-law model, but the improvement is not statistically conclusive. Though speculative, this model can provide an explanation for both the hard-band variability and the lack of a strong Fe Kα emission line if the scattering medium is highly ionized. This model will be discussed and developed further in 4.6.1.2. § 4.5.1.5 A Jet Contribution to the X-ray Flux? For completeness, we have also considered other possible contributions to the X- ray flux. As mentioned in 4.1, a parsec-scale radio jet was detected in Mrk 231 (Ulvestad § et al. 1999a) with apparent speed, β 0.14. The radio flux of the mas components app ∼ was found to be variable with a maximum flux density of 60 mJy (Ulvestad et al. 1999b). We investigated the possibility that this small-scale radio jet could contribute to the X- ray flux in the Chandra spectrum by estimating the synchrotron-self Compton (SSC) flux expected from the total radio flux density (0.1 Jy) of the nucleus. From the jet- to-counterjet ratio (> 45), Ulvestad et al. (1999a) estimate β & 0.48 and θ . 10◦ if relativistic boosting is responsible for the one-sided radio structure of the source; though these authors favor free-free absorption as the cause of the one-sided structure, the relativistic boosting scenario would generate the largest jet contribution to the X-ray flux. Assuming β = 0.48 and θ = 10◦, the estimated SSC flux from 0.5–8.0 keV (using 22 2 1 eq. 1 from Ghisellini et al. 1993) is 10− erg cm− s− , many orders of magnitude ∼ 62 lower than the observed Chandra flux (see Table 4.1). Therefore, SSC emission cannot explain observed X-ray properties. Another possible source of X-rays from the jet is synchrotron emission. Applying the X-ray-to-radio flux density ratios, F1 keV/F15 GHz, for the knots in the jet of M87 (Biretta et al. 1991; Marshall et al. 2001), we obtained the predicted range of X-ray fluxes 5 2 1 1 for Mrk 231, F1 keV =(0.3–9.4) 10− photon cm− s− keV− . Though this encompasses × 5 2 1 1 the X-ray flux density that we observe, F 2 10− photon cm− s− keV− , the 1 keV ≈ × steep slope, Γ 2.5 (Marshall et al. 2001), observed for the knots in M87 is inconsistent ∼ with these data (Γ 0.5). Therefore, X-ray emission from the jet is unlikely to contribute ∼ significantly to the observed flux in X-rays.

4.5.1.6 Spectral Checks The HRMA preferentially scatters higher energy photons to larger radii. Since the spectrum of the nucleus is so hard, our analysis was also performed on a spectrum extracted from a 300 radius aperture with background from an annulus with an inner radius of 3.005 and an outer radius of 500. This source aperture captures more than 90% of the encircled energy at 6.4 keV. The results from fitting the nuclear emission were consistent with those above, though the soft, thermal emission was not as well constrained. The larger aperture includes more of the extended starburst emission, and so this result is not surprising as the thermal starburst emission from the larger region is unlikely to have a single temperature. To investigate further the effects of pile-up on our results, the nuclear spectrum was also fit with the simulator-based fitting tool LYNX (Chartas 2000a). LYNX utilizes XSPEC (Arnaud 1996) to generate the incident spectrum, the PSF simulator MARX (Wise et al. 1997) to simulate the mirror response, and the Penn State ACIS simulator (Townsley et al. 2001a) to propagate X-rays through the detector. The χ2 statistic between the simulated and observed spectra is minimized to obtain the best-fitting input model parameters. For this analysis, we used model 4 of Table 4.2, but similar results would have been obtained using other statistically acceptable models. To simulate a spectrum without pile-up, we propagated the best-fitting incident model through LYNX allowing only one incident photon per frame. From this analysis, the average pile-up fraction for the observation is estimated to be 7.7%. For a simple power-law fit to our data, pile-up has the effect of flattening the slope slightly: an incident spectrum which would be fit with a photon index of Γ = 0.4 0.07 without pile-up would be measured to  have Γ = 0.3 0.08 with pile-up for our observed count rate. This level of pile-up would  cause 4% error in the flux measurements. This systematic error is much less than the ∼ statistical errors, and so pile-up has not significantly affected our spectral analysis. The final LYNX output file can also be used to estimate the extent to which the wings of the PSF of the nucleus would contaminate the spectrum of the extended galaxy emission (see 4.5.2). The simulated radial surface brightness profile in the 2–8 keV band § is shown in Figure 4.3b; the nucleus is expected to contribute approximately 6 counts in the soft band and 16 counts in the hard band to the annulus between 300 and 2500. These contributions are negligible. 63

4.5.2 Extended Galaxy Emission

From the images presented in 4.3, emission extending 2500 from the nucleus § ∼ is clearly evident. For the initial spectral fitting of this emission, a spectrum for the entire galaxy was extracted using a circular source cell with a 2500 radius. The nucleus was excised with a 300 radius circle, and the background spectrum was extracted from a region with an inner radius of 4000 and an outer radius of 6500. The rmf and arf were generated as described in 4.5.1. § The extended emission is soft, and the majority of the galaxy is not detected above 2 keV. Therefore, the data were filtered from 0.5–2.0 keV to maintain an ∼ adequate signal-to-noise ratio for the spectral fitting. This also eliminates any significant contamination from the nucleus to the spectrum of the extended galaxy (see 4.5.1.6). § As a first model, a Raymond-Smith plasma with solar abundances was fit. The resulting temperature, kT 0.4 keV, was reasonable, but the value of χ2/ν = 143.7/31 was ≈ clearly unacceptable. Substantial positive residuals between 1.0–2.0 keV indicated that an additional component was required, and so a second Raymond-Smith component was added. The new fit was statistically acceptable (χ2/ν = 37.2/29). The two temperatures, +0.07 +0.22 kT = 0.30 0.05 keV and kT = 1.07 0.18 keV, were also within the range typically seen in starburst− galaxies (0.2–3.0 keV; −e.g., Ptak et al. 1999). The spectrum of the galaxy fit with this model is presented in Figure 4.10. Some additional systematic positive residuals above 1 keV suggested that a power- law component could also fit the data, and so a power law was substituted for the second Raymond-Smith plasma component. The temperature of the Raymond-Smith plasma +0.07 increased relative to the cooler one in previous fits, kT = 0.80 0.11 keV. The best-fitting photon index for the power law, Γ = 2.6 0.7, was consistent with− emission arising from  a population of X-ray binaries, the most likely source for such a component. However, the value of χ2/ν = 40.6/29 was not a significant improvement over the previous model. To investigate whether the spectral components were spatially distinct, we divided the extended galaxy into an inner and outer region with roughly equal counts in each annulus. The inner region had an outer radius of 1000 with the nuclear region excised, and the outer region excluded the inner region. The background regions for both were the same as for the complete region. Based on the counts in the spectrum, the data from 0.5–3.0 keV were kept for the spectral analysis of the inner region. Briefly, the inner region was fit equally well with either a two-temperature Raymond-Smith plasma or a power law plus Raymond-Smith component. Both fits produced an acceptable value of χ2/ν = 14/12 (see Table 4.2). The spectrum of the outer region showed less hard X-ray emission, and so the data were filtered as for the entire extended region, 0.5 < E < 2.0 keV. The Raymond-Smith plasma plus power law was slightly preferred over the two-temperature Raymond-Smith plasma (∆χ2 = 1.6 for the same number of degrees of freedom). 64

Fig. 4.1 X-ray images of Mrk 231 adaptively smoothed at the 2.5σ level displayed on a logarithmic scale. The images have been filtered to eliminate the background flare leaving an effective exposure of 37.1 ks. The white cross marks the precise radio coordinates; the X-ray centroids in each band are within 100 of this position. The contours indicate the 2.5σ significance contours. (a) Full-band (0.35–8.0 keV) image. (b) Soft-band (0.35– 2.0 keV) image. (c) Hard-band (2.0–8.0 keV) image. 65

Fig. 4.2 Surface plots of the adaptively smoothed images of Mrk 231 in the (a) soft-band (0.35–2.0 keV) and (b) hard-band (2.0–8.0 keV). The nuclear region clearly contributes significantly to the emission at all energies, and the comparable peak intensity in each band indicates the unusually hard nature of the point source. Note the extended emission apparent in the soft band that is not seen in the hard band. 66

Fig. 4.3 (a) Radial surface-brightness profiles of the emission from Mrk 231. Each data point represents the counts per pixel in an annulus with the average radius marked; the annulus width is indicated by the horizontal error bars. The bin size was chosen so each bin contained at least 50 counts, and the errors in the vertical direction are comparable to the size of the symbols. The filled circles indicate the surface brightness in the full-band image while open pentagons and open triangles show the surface brightness in the soft and hard bands, respectively. The thick, gray line is the simulated surface-brightness profile for the hard band based on modeling with MARX and LYNX (see 4.5.1.6). The § hard-band emission out to 1800 is consistent with the PSF. The dot-dashed, dashed, and dotted lines indicate the surface brightness of the background for the full, soft, and hard bands, respectively. (b) Background-subtracted, hardness-ratio profile as a function of radius. The nucleus is significantly harder than the soft extended emission. Vertical error bars include the errors in the hardness ratio as well as the background. Where the hard-band emission is consistent with the background, the calculated hardness ratio is given as an upper limit. The dashed line indicates the hardness ratio of the background outside of the galaxy with the dotted lines showing the errors for that measurement. 67

Fig. 4.4 (a) Archival HST WFPC2 image taken with the F439W (blue) filter (Surace et al. 1998) overlaid with the contours from the adaptively smoothed, full-band X-ray image. The optical image is displayed with a logarithmic scale, and X-ray contours represent factors of two in counts. The extended structure in the X-rays does not appear to correlate with the crescent-shaped star-forming region to the South of the nucleus evident in the blue image. (b) Narrow-band Hα image (Hamilton & Keel 1987) overlaid with the contours from the adaptively smoothed, soft-band X-ray image. The optical image is displayed with a logarithmic scale, and the contours represent factors of two in counts. Note the region of enhanced emission to the Northeast of the nucleus evident in both the full and soft-band X-ray contours. 68

Fig. 4.5 Adaptively binned X-ray hardness-ratio map showing the ratio of 2.0–8.0 keV to 0.35–2.0 keV counts displayed with a linear scale; white indicates the softest regions (see 4.3). Only the region of the galaxy displayed in this figure has sufficient counts § after background subtraction to measure the hardness ratio. The nucleus has the highest surface brightness which results in the finest resolution (or smallest pixels) in the map. The galaxy is significantly softer than the nucleus; within the errors, the region at radii 200 from the nucleus has a consistent hardness ratio, HR 0.1–0.2. Note the & disk ≈ hard region to the West of the nucleus spatially coincident with the reddest B R − contours from Hutchings & Neff (1987). The overlaid contours show the intensity from the adaptively smoothed full-band X-ray image. Contour levels represent factors of two in counts. 69

Fig. 4.6 Full-band image of the 2000 2000 region immediately surrounding the nucleus × of Mrk 231 displayed with a linear scale. The precise radio coordinates of the nucleus are marked with a white cross, and the putative non-nuclear source, possibly an ultra- luminous X-ray binary or supernova remnant, is enclosed in the concentric black circles. The background region was extracted from the annulus with an outer radius of 200 (dashed line) and an inner radius of 100 (solid line). 70

Fig. 4.7 Light curves for the nucleus of Mrk 231 extracted from a circular region with a 200 radius. Each time bin has a width of 3975 s; the background is expected to contribute . 1 count per bin. Error bars are the square root of the number of counts in each bin. The background flare region (see 4.2) is indicated with vertical dotted lines. (a) § Full-band and soft-band light curves shown as solid and dot-dashed curves, respectively. While the full-band light curve shows some indication of decreasing flux over the course of the observation, the soft-band light curve is consistent with a constant count rate. (b) Hard-band light curve (dashed line) clearly indicating a count rate decrease of 45%. ≈ Note that the vertical axis does not extend to zero. 71

Fig. 4.8 The solid curve shows the K-S probability as a function of the lower limit on the energy band to which the K-S test was applied; all energy bands extend to an upper bound of 8.0 keV. The most significant variability is indicated by the minimum value of 4 the K-S probability, 4.2 10− , for an energy band of 2.75–8.0 keV. The dotted curve × (associated with the right vertical axis) shows the number of counts used in the K-S test in each energy band. Above an energy lower limit of E 4.25 keV, the calculated ≈ significance decreases sharply as the number of counts used in the K-S test drops. 72

Fig. 4.9 (a) Observed-frame Chandra ACIS-S3 spectrum of the nucleus of Mrk 231 fit with a reflected power law, a scattered and absorbed power law, and a Raymond-Smith plasma (model 4 in Table 4.2). In both (a) and (c), plain crosses are the data points, and the solid curve represents the best-fitting model. The ordinate for the lower panel, labeled χ, shows the fit residuals in terms of σ with error bars of size one. (b) Best- fitting reflection-dominated model for the nuclear data. (c) Observed-frame Chandra ACIS-S3 spectrum of the nucleus of Mrk 231 fit with three absorbed power laws and a Raymond-Smith plasma (model 5 in Table 4.2). (d) Best-fitting model with multiple absorbed and scattered power laws for the nuclear data. 73

Fig. 4.10 Observed-frame Chandra ACIS-S spectrum of the extended galaxy emission of Mrk 231 fit with two Raymond-Smith plasmas (see Table 4.2). Plain crosses are the data points, and the solid curve represents the best-fitting model. The ordinate for the lower panel, labeled χ, shows the fit residuals in terms of σ with error bars of size one. 74 c ) ν | 2 χ 0.44 0.31 0.73 0.76 0.37 ( P /ν 2 χ a 80.4/79 83.7/78 69.0/77 66.0/75 88.2/81 231 4 0 3 67 9 7 . . . . . 12 43 09 37 3 1 9 1 10 ...... 1 32 1 1 2 6 . . . 33 0 1 223 0 0 > 2637 · > 0 > > 0 0 Mrk − − − +0 +1 +89 +427 − − − +1 +1 + − · + − + + − − 6 · 1 1 . alues . . 1 3 12 21 . . . . 34 2 1 for V 697 . 1 1 0 3 59 5737 1 ) ) ) ) Data 2 2 2 2 − − − − a cm cm cm cm eV) eV) eV) eV) · · 22 22 22 22 R R (k (k (k (k · Chandr T T T T (10 (10 (10 (10 arameter k k k k P H H H H the N N N N to 09 . 20 30 11 28 . . . . Fits e e e e e 0 0 0 1 1 1 1 1 − − +0 +0 . . . . .  Γ 2 2 2 2 2 del 31 52 . . 46 0 . 2 Mo 0 d for plasma and arameters P ws la plasma and b 4.2. er w w, w, w ymond-Smith del o plasma plasma la la la p able Ra Mo T er er er ed w w w o o o and p p p w w ymond-Smith ed la la absorb Ra er er ymond-Smith ymond-Smith w w o o Ra Ra and absorb P Reflected Three Reflected P Nucleus 1. 4. 5. 3. 2. 75 2– . c 0 ) 6 ν − | = 2 31 . 10 obtained χ T 0.77 0.14 0.07 0 0.32 0.65 ( k e < normaliza- P b range ould /ν w absolute 2 the χ 9.9/14 37.2/29 40.6/29 13.9/12 13.6/12 11.5/14 143.7/31 than greater 07 22 07 08 41 20 12 25 14 05 12 11 10 21 05 11 23 15 ...... within 0 0 0 0 0 0 0 0 0 e − − − − − − − − − +0 +0 +0 +0 +0 +0 +0 +0 +0 or lie other 0.4 alues 2 80 84 81 07 21 98 26 30 34 ...... V χ to 0 0 0 1 1 0 0 0 0 ?? § of terest see in ; eV) eV) eV) eV) eV) eV) eV) eV) eV) eV) alue 2 of v − (k (k (k (k (k (k (k (k (k (k e b T T T T T T T T T T cm ued constrained arameter k k k k k k k k k k this 4.5.1 to P § 20 tin ere 10 w see 55 59 12 66 74 36 that ...... × 0 0 1 1; − − − +0 +0 +2 . Γ 03 2 . 58 45 81 1 4.2—Con . . . = 2 2 3 parameters plasmas of correct, Γ all able at T ere w w w w la la la plasma plasma plasma fixed taking freedom. del er er er absorption w w w as of o o o mo w p p p ymond-Smith eV en Ra k and and and 0 giv index . degrees all confidence Galactic 2 b eV eV ymond-Smith ymond-Smith ymond-Smith of k < the k for del 0 90% 0 Ra Ra Ra . er if E fixed . plasma plasma plasma plasma 2 b 3 , Mo photon e y for v < < w um n ha E E 4.5.1. are eratures erature erature erature § er-la the w dels Emission, o see is temp probabilit p mo errors ν Region, o-temp o-temp o-temp Region, ymond-Smith ymond-Smith ymond-Smith ymond-Smith eV; k All The All The The Tw Tw Ra Ra Ra Tw Ra a b c d e tion. 3.0 where 1. Inner Outer Extended 2. 3. 2. 2. 1. 1. 76

Fig. 4.11 Rest-frame HST ultraviolet spectrum of Mrk 231. This 770 s FOS calibration observation was taken with the 100 aperture and the G160L grating. The data have been smoothed with a boxcar filter, and they have not be corrected for Galactic reddening.

4.6 Discussion

4.6.1 Mrk 231 as a Broad Absorption Line QSO 4.6.1.1 Spectroscopic Classification The definitive classification of Mrk 231 as a BAL QSO has been debated in the literature for almost three decades. The canonical definition of a BAL QSO, as published in Weymann et al. (1991), is based on a conservative measurement of the equivalent width of the C iv resonance absorption line system called the balnicity index (BI). All QSOs 1 with BI > 0 km s− are considered to be BAL QSOs. Though Mrk 231 is a bright optical source, the large amount of intrinsic reddening in its spectrum renders it relatively faint in the ultraviolet, and a high-quality spectrum of the C iv region is not available in the literature. A search of the HST archive revealed a Faint Object Spectrograph (FOS) calibra- tion spectrum of Mrk 231 taken on 1996 Nov 21 covering the crucial C iv region. This spectrum, displayed in Figure 4.11 and unpublished to our knowledge, clearly shows ab- sorption blue-shifted from the C iv and N v emission lines, as well as possible Al iii and 1 Si iv broad absorption. A rough measure of the balnicity index, BI & 600 km s− , is a quantitative indicator of the BAL QSO nature of this object. In addition, Mrk 231 fits within the BAL QSO class based on its X-ray properties. 77

The quantity αox, the spectral index of a power law defined by the flux densities at rest-frame 3000 A˚ and 2 keV, is a useful parameter for measuring the X-ray power of a QSO relative to its ultraviolet continuum emission. A large, negative αox indicates relatively weak soft X-ray emission; the mean value of α for radio-quiet QSOs is ox ≈ 1.48 (e.g., Laor et al. 1997) with a typical range from 1.7 to 1.3 for objects that − − − do not suffer from X-ray absorption (e.g., Brandt, Laor & Wills 2000). Weakness in soft X-rays is plausibly explained by intrinsic X-ray absorption, which strongly depresses the observed flux at low-to-moderate energies. A strong correlation of large, negative values of αox with the absorption-line equivalent width of C iv supports this hypothesis (Brandt et al. 2000). As expected, BAL QSOs populate the extreme end of this correlation: they are the weakest soft X-ray sources as well as the QSOs with the most extreme ultraviolet absorption. Recent spectroscopic observations of three AGN with large, negative αox values and strong C iv absorption, PG 1411+442, PG 1535+547, and the BAL QSO PG 2112+059, found direct evidence of intrinsic X-ray absorption (Brinkmann et al. 1999; Gallagher et al. 2001b). For comparison with the BAL QSOs of the Brandt et al. (2000) sample, we have measured the equivalent width of the C iv absorption blueward of the expected location of the emission line, EW 10 A.˚ Calculating α to complete the comparison is a ≈ ox difficult matter in this complicated object. Though the rest-frame 2 keV flux density can be determined from the best-fitting X-ray model, the ultraviolet spectrum suffers from severe extinction. A na¨ıve measurement of αox from the observed 3000 A˚ flux density yields α = 1.72; however, this can only be considered an upper bound on the value. ox − Correcting for AV = 2 mag (e.g., Boksenberg et al. 1977) using the extinction curves of Sprayberry & Foltz (1992) results in α = 2.08; this value includes both starburst and ox − scattered flux as well as the reddened nuclear continuum before the reddening correction and is therefore a reasonable lower limit to αox. Regardless of the exact value of αox, the strong C iv absorption coupled with weakness in soft X-rays places Mrk 231 at the heavily absorbed end of the Brandt et al. (2000) correlation, which is populated primarily by BAL QSOs.4

4.6.1.2 X-ray Variability and Spectral Results With this Chandra observation, we can demonstrate convincingly that the 2– 8 keV X-ray luminosity is dominated by the unresolved active nucleus in Mrk 231. The hard-band radial profile is consistent with the Chandra PSF, and the significant, hard-band variability provides compelling evidence that Chandra is probing within light hours of the central black hole. At the observed rate the flux would decrease by a factor of 2 in 7 hr. The 45 14% decrease in count rate in the hard band represents ≈  42 1 a luminosity difference, ∆L 10 erg s− ; variability of this magnitude can only X ≈ reasonably originate in the immediate vicinity of the active nucleus. In addition, the compact spatial extent and very hard spectrum of the dominant energy source preclude a significant contribution by thermal starburst emission. As an alternative, a concentrated

4Though Brandt et al. (2000) did not correct for intrinsic reddening in their ultraviolet mea- surements, several of the BAL QSOs in their sample, e.g., PG 2112+059, do not have reddened continua. Those that do, such as PG 1700+518, would only have more negative values of αox. 78 population of ultraluminous X-ray binaries cannot reasonably reproduce the variability, 42.2 1 the spectrum, or the luminosity in the 2.0–8.0 keV band, L2 8 = 10 erg s− . The measured range of α , 2.08 < α < 1.72,− suggests a source that is ox − ox − under-luminous in the X-ray band by a factor of 6–36, and absorption remains the most likely explanation for this faintness. Though the moderate signal-to-noise ratio of these data is not sufficient to convincingly constrain the multi-component spectral properties of Mrk 231, progress has been made. Assuming a single, direct, highly absorbed continuum as the sole source of the nuclear X-rays would require a bizarre and contrived absorber in order to force the model to match the observed flatness of the X-ray spectrum (see 4.5.1.2). Though a reflection-dominated model such as that examined by Maloney & § Reynolds (2000) remains feasible based on the spectral fitting alone, our detection of significant variability is inconsistent with their physical picture. The observed X-ray variability during the Chandra observation strongly suggests that continuum emitted on small scales is contributing significantly to the hard-band flux. This continuum cannot reasonably be considered to be direct as the X-ray flux is not sufficient for this luminous QSO, and so we prefer an indirect (i.e., reflected or scattered) origin for this variable component. For a physically reasonable reflecting medium, the size scale is too large to account for the short observed time scale of the variability. The smallest possible source of a reflecting medium would be the BAL wind (e.g., Elvis 2000) at distances typical of the Broad Line Region (e.g., Ogle 1998). Even if the BAL wind resides at the smallest radii modeled (R 1016 cm; e.g., Murray et al. 1995), however, the shortest expected time ∼ scale of variability for the reflected continuum would be days to weeks as opposed to the hours observed. Therefore, even assuming the smallest reasonable scale for reflecting material, a reflection-dominated continuum is not favored by our observation of hard- band variability. Our alternate model for the nuclear continuum is comprised of three scattered power-law spectra absorbed by different column densities of gas, and this model also provides a statistically acceptable fit to the data. We interpret these three power-law components as scattered rather than direct based on luminosity arguments; after cor- recting for absorption, none of them could provide the 2–10 keV X-ray luminosity of 44 1 LX & 10 erg s− expected from the infrared power (e.g., Risaliti et al. 2000). The short time scale of the variability provides a spatial scale for at least one of the scattering 15 regions of R . 10 cm. An electron mirror Thomson scattering X-rays would require 9 3 densities of n & 10 cm− on these small scales (assuming a Thomson optical depth of 0.5). Given the implied X-ray luminosity, particle density, and size, the ionization state 2 5 1 of this gas can be estimated: ξ = L /(nR ) 10 erg cm s− . X ∼ This scenario supplies a coherent explanation for the hard-band variability as well as the lack of detected Fe Kα emission. The variability above 2 keV can be explained ∼ by flux changes of a single component, a scattered power law with intrinsic absorption of 22.5 2 NH= 10 cm− . Strong Fe Kα line emission would not be expected from scattering off 5 1 of a highly ionized plasma (ξ 10 erg cm s− ) where the Fe atoms are largely stripped ∼ of electrons (Kallman & Bautista 2001). Our lack of detection of any direct component to the top of the Chandra band-pass supports the conclusion of Maloney & Reynolds (2000) that an intrinsic absorber with a 79

24 2 Compton-thick column density of & 10 cm− blocks these X-rays. Furthermore, for a reasonable geometry, the blockage of the direct emission combined with the presence of the scattered emission implies that the Compton-thick absorber is on a roughly similar scale to the scattering mirror; absorbers much larger than the mirror would see it as a point source and would thus block it entirely (see Figure 4.12 for a possible structure). Given the small scales implied for this absorber as well as the requisite column density, this absorber could be identified with the “hitchhiking gas” proposed by Murray et al. (1995), which shields the BAL wind from becoming completely ionized by extreme ul- traviolet and soft X-ray photons. The hitchhiking gas would not contribute a reflected continuum because it is so highly ionized, though it could scatter X-rays into the line of sight. Thus the hitchhiking gas could provide both the Compton-thick absorber along the direct line of sight and the small-scale scattering medium that preserves short timescale variability. For the absorbers of the multiple scattered power laws, the BAL wind is a rea- sonable candidate. As mentioned above, radiative acceleration models put the BAL 16 2 wind on much larger scales (& 10 cm− ) than our proposed scatterer, and the col- 23 2 umn densities for this wind are expected to be 10 cm− (Proga et al. 2000). This ∼ column density is similar to the X-ray column densities measured for BAL QSOs from X-ray spectroscopy (Mathur et al. 2001; Green et al. 2001; Gallagher et al. 2002a), as well as the column densities of two of the three scattered power-law components in our modeling. Multi-epoch observations of Mrk 231 are certainly required to investigate the possibility of multiple scattered components further. These individual absorbed power- law components might be expected to vary with time delays dependent on the physical locations of their scattering media; in this case, flux variations in the different energy bands < 1.5 keV, 1.5–5 keV, and > 5 keV would provide additional support for this proposed model. Based on ROSAT data, Boller, Brandt & Fink (1996) and Laor et al. (1997) found a significant correlation between the X-ray photon indices of radio-quiet AGN and the value for the FWHMHβ, where those AGN with the narrowest broad Balmer lines, Narrow-Line Seyfert 1 galaxies (NLS1s), tended toward the steep end of the dis- tribution of Γ. Subsequent work with ASCA in the 2–10 keV band has extended this finding to higher energy, although the spread of Γ in the 2–10 keV band is much smaller (e.g., Brandt, Mathur, & Elvis 1997; Reeves & Turner 2000). Mrk 231 has relatively 1 narrow Balmer emission lines, FWHMHα = 2870 km s− (Boroson & Meyers 1992) and 1 FWHMHβ = 3000 km s− (L´ıpari et al. 1993), particularly for a QSO of its luminos- ity. As a low-ionization BAL QSO, Mrk 231 shares several spectral characteristics with NLS1s such as weak [O iii] lines and strong optical Fe ii emission. This similarity sug- gests a physical link between the low-ionization BAL QSOs and the NLS1s; perhaps both populations are radiating at a higher fraction of the Eddington luminosity (e.g., Brandt & Gallagher 2000). In order to pursue this potential link, a direct measurement of the underlying photon index of a low-ionization BAL QSO would offer powerful evidence for a connection. Unfortunately, with these data and the evident complexity of the X-ray spectrum, we are not able to constrain the underlying photon index of the X-ray spec- trum. The extreme X-ray faintness of other low-ionization BAL QSOs (e.g., Gallagher et al. 1999; Green et al. 2001) may prevent a direct measurement of Γ in these objects 80 for many years. The complexity we have found in the spectrum of Mrk 231 should be remembered when interpreting X-ray data on more distant (and therefore fainter and unresolved) low-ionization BAL QSOs.

4.6.2 Mrk 231 as a Starburst Galaxy In this observation of Mrk 231, Chandra has resolved the extent and structure of an X-ray luminous starburst. Over the entire galaxy, the 0.5–2.0 keV luminosity, L0.5 2.0, 41.2 1 − equals 10 erg s− , with 30% from the inner 200 (1.6 kpc). The strong Fe L shell ∼ emission of the nuclear spectrum provides evidence for a 1 keV thermal component in ∼ addition to the non-thermal spectrum. This thermal X-ray component supports the radio evidence for star formation in the subkiloparsec gas disk (Carilli et al. 1998). In this disk, the star-formation rate claimed by Taylor et al. (1999) based on their radio observations 1 is 220 M yr− . Such an active starburst region might be expected to be more X-ray 40.7 1 luminous; M82 has an 0.5–2.0 keV luminosity of 10 erg s− with a star-formation ∼ rate an order of magnitude lower (e.g., Ptak et al. 1997). Though its total flux from star formation places Mrk 231 at the X-ray luminous end of the starburst population, it is still about an order of magnitude fainter from 0.5–2.0 keV than the most X-ray luminous starburst, NGC 3256 (Moran et al. 1999). The starburst emission certainly has more than one component, though discriminating between a two-temperature thermal plasma and a thermal plasma plus power-law model is not possible at present. The temperature of the thermal plasma within the nuclear region is 1 keV, and the data are consistent ∼ with such a component contributing at all radii in the galaxy. However, the best-fitting plasma temperature at a given radius should only be considered an emission-measure- weighted average as the actual X-ray spectrum is likely to be complex. Higher counting statistics and better low-energy calibration of ACIS-S3 below 0.5 keV would contribute significantly to understanding the distribution of hot gas in this starburst in more detail. The one possible off-nuclear source is consistent with the luminous point sources seen in other starburst galaxies. One source hardly characterizes a population, however, and further observations would allow a more definitive description of the luminous X-ray binary population of this galaxy.

4.7 Future Work

Mrk 231 was the first example of a low-ionization BAL QSO with X-ray spec- troscopy, and it may be one of very few for many years as this class populates the X-ray faintest end of the BAL QSO distribution (Green et al. 2001). From X-ray spectroscopy of the handful of BAL QSOs bright enough for such analysis, the typical BAL QSO spec- trum appears to be a normal QSO continuum absorbed by gas with column densities of 22 23.5 2 NH = 10 − cm− (e.g., Gallagher et al. 2002a, 2001b; Green et al. 2001). Mrk 231 does not fit within this paradigm as it appears to have Compton-thick obscuration with only indirect X-ray reaching the observer. If Mrk 231 is typical of the class of low- ionization BAL QSOs, their extreme faintness may result from similar spectra. Spectra such as that exhibited by Mrk 231, whether reflection-dominated or composed of several, absorbed power laws, are much fainter than a direct, absorbed continuum and have a

81

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Fig. 4.12 Diagram of a possible geometry for the nucleus of Mrk 231 (see 4.5.1.4 and § 4.6.1.2). The filled black circle represents the black hole and X-ray emitting region. The § 24 2 region marked “Hitchhiking Gas” has NH & 10 cm− , thus completely blocking direct X-rays at least up to 8 keV. Two lines of sight are labeled: (1) the direct continuum, ∼ which is completely blocked, and (2) the scattered continuum exhibiting the observed X- 22 2 ray variability. This second component is observed to be absorbed with NH 10 cm− , 16 ∼ perhaps by the BAL wind launched at distances & 10 cm. Note that the hitchhiking gas would have to be virtually dust free to enable any optical and ultraviolet continuum to reach the observer. The other two scattered lines of sight implied by the spectral fitting ( 4.5.1.4) are not shown in this figure, but they would have similar paths to (2) § (though their scattering regions could be on larger scales). fundamentally different spectral shape. Though spectroscopy of other objects may be difficult, a statistical comparison of the hardness ratios of low-ionization BAL QSOs and normal BAL QSOs can investigate this hypothesis. Furthermore, Mrk 231 is likely to be the only host galaxy of a BAL QSO imaged in X-rays for many years; BAL QSOs tend to lie at much greater distances. For comparison, IRAS 07598+6508, the next nearest known BAL QSO, has a redshift of z = 0.149. Even with relatively good quality data and the advantage of excellent spatial resolution, the X-ray data for Mrk 231 are complex; similar spectra from unresolved sources at higher redshifts may prove intractable and must be treated with great care. Additional observations of Mrk 231 with Chandra to increase the signal-to-noise ratio for both the extended galaxy and nuclear spectra would allow for study of variability on longer time scales where the amplitude of variability may well be larger and could correlate with observed changes over time in the BALs. Mrk 231 has been put forth as a possible local example of energetic starburst/QSO phenomena common in the early Universe. It would be interesting to examine the hard- ness ratio and flux evolution as a function of redshift of such a galaxy for comparison with sources found in deep X-ray surveys. At present however, the spectrum above 8 keV is highly uncertain, thus making such analysis above z 1 speculative. The question ≈ of the true power of this QSO remains, and the far-infrared and hard X-rays are the most likely regimes for addressing this issue. There was no OSSE detection of Mrk 231 (Dermer et al. 1997); however, this limit is not very sensitive. An 80 ks observation with the Phoswich Detection System (15–300 keV) on Beppo-SAX scheduled for Cycle 82

5 should provide significant insight into the true power in the hard X-ray regime, and further information on whether the direct continuum is beginning to show through at the hard end of the ACIS bandpass.

Acknowledgments: We thank Dan Weedman for sharing his enthusiasm for and knowl- edge of this extraordinary galaxy. Leisa Townsley and Pat Broos generously provided expertise and software pertaining to ACIS-S3 in general and CTI correction in particular. We thank Bill Keel for graciously providing optical images of Mrk 231. Mike Eracleous, Mike Brotherton, Franz Bauer, and Nahum Arav also contributed advice and informa- tion. The insightful comments of an anonymous referee improved this paper. NASA grant NAS 8-38252, P.I. G. P. Garmire, supports the ACIS Instrument Team. S. C. Gallagher gratefully acknowledges support from NASA GSRP grant NGT5-50277 and from the Pennsylvania Space Grant Consortium. W. N. Brandt thanks NASA LTSA grant NAG5-8107. 83

Chapter 5

X-ray Spectroscopy of QSOs with Broad Ultraviolet Absorption Lines

5.1 Introduction

Since the first surveys with ROSAT , Broad Absorption Line (BAL) QSOs have been known to have faint soft X-ray fluxes compared to their optical fluxes (Kopko, Turnshek, & Espey 1994; Green & Mathur 1996). Given the extreme absorption evident in the ultraviolet (UV), this soft X-ray faintness was assumed to result from intrinsic absorption. Based on this model, the intrinsic column densities required to suppress the X-ray flux, assuming a normal QSO spectral energy distribution, were found to 22 2 be 5 10 cm− (Green & Mathur 1996). Due to the 2–10 keV response of its & × detectors, a subsequent ASCA survey was able to raise this lower limit for some objects 23 2 by an order of magnitude, to 5 10 cm− (Gallagher et al. 1999). In all of these & × studies, the premise of a typical underlying QSO spectral energy distribution and X- ray continuum was maintained. The strong correlation found by Brandt, Laor, & Wills (2000; hereafter BLW) between C iv absorption equivalent width (EW) and faintness in soft X-rays further supported this assumption. ASCA observations of two BAL QSOs, PHL 5200 (Mathur et al. 1995) and Mrk 231 (Iwasawa 1999; Turner 1999), provided suggestive evidence that intrisic ab- sorption was in fact to blame for X-ray faintness. However, the limited photon statistics and bandpass precluded a definitive diagnosis. The observation of PG 2112+059 with ASCA provided the first solid evidence for intrinsic X-ray absorption and a normal un- derlying X-ray continuum in a BAL QSO (Gallagher et al. 2001b). Subsequently, a few more observations of BAL QSOs with ASCA and Chandra have upheld this result (Green et al. 2001; Mathur et al. 2001; Gallagher et al. 2001a). Observations of three mini-BAL QSOs,2 PG 1411+442 (Brinkmann et al. 1999), RX J0911.4+0551 (Chartas et al. 2001), and PG 1115+080 (Gallagher et al. 2001a), were also consistent with the absorption scenario. In this paper, we have gathered the results from the 0.5–10 keV spectral analyses of all the available QSOs with broad UV absorption lines, including BAL QSOs, mini- BAL QSOs, and one Narrow-Line Seyfert 1 galaxy, with more than 200 counts. Though additional data exist in the literature, these are the only sources with enough counts for

1The work presented in this chapter has been previously published as Gallagher, S. C., Brandt, W. N., Chartas, G. & Garmire, G. P. (2002) X-ray spectroscopy of QSOs with broad ultraviolet absorption lines. Astrophysical Journal 567:37. 2 3 1 Mini-BAL QSOs contain UV absorption lines with velocity widths of 10 km s− , but they do not formally meet the BAL QSO criteria established by Weymann et ∼al. (1991). 84 spectral analysis on an individual basis. From our experience, drawing strong conclusions about the nature of an X-ray spectrum based on fewer counts can lead to interpretations that are not upheld by higher quality data.

5.2 Sample Data

Our sample is listed in Table 1, along with some relevant observational parameters. The X-ray data were taken by the ASCA and Chandra observatories, as indicated in the “Notes” column to Table 1, and Figure 5.1 shows those spectra not previously published. The analyses for the sample objects are described in the references listed in Table 1, with the exceptions of PG 1411+442 (Brinkmann et al. 1999) and PHL 5200 (Mathur et al. 2001). The spectral analyses for these two objects have been redone according to the method detailed in Gallagher et al. (2001b) to ensure consistency with the other results. For PHL 5200, corrective measures for dealing with the degradation of the ASCA CCD detectors, as outlined in Weaver, Gelbord, & Yaqoob (2000; Appendix A), were followed. Our derived photon index for the underlying X-ray power-law continuum of PHL 5200, +0.60 Γ = 2.17 0.47, was significantly flatter than that of the preferred model of Mathur et al. (2001). −

The general strategy for all of the X-ray analysis was to first try to fit the data with a simple power-law model (including Galactic absorption) and subsequently add intrinsic absorption of three sorts: neutral, partially covering neutral, and ionized. In all cases, including intrinsic absorption provided statistically acceptable models and im- proved the χ2 of the fits significantly. For the gravitational lenses, APM 08279+5255, RX J0911.4+0551, PG 1115+080, and H 1413+117, the lensing galaxies could provide some of the observed absorption. However, this contribution is unlikely to be significant. 22 2 At the redshifts of the known lenses, the column density would have to be & 10 cm− , much larger than the typical NH through the outer regions of a galaxy, to affect strongly the best-fitting column density. In general, distinguishing between a partially covering and ionized absorber is not possible statistically; with CCD resolution and the current statistics, the two models can fit the data equally well. Typically, the ionized-absorber models produce column densities within a factor of 0.5–2.0 of the partial-covering mod- els. However, though an ionized absorber was included for completeness in fitting each of the spectra, the ionized-absorber models, such as absori in the spectral fitting package XSPEC (Arnaud 1996), assume a single-zone, optically thin plasma. Such assumptions 22 2 are not in general appropriate for the large X-ray column densitities, NH& 10 cm− , exhibited by QSOs with broad UV absorption (H. Netzer, 2001, priv. comm.). There- fore, when a complex absorber model provided a statistically significant improvement over a simple neutral absorber, the partial-covering model was chosen. The column densities and covering fractions for the best-fitting absorption models, as well as the best-fitting photon indices, are listed in Table 1. 85 al. the C C et of e C C A Brandt analysis. 10 GL, A GL, Numerical y GL, GL, A A Notes Ellison sources 5 B, mB, B, mB, NLS1, B, B, mB, X-ra . (1993); d The the al. 6 5 9 9 10 10 · 10 EW x(corr) · of o et ˚ A) · 21 46 55 absorption. 4.4 terest. 19 ( 8.5 4.6 lens. α iv Lines in C of and e 5 6 7 8 7 7 7 8 c x trinsic b o Hamann 70 52 60 74 44 66 55 44 in ...... 9 α to 1 1 1 1 1 1 1 1 vitational description x(corr) − − − − − − − − for o Absorption gra α / x o 2.26/ 1.72/ 1.74/ 2.17/ 2.01/ 2.03/ 1.62/ 1.66/ (2001a). GL: α calculate − − − − − − − − erscripts. violet detailed (1991). 1; al. to corrected sup 23 66 21 61 90 95 23 47 26 98 26 73 91 92 26 60 et al. the ...... Ultra normalization 0 0 0 0 0 0 0 0 b +0 +0 +0 +0 +0 +0 +0 +0 − − − − − − − − et Γ used for Seyfert 87 99 20 02 17 86 39 97 ...... 1 1 2 2 2 1 1 1 ectrum Broad except umerical sp n b Gallagher eymann 4 39 17 12 26 23 20 12 02 07 12 y ...... w-Line 0 0 0 0 0 with densities · · · W reference − − − − − +0 +0 +0 +0 +0 · · · 8 with ering · · · 5.2; v X-ra 71 64 97 91 79 § . . . . . raction flux Narro the 0 0 0 0 0 F Co QSOs parameters the ˚ A see 3 b of (1987); all H ) 2 indicated NLS1: N 1 2 4 5 5 5 al. 6 5 − from ...... 3000 . . 5 indicate 3 2 0 18 12 22 6 +8 − − − − +3 +2 +0 erties +28 − +18 − +26 − et are − +7 cm (2001); 0 8 1 . . . the 12 taking 19 20 40 22 19 6 3 1 QSO; trinsic al. Prop for (10 alues In column et v (C). calculated a Basic this EW z ere Neugebauer confidence 7 mini-BAL 3.870 2.800 1.722 2.548 0.090 0.038 0.457 1.981 w in Chartas references 2 5.1. Chandr 90% mB: B the 19.2 18.0 15.8 17.0 15.0 15.3 15.5 18.0 x(corr) (1997); and o for able absorption α prep.; T erscripts 1 2 al. QSO; en iv in (A) for et sup giv C 1 3 4 4 indicate al., 1 a BAL of are 3 et alues ASCA B: Bade v 6 Name 08279+5255 (2000). are 5200 J0911.4+0551 2112+059 1535+547 1411+442 1115+080 erscripts Errors Sources Numerical The Key: al. a b c d e 1413+117 Chartas 1 et (1999); sup data PG PG H APM PHL RX PG PG 86

5.3 Results

As obtaining even moderate-quality spectra of these faint X-ray sources requires a significant amount of observing time with the current generation of X-ray observatories, a comprehensive and intricate X-ray picture of the population of QSOs with broad UV absorption lines awaits the accumulation of a significantly larger body of data. However, as the number of these QSOs observed in X-rays grows, a consistent basic picture is beginning to emerge.

5.3.1 Normal Underlying QSO Continua Most fundamentally, the days of assuming, without evidence, that QSOs with broad UV absorption lines have the underlying X-ray continua of normal QSOs are over. The spectroscopic measurement of the photon index, Γ, of the power-law continuum for each of these objects versus redshift is shown in Figure 5.2a. Excluding the apparent outlier H 1413+117 (see 5.3.2.1), the mean photon index for our sample is Γ = 2.01 §  0.13, entirely consistent with the range of Γ measured for large samples of low and high luminosity radio-quiet QSOs (e.g., George et al. 2000; Reeves & Turner 2000). Furthermore, there is no evidence for strong systematic changes in spectral slope as a function of redshift. In fact, the well-measured continuum slope of APM 08279+5255 +0.26 (Γ = 1.86 0.23) extends constraints on radio-quiet QSO X-ray continuum shape evolution to significan− tly higher redshift, z = 3.87, than previously possible. For evaluating the spectral energy distribution of a QSO, the quantity αox mea- sures the relative flux densities at 2 keV and 3000 A.˚ Radio-quiet QSOs without in- trinsic absorption have α = 1.51 0.14 (BLW), while absorbed QSOs tend to have ox −  α 1.7. The α values for this sample are found to be generally indicative of ab- ox . − ox sorption (see Table 1 and Figure 5.2b), though the measured X-ray column-density is not obviously correlated with αox. Though our values for the 3000 A˚ flux density were not corrected for possible intrinsic reddening, this is not expected to be a significant problem for these blue QSOs. Furthermore, though variability between the optical and X-ray observations could result in an inaccurate αox, this quantity is only sensitive to large changes in flux density; e.g., a factor of two increase in the 3000 A˚ flux density only corresponds to ∆αox= 0.1. Correcting the X-ray spectra for the fitted intrinsic absorption brings these objects within the range of typical αox, indicating that the un- derlying spectral energy distributions are not anomalous (see Figure 5.2b). Both of the assumptions of the early BAL QSO X-ray work, normal X-ray continuum and typical αox, are thus supported.

5.3.2 Complex Absorption A simple neutral absorber blocks almost all X-ray flux up to an energy set by the 23 2 column density; for N = 10 cm− , X-rays up to 3 keV in the rest frame are absorbed. H ∼ A signature of more complex absorption, for example partial covering and/or ionized gas along the line of sight, is a much less abrupt decline of the flux at the lowest energies (see Figure 5.1b for an example). For five of the eight sources, a partial-covering absorber 87 provided a statistically significant improvement over a neutral absorber, and the column densities listed are for that model. For the remaining three QSOs, APM 08279+5255, H 1413+117, and PG 2112+059, the current data do not permit the nature of the absorber to be examined in detail. For H 1413+117 and PG 2112+059, the signal-to-noise ratio of the X-ray data below 1 keV is insufficient for studying the structure of the absorption, and for APM 08279+5255, the large redshift (z = 3.87) pushes the signatures of a complex absorber below the Chandra ACIS bandpass (0.5–8.0 keV).

5.3.2.1 H 1413+117 The famous gravitationally lensed “Cloverleaf” QSO, H 1413+117, does not fit within the basic picture outlined above for the other absorbed QSOs in this sample. Though significant intrinsic absorption is measured, the best-fitting photon index is unusually flat for a radio-quiet QSO, Γ 1.4. Furthermore, even after correcting for ≈ intrinsic absorption, α remains low at 1.74. Finally, including a narrow, 6.4 keV Fe Kα ox − emission line in the model improves the fit at the > 99% confidence level according to the F -test (∆χ2 = 9.1 for one additional fitting parameter; Figure 5.1c). In the rest +354 frame, this strong line has an equivalent width of EW = 650 517 eV. The flat rest-frame 1.8–18 keV slope and strong Fe Kα emission line together −suggest that the spectrum is reflection-dominated, and thus the direct X-rays are completely obscured. In terms of absorption-line profiles, H 1413+117 is similar to PHL 5200, with smooth profiles starting at near-zero velocities (Turnshek et al. 1988), and both QSOs also exhibit high continuum polarization (Goodrich & Miller 1995). However, the X-ray properties of these two QSOs are quite different as PHL 5200 shows no evidence for reflection. The difference between H 1413+117 and the other BAL QSOs may lie in its UV spectrum; Hazard et al. (1984) found broad Al iii absorption, which has never been found without Mg ii absorption, the signature of a low-ionization BAL QSO (LoBAL QSO; N. Arav and M. Brotherton 2001, priv. comm.). Those BAL QSOs with broad absorption in Mg ii have notably lower αox values than BAL QSOs with only high-ionization absorption lines (Gallagher et al. 1999; Green et al. 2001). This could plausibly result from either an intrinsically X-ray weak con- tinuum, or from heavy, perhaps Compton-thick, obscuration (e.g., Green et al. 2001). In the latter case, only indirect X-rays from the active nucleus, either scattered around the obscuration by a small-scale electron mirror or reflected off of more distant, neutral material, can reach the observer, and the intrinsic X-ray power of the nucleus cannot be directly measured. Given the extreme X-ray faintness of this class of objects, very few spectra are likely to be obtained. The two examples available at present are Mrk 231 (Maloney & Reynolds 2000; Gallagher et al. 2002b) and H 1413+117. For these two objects, two pieces of evidence indicate that the direct X-ray continuum is not being observed: the spectral shape is unusually flat (Γ . 1.5), and the value of αox is still low after correcting for all the direct signs of absorption. 88

5.4 Discussion

In every case where the analysis is sensitive to structure in the X-ray absorption, we find evidence for complexity. While we model this complexity as intrinsic, neutral absorption partially covering the X-ray continuum source, the actual nature of this ab- sorbing gas could be much more complicated. X-ray scattering in the nucleus effectively causes partial-covering absorption, and if the X-ray absorbing gas is physically related to that absorbing the UV continuum, then it should also be ionized. Determining the ionization state of the X-ray absorbing gas and the continuum scattering fraction is es- sential for accurately measuring the column density. Furthermore, the current modeling assumes no velocity dispersion in the absorbing gas. In such gas, the opacity in the soft X-rays arises primarily from bound-free metal edges, and these are used to determine the absorption column density. With a significant velocity dispersion, the opacity of bound-bound absorption lines can increase significantly. Therefore, the assumption of zero velocity dispersion, clearly na¨ıve for BAL QSOs, may result in an overestimate of the amount of X-ray absorbing gas. The significant correlation found by BLW of C iv absorption EW versus weakness in soft X-rays holds across several orders of magnitude of absorption EW. The inference that the weakness in soft X-rays is the result of intrinsic absorption has been substan- tiated spectroscopically for all but one of the objects in our sample (H 1413+117, see 5.3.2.1). However, directly identifying the EW of the UV absorption with the column § density of the X-ray absorbing gas is problematic (see Figure 5.2c). BLW speculated that at the absorbed end of their correlation, the EW of C iv is dominated primarily by velocity dispersion rather than column density. In this case, BAL QSOs, by definition (Weymann et al. 1991), would be expected to have the most extreme C iv EW proper- ties, as they do. As seen in Figure 5.2c though, BAL QSOs may not necessarily always have the largest X-ray absorption column densities. In terms of gross X-ray properties, it has become apparent that drawing a distinc- tion between BAL QSOs and other QSOs with broad UV absorption lines is largely a matter of semantics. The photon indices and the column densities measured with the cur- rent X-ray observatories overlap. Future, higher spectral resolution X-ray observations may reveal features warranting separate classifications, but at present the differences do not justify treating the BAL and mini-BAL QSO populations, for example, as distinct.

5.5 Future Work

Though X-ray spectroscopy of QSOs with broad UV absorption lines is finally addressing some of the questions first raised by ROSAT observations, many issues remain unresolved. Primarily, the velocity structure and ionization state of the X-ray absorbing gas are unknown. Without this crucial information, the mass ejection rate of gas from the nuclei of these QSOs cannot be determined. This parameter is fundamental for understanding the balance of accretion and outflow in the central engine. To address these important questions, X-ray spectroscopy at grating resolution of the X-ray brightest QSOs with broad UV absorption is the next step. Investigating the physical basis for 89 the extreme X-ray faintness of LoBAL QSOs may need to await the next generation of X-ray observatories.

Acknowledgments: We thank F. Hamann for providing equivalent-width data, C. Vig- nali for helpful discussions, and M. Brotherton for constructive comments that improved this paper. NASA grant NAS 8-38252, P.I. GPG, supports the ACIS instrument team. SCG gratefully acknowledges support from NASA GSRP grant NGT5-50277 and from the Pennsylvania Space Grant Consortium. WNB thanks NASA LTSA grant NAG5- 8107. 90

Fig. 5.1 Observed-frame Chandra ACIS-S3 spectra. Both (a) APM 08279+5255 and (b) PG 1115+080 have been fit above rest-frame 5 keV with a power-law model that has then been extrapolated back to lower energies. The residuals at low energies are signatures of significant intrinsic absorption. (c) H 1413+117. The continuum has been fit with a power-law model with intrinsic absorption. Significant positive residuals near rest-frame 6.4 keV are indicative of a strong neutral Fe Kα emission line with rest-frame EW 650 eV. ≈ 91

Fig. 5.2 For all plots, data points are for the QSOs listed in Table 1 unless otherwise indicated. Filled stars mark BAL QSOs, filled triangles mark mini-BAL QSOs, and the filled circle marks the NLS1, PG 1535+547. (a) Photon index versus redshift for the QSOs listed in Table 1. The filled square is the result from Green et al. (2001) for a simultaneous fit to six BAL QSOs versus their median z. The dashed line represents the mean of the George et al. (2000) ASCA sample of radio-quiet QSOs, and the dotted lines indicate the dispersion. (b) Histograms of αox values. The dotted histogram shows the distribution of the BLW sample of low-redshift Palomar-Green QSOs without known X-ray absorption. The shaded histogram shows the distribution of αox for our sample as listed in Table 1. The filled symbols show αox(corr), the αox values for our sample corrected for intrinsic X-ray absorption. (c) Intrinsic X-ray column density versus rest- frame C iv absorption EW. 92

Chapter 6

First Results from an Exploratory Chandra Survey of Large Bright Quasar Survey Broad Absorption Line QSOs

6.1 Introduction

Since the first surveys with ROSAT , Broad Absorption Line (BAL) QSOs have been known to have faint soft X-ray fluxes compared to their optical fluxes (Kopko et al. 1994; Green et al. 1995; Green & Mathur 1996) Given the extreme absorption evident in the ultraviolet, this soft X-ray faintness was assumed to result from intrinsic absorption. Due to the 2–10 keV response of its detectors, a subsequent ASCA survey was able to 23 2 raise this lower limit for some objects by an order of magnitude, to 5 10 cm− & × (Gallagher et al. 1999). In all of these studies, the premise of a typical underlying QSO spectral energy distribution and X-ray continuum was maintained. The strong corre- lation found by Brandt, Laor, & Wills (2000; hereafter BLW) between C iv absorption equivalent width (EW) and faintness in soft X-rays further supported this assumption. The observation of PG 2112+059 with ASCA provided the first solid evidence for intrinsic X-ray absorption and a normal underlying X-ray continuum in a BAL QSO (Gallagher et al. 2001b). Subsequently, a few more observations of BAL QSOs with ASCA and Chandra have upheld this result (Green et al. 2001; Mathur et al. 2001; Gallagher et al. 2001a). Gallagher et al. (2002a) compiled the results from the first moderate-sized survey of BAL QSOs with enough counts for spectroscopic analysis. They concluded from the spectroscopic evidence that the intrinsic ultraviolet-to-X-ray spectral energy distribu- tions of BAL QSOs are consistent with those of typical QSOs. Furthermore, complex, 23 2 intrinsic absorption with N = (0.1 4) 10 cm− is generally evident in the X-ray H − × spectra. Constraining the physical state of the X-ray absorber, e.g., the column density, ionization state, covering fraction, and velocity structure, is the ultimate goal of such studies. At present however, only a handful of these generally faint targets have provided data of sufficient quality for spectral analysis. This sample is generally heterogeneous, and thus far there is no outstanding predictor of observed 0.5–10.0 keV flux based on the ultraviolet spectral properties. In other words, there has been no obvious connection between the characteristics of the ultraviolet and X-ray absorbers as of yet, though all QSO with broad ultraviolet absorption also apparently have X-ray absorption. In an effort to remedy this situation, we are in the process of compiling the multi- wavelength properties of a complete sample of BAL QSOs from the Large Bright Quasar Survey (LBQS; Hewett, Foltz & Chaffee 1995). The sample will ultimately include all of the BAL QSOs from the LBQS with z > 1.35, the redshift at which the definitive C iv 93

BAL is shifted into the wavelength regime accessible with ground-based spectroscopy. One of the ultimate goals of this survey is to understand the connection between ultra- violet and X-ray absorption in luminous QSOs, and hence gain some insight into the mechanism for launching and maintaining energetic winds. The targets are optically 1 bright with BJ magnitudes from 16.7–18.5, and they are drawn from a homogeneous, magnitude-limited survey that has effective, well-defined and objectively applied selec- tion criteria (e.g., Hewett, Foltz & Chaffee 2001). The redshifts range from 1.41–2.88, and several low-ionization BAL QSOs (LoBAL QSOs), those with Mg ii absorption in addition to the more typical C iv, are included. In this chapter, we present exploratory Chandra Advanced CCD Imaging Spectrometer (ACIS; G. P. Garmire et al., in prepara- tion) observations of the first 15 targets. These short (5–7 ks) observations are intended to determine the basic X-ray fluxes and rough spectral shapes of the sample, and for those targets that were not detected, sensitive upper limits can be set with these ex- posure times. In addition to presenting new X-ray results, we also compare these data to some of the spectroscopic and polarimetric information available in the literature to search for correlations.

6.2 Observations and X-ray Data Analysis

Our observed sample is listed in Table 6.1, along with some relevant observational parameters. All of these objects were observed in the spectroscopic BAL QSO survey of Weymann et al. (1991, hereafter WMFH), and so a significant amount of data is available from the literature both from this work and others, including several broad- band polarimetric studies (Hutsem´ekers, Lamy & Remy 1998; Hutsem´ekers & Lamy 2000; Lamy & Hutsem´ekers 2000). Each target was observed at the aimpoint of the back-illuminated S3 CCD of ACIS in faint mode. The data were processed using the standard Chandra X-ray Center (CXC) aspect solution and grade filtering. The optical position of each target was obtained from the LBQS final catalog paper (Hewett et al. 1995). Using Digitized Sky Survey (DSS) images this position was refined by determining the J2000.0 coordinates of the optical centroid of the QSO image.2 From the Chandra level 2 events file, a circle of radius 200 with the optical position as the center was searched for X-ray photons in the 0.35–8.0 keV bandpass using the TARA IDL software package (Broos & Townsley 2000). TARA is a powerful, inter- active data visualization tool written in IDL that allows for live-time filtering on any column in a FITS format X-ray event file. The light curves within each aperture were checked to ensure that spurious events such as flaring pixels were not contributing to the counts. An object was considered detected if more than 3 counts were observed within this aperture. This low detection threshold is reasonable considering the extremely low

1This blue magnitude is related to the more standard Johnson blue magnitude using the following relation: BJ = B 0.28(B V ) (Hewett et al. 1995). 2 − − The Digitized Sky Surveys were produced at the Space Telescope Science Institute under U.S. Government grant NAG W-2166. The images of these surveys are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the UK Schmidt Tele- scope. The plates were processed into the present compressed digital form with the permission of these institutions. 94

Table 6.1. Basic Properties and Observation Log

c d LBQS f2500 Galactic NH Observation Exposure a b 16 20 2 Name z BJ (10− fλ) (10 cm− ) Date Time (s)

0019+0107 2.128 18.1 6.50 3.2 2001-12-08 4799 0021 0213 2.294 18.7 3.76 2.9 2001-08-20 6746 − 0025 0151 2.078 18.1 6.11 2.8 2001-10-31 4695 − 0029+0017 2.229 18.6 2.81 2.8 2001-06-20 6686 1029 0125 2.039 18.4 1.60 4.8 2001-05-30 4502 − 1212+1445 1.638 18.0 4.61 2.7 2001-05-30 4502 1216+1103 1.616 18.4 3.02 2.1 2001-03-18 4638 1231+1320L 2.380 18.8 6.94 2.6 2001-08-07 6708 1235+0857 2.880 18.0 6.13 1.7 2001-03-18 5627 1239+0955 2.020 18.2 5.46 1.6 2001-07-25 5160 1331 0108L 1.870 17.9 11.99 2.2 2001-03-18 5062 − 1442 0011 2.202 18.2 2.59 3.6 2001-05-30 4278 − 1443+0141 2.448 18.2 1.95 3.3 2001-03-23 5937 2154 2005 2.035 18.1 4.77 2.7 2001-09-04 5002 − 2201 1834L? 1.810 17.8 9.42 2.8 2001-08-03 5080 −

aThe names marked with “L” indicate low-ionization BAL QSOs with Mg ii absorp- tion. The classification of 2201 1834 is uncertain because Mg ii is not covered in the − WMFH spectrum; however, its reddened continuum and strong C iv absorption are typ- ical of LoBAL QSOs. bThe blue filter used in the LBQS survey. See 6.1. § c 2 1 The value for f2500 (erg cm− s− ) was measured from the continuum fit at rest-frame 2500 A˚ using the spectra of WMFH. dGalactic column densities from Dickey & Lockman (1990). 95

Fig. 6.1 The differences between the Digital Sky Survey and Chandra positions for the detected BAL QSOs. instrument and sky background of the Chandra ACIS-S3 combination. Given the typi- 7 1 cal observed background count rate of 4.2 10− ct pix s− , within a circular aperture × with a 100 (2-pixel) radius, 0.03 counts would be expected in a 5 ks exposure. For this 6 background rate, 3 counts in the source cell would occur with 4.4 10− probability × according to Poisson statistics. The validity of Poisson statistics for this situation has been established empirically by Vignali et al. (2001). For each detected source, the X-ray centroid within the 200 source cell was calculated with TARA. The offsets between the DSS and Chandra positions are plotted in Figure 6.1. In general, they were within 100, the absolute astrometric accuracy of Chandra. Of the 15 observed BAL QSOs, 11 of 15 were detected with counts ranging from 4–84 photons. After determining the position of the X-ray centroid, the counts in the full (0.35– 8.0 keV), soft (0.35–2.0 keV), and hard (2.0–8.0 keV) bands were extracted from a source cell with a radius of 2.005. The lower limit of 0.35 keV was chosen to avoid a large spectral feature in the instrumental background; above 8.0 keV, the effective area of the Chandra mirrors drops considerably and the background increases significantly. The hardness ratio (HR), defined as the 2.0–8.0 keV counts divided by the 0.35–2.0 keV counts, was then calculated; this parameter provides a coarse quantitative measure of the spectral shape. In order to transform the HR into a physical parameter, the photon index, for characterizing the X-ray spectrum, observations of a power-law X-ray spectrum were simulated using the X-ray spectral modeling tool, xspec (Arnaud 1996). For these simulations, the requisite redistribution matrix files and ancillary response files were created by the CIAO 2.1 tools “mkrmf” and “mkarf” for the ACIS-S3 aimpoint. The 96 input models had photon indices ranging from 1.0–3.0; Galactic absorption as listed in − Table 6.1 was also incorporated. For reference, a typical QSO with no intrinsic absorption and a photon index, Γ = 2.0, would have been observed to have HR=0.13–0.15 for the 20 2 range of Galactic column densities of (1.6–4.8) 10 cm− in this sample. × While typical QSOs have rest-frame 2–10 keV photon indices of Γ = 1.9 0.27  (e.g., Reeves & Turner 2000), a QSO exhibiting intrinsic absorption will have an ap- parently harder X-ray spectrum due to the lack of low-energy X-rays. Thus, a larger value of the HR will correspond to a lower value of Γ, which might result from intrinsic absorption. Once the value of Γ corresponding to the observed HR and Galactic NH was determined, the model power-law spectrum was normalized using the full-band count rate. Given a photon index and normalization, the derived X-ray properties, FX (2– 8 keV flux) and F2 keV (flux density at rest-frame 2 keV) could be calculated. Finally, αox, the slope of a hypothetical power law connecting the flux densities at 2500 A˚ and 2 keV, was determined. An alternative strategy would have been to assume Γ = 2.0 and use simulations to match HR to an intrinsic column density. As the signature of intrinsic 23 2 absorption is an exponential cut off, for NH& 10 cm− , the flux density at 2 keV is extremely sensitive to the exact value for NH. Therefore, allowing Γ to vary will have less effect on the calculated value of αox for the range of Γ consistent with the observed HR. Note that the 2500 A˚ flux density has not been corrected for intrinsic reddening in the QSO. Apparently unreddened QSOs exhibit a range of continuum steepness (Richards et al. 2002), and so it is not clear how to measure or correct for reddening. In any case, these LBQS QSOs are generally blue, and αox is a robust parameter. An increase in f of a factor of 2 would result in ∆α 0.1. The values for the X-ray properties are 2500 ox≈ presented in Table 6.2.

Based on the X-ray spectral analysis of Gallagher et al. (2002a), the observed X-ray continuum above rest-frame 5 keV for their spectroscopic sample of BAL QSOs was found to be largely free from the effects of intrinsic absorption. In each case, fitting this hard continuum gave Γ 2.0, consistent with the photon indices from the final ∼ model fits which included the complex absorption. If this situation generally holds true, the intrinsic 2 keV flux density can be reasonably estimated by using the X-ray continuum above 5 keV to normalize a typical QSO power-law spectrum. Following this assumption, the observed counts in the rest-frame 5 keV to observed-frame 8 keV bandpass were used to normalize a Γ = 2.0 power-law model. With this normalization, an “absorption-corrected” value for the 2 keV flux density could be used to calculate an intrinsic αox, αox(corr). The final column in Table 6.2 lists these values for the detected BAL QSOs; α (corr) ranges from 1.36 to 1.82, a reasonable span considering the ox − − statistical errors in this calculation.

6.3 Testing for Significant Correlations

From the available ultraviolet and optical data in the literature, correlations were sought that might indicate a connection between the X-ray properties and those in 97 σ er d to 1 zero · · · · eV upp · · · · are k 1.66 1.81 1.71 1.44 1.48 1.52 1.36 1.55 1.82 1.62 1.51 for · · · · (corr) x 5 − − − − − − − − − − − o the α errors (1991) 20 06 31 22 . . . . x 2.08 2.28 2.23 1.54 1.68 2.35 1.65 1.85 2.22 1.75 1.90 2 2 2 2 o The − − − − − − − − − − − − − − − α rest-frame < < < < Nousek ts). c from & non-detections, ) 1 or − coun ws s F pass 2 eV − 10 17 28 55 74 92 14 97 80 13 50 51 31 06 65 06 26 29 27 30 20 96 ...... k 1 0 1 0 0 1 1 0 0 0 1 94 44 95 68 Burro . . . . cm +1 +0 +2 +0 +1 +3 +1 +1 +1 +0 +1 − − − − − − − − − − − X 0 1 0 1 band F (1986). 28 30 69 93 05 28 37 70 31 18 75 ...... < < < < erg 3 0 2 2 2 8 3 1 0 5 3 the Kraft, 14 − in erties of ectrum. Gehrels (10 sp ts)/(0.35–2.0 rate Prop w t y tables from 45 32 49 10 34 09 30 73 13 45 72 87 95 14 56 11 44 54 34 67 b ...... coun 0 0 1 0 0 0 0 0 0 0 17 · · · · er-la . − − − − − − − − − − +0 +0 +3 +0 +0 +0 +0 +1 +0 +0 · · · · els coun 2 the w X-ra · · · · o eV 86 33 00 29 57 35 64 00 14 00 Ratio > ...... lev k p 0 0 2 0 0 0 0 1 0 1 the Hardness 0 . using 2 6.2. a ) = 1 using − Γ able s (2.0–8.0 rate. T a confidence t ct as 3 calculated eV 91 33 61 24 88 55 05 00 56 56 04 25 60 98 47 26 94 32 32 91 85 36 − ...... 90% k coun 0 0 0 1 0 0 2 1 0 0 1 51 92 59 08 . . . . − − − − − − − − − − − +1 +0 +0 +1 +1 +0 +2 +1 +0 +0 +1 (10 0 0 0 1 are calculated 9 for . 08 35 58 49 59 28 44 60 71 19 the defined ...... < < < < normalize is el) 1 14 1 6 3 0 1 2 3 2 1 (1991). are is 2.0–8.0 to Rate lev t from ons eV (corr) k here x Ly errors o Coun α of 8.0 of 3 ratio . 7 confidence . 1 propagated § frame alue v 0213 0151 0125 0108L 0011 2005 1834L? detections, (90% ed − − − − − − − Name or LBQS wing Hardness The kground. F Errors a b c d 1216+1103 1235+0857 1443+0141 0021 0025 0029+0017 1029 1212+1445 1231+1320L 1239+0955 1331 1442 2154 2201 0019+0107 limits follo observ bac 98 other parts of the spectrum. In particular, parameters which characterize the ultraviolet absorption were investigated. The first is the “BALnicity index” (BI), a conservative value of the C iv absorption equivalent width that includes any contiguous absorption 1 that falls between 3000–25,000 km s− blueshifted from the systemic redshift, if the 1 absorption (at least 10% below the continuum) exceeds 2000 km s− in width. The second absorption parameter is the “detachment index” (DI), a dimensionless measure of the velocity of the onset of absorption normalized to the width of the emission line. Both are typically measured for C iv absorption and were originally defined and tabulated by WMFH. If the extent of X-ray weakness relative to that expected for a normal QSO is due to intrinsic absorption, αox might be expected to correlate with either or both of these parameters. To test for such a relationship, we calculated three non-parametric statistics: the global χ2 of the Cox Proportional Hazard Model (Isobe et al. 1986), Kendall’s τ, and Spearman’s ρ (Press et al. 1997, Chapter 14). From the distribution of these statistics expected from unrelated data, the probability that two variables are correlated can be estimated; the results of this analysis are presented in Table 6.3.

Table 6.3. Results from Bivariate Statistical Testsa

Cox Hazard Model Generalized Kendall Spearman’s ρ Variablesb Global χ2 Prob.c τ Prob.c ρ Prob.c

BALnicity Index/αox (15) 0.121 0.723 0.847 0.397 0.168 0.529 Detachment Index/αox (15) 2.605 0.107 1.483 0.138 0.380 0.155 Polarization/α (11) 1.562 0.211 1.014 0.311 0.422 0.182 ox − αB/αox (11) 1.103 0.294 0.590 0.556 0.208 0.510 Hardness Ratio/αox (11) 0.728 0.009 2.499 0.013 0.656 0.038 Hardness Ratio/z (11) 6.780 0.009 2.499 0.013 0.774 0.014 −

aBivariate statistical tests were performed using ASURV Rev 1.2 (La Valley et al. 1992) which implements the methods described in Isobe, Feigelson & Nelson (1986). bThe number of data points is given in parentheses. cThe probability that the given variables are not correlated.

As can be seen in Figure 6.2a, there is no evident correlation between the C iv absorption-line BI and αox; this is supported by the high probability that the null hy- pothesis is consistent with these data (see Table 6.3). This result is perhaps surprising 99 given that Brandt et al. (2000) found a highly statistically significant correlation be- 3 tween C iv absorption EW and αox. However, their sample, the Bright Quasar Survey (Schmidt & Green 1983) objects with z < 0.5, encompassed a much larger dynamic range of both αox and absorption EW; the majority of their sample exhibited neither weakness in X-rays nor C iv absorption. This BAL QSO sample is probing the most ex- treme end of C iv absorption, where the straightforward relationship between weakness in soft X-rays and ultraviolet continuum absorption apparently does not hold. Given the well-documented complexity of BAL profiles, which can include contributions from scattered flux as well as exhibit severe saturation (e.g., Ogle et al. 1999; Arav 2001), the observed lack of correlation is therefore understandable. The second absorption-line parameter, DI, also shows no evidence for a statisti- cally significant correlation with αox. An anti-correlation was predicted based on the radiation-driven wind models for QSOs of Murray et al. (1995). In their picture, X-ray 22 23 2 absorbing gas with N = 10 10 cm− is required to shield the disk wind, driven H − by resonance-line pressure from extreme ultraviolet photons, from becoming completely ionized by the soft X-rays generated near the central engine. The larger column den- sities (resulting in more negative values of αox) would be expected along lines of sight skimming the disk. The wind seen through these observing angles would have a larger transverse velocity component and therefore lower velocities of onset (or smaller values for DI). At smaller inclination angles where the line of sight intercepts the wind after it has been accelerated, the outflow is more radial, and thus the detachment index would be larger. Our data at this point do not support this simple picture, though perhaps the DI is not an appropriate measure for comparison given the complex and diverse structure in the absorption lines. The broad-band polarization percentage, P , also shows no evidence for a signif- icant correlation with αox (see Figure 6.3a). Gallagher et al. (1999) had noted that of their sample of ASCA-observed BAL QSOs, all of those with high polarization were de- tected. They speculated that if polarization is an indication of the amount of scattering occuring in the nucleus of a BAL QSO, then those with more scattering in the optical might also exhibit relatively higher X-ray fluxes due to X-rays scattered into the line of sight. Green et al. (2001) found no such correlation from their Chandra data. Though we concur with the result of Green et al. (2001), it is worth noting that at least some of their sample have high polarizations due to dust-scattering; dust is not an efficient medium for scattering X-rays through large angles. Though the slope of the continuum blueward of C iii], αB, is not correlated signif- icantly with αox, there are some interesting trends in the data (see Figure 6.3b). First, all of the BAL QSOs with the reddest continua (largest values of αB) reside at the X-ray weak end of the sample. However, not all of the X-ray weakest BAL QSOs are red. If αox can be associated with the amount of intrinsic X-ray absorption, this does suggest that the X-ray absorbing gas does not necessarily cause continuum reddening.

3We chose to use BI instead of C iv EW for this analysis as this value is insensitive to assump- tions about the shape of the C iv emission line. Nevertheless, CIV absorption EW parameterizes the same absorption; using CIV absorption EW (from Hamann et al. 1993) instead of BI also results in a high probability that this variable is unrelated to αox. 100

The search for correlations between absorption line and continuum properties with αox did not result in any obvious patterns, but those with the HR proved more fruitful. Though individual measurements of the HR suffer from large statistical errors, the ensemble of HRs from the sample offers promise for investigating the relationship of observed X-ray continuum shape to other properties. Tests of HR versus αox and versus z both resulted in statistically significant correlations (see Table 6.3). According to the evaluation of Kendall’s τ, both of these correlations are significant at the & 98% confidence level for HR versus αox and z. The results for the other statistical tests also indicated evidence for the same correlations at a similar level. The correlation of αox with HR makes physical sense if αox results from intrinsic absorption. As mentioned previously in 6.2, larger values of the hardness ratio are expected for more absorbed § X-ray continua. This understanding is further supported by the relationship between HR and z; the HR decreases as z increases. Again, this is consistent with intrinsic absorption. At higher z, the spectral signatures of intrinsic absorption are pushed to lower observed energies, and therefore occupy less of the soft bandpass.

6.4 Discussion

The sample of BAL QSOs presented in this paper offers the advantage of being much more homogeneous than those in previous hard-band surveys. For comparison, while the Green et al. (2001) sample had comparable sensitivity, the large range in redshift (z = 0.148–2.371) and ultraviolet luminosity (almost a factor of 40) makes general conclusions about BAL QSOs as a whole more difficult to derive. Our sample ranges over narrower spans of both properties with z = 1.616–2.880 and ultraviolet luminosities spanning a factor of 13. In addition, our survey objects have uniform, ≈ high-quality spectroscopic data available for meaningful comparisons between properties in different wavelength regimes. From the days of ROSAT and ASCA, the detection fraction of BAL QSOs has increased substantially. Of the 15 observed LBQS BAL QSOs, 11 were detected. Given the sensitivity of Chandra, meaningful upper limits on X-ray flux and αox could be set for those that were not detected. For the observed values of αox, our range encompasses values from 2.35 to 1.54 with four upper limits with a median value of 2.08. In − − − comparison, the mean and standard deviation for QSOs without evident ultraviolet ab- sorption is 1.56 0.14 (Brandt et al. 2001). For the detected BAL QSOs, calculating −  α (corr), the “X-ray absorption-corrected” α , gives a range of 1.82 to 1.36 with a ox ox − − median of 1.52. Even before renormalizing the X-ray continuum using the flux above − rest-frame 5 keV, not all of the BAL QSOs are observed to have αox values that would group them as X-ray weak, e.g., with α < 1.7. While the calculated X-ray fluxes do ox − suffer from large errors due to the generally low count rates, αox is not extremely sensi- tive to these errors. Another possible source of the spread is variability of the ultraviolet continuum since some of these fluxes were last measured more than 20 years ago (Hewett et al. 1995). As found by Green et al. (2001), the LoBAL QSOs in this sample (indicated by stars in Figures 6.2–6.4) also inhabit the X-ray weak end of αox. Attempting to correct the ultraviolet continua for dust extinction in these red QSOs would only push 101 them to more negative values. Given this evident distinction in the X-ray as well as the ultraviolet properties, it is important to distinguish the samples of low and high- ionization BAL QSOs when trying to generalize about the population of BAL QSOs as a whole. From the X-ray data alone, two pieces of evidence are consistent with the in- terpretation of intrinsic absorption as the cause of X-ray weakness. The first is the significant correlation of HR and αox. An X-ray spectrum exhibiting intrinsic absorption by neutral, partial-covering, or ionized gas will have fewer counts in the soft band, and thus a higher value for the hardness ratio. At larger redshifts, the absorber part of the spectrum will be shifted to lower and lower energies thus occupying a smaller fraction of the observed bandpass. Therefore, for a given column density, an absorbed QSO at a lower redshift will have a higher value of HR than the same QSO at a higher z. Exactly this trend is observed in Figure 6.4. However, making a straight identification of αox with column density is problematic for two significant reasons. For the first, Gallagher et al. (2002a) have shown that the intrinsic absorption is complex, but with the present spectral resolution and signal-to-noise ratio, the specific nature of that complexity is poorly constrained. Thus the assumption of neutral absorption, while useful for refer- ence, is probably not valid. The effect of ionization state, covering fraction, and velocity structure on measured column density from a CCD spectrum remains to be quantified. Empirically, this complexity manifests itself as additional flux at soft energies from what would be expected from a completely neutral absorber. Given that situation, the HR and αox data presented in this sample should not be translated into measurements of intrinsic column density. With that caveat, for reference, the tracks for a QSO with Γ = 2.0 and different levels of intrinsic neutral absorption are indicated in Figure 6.4b with dashed lines. The correlation of HR with z is steeper than any of the tracks; this could result from an inherent spread in intrinsic column densities.

6.5 Future Work

Many BAL QSO studies have observed the most notable of objects. While they frequently offer interesting case studies, this situation can lead to spurious connections and inhibit understanding of populations as a whole. For example, a disproportionate number of LoBAL QSOs, compared to their numbers in optically selected samples, have been observed at all wavelengths. This contributes to skewing general conclusions about QSOs and their winds based on these objects. Another selection bias results from the fact that LoBAL QSOs are generally discovered at lower redshifts than typical BAL QSOs as the Mg ii transition shifts into the wavelength regime accessible from the ground at z 0.25 rather than z 1.35 for C iv. Though HST has mitigated this somewhat by ∼ ∼ enabling the identification of some BAL QSOs at lower z, most BAL QSOs are found with ground-based spectroscopy. At their relatively smaller distances, LoBAL QSOs are much more likely to be detected in other wavelength regimes, e.g., by the Intrared Astronomical Satellite (IRAS). For non-detections, the upper limits of these relatively nearby QSOs offer much tighter constraints, e.g., on X-ray luminosity. Certain notable connections, for example that LoBAL QSOs often have lower [O iii] luminosities (e.g., Boroson & Meyers 1992; Turnshek et al. 1997), may also hold for typical BAL QSOs. 102

The [O iii] emission line is not accessible with optical spectrographs for most known high BAL QSOs, and therefore their [O iii] luminosities (as well as those of most QSOs) have not been systematically characterized. Finally, there is a notable population of ultra- luminous infrared galaxies with LoBAL QSOs (Low et al. 1989). However, the mid and far-infrared properties of BAL QSOs as a whole are completely unknown; they were not accessible with IRAS. Samples of BAL QSOs planned for observations in the future, for example with the SIRTF Guaranteed Time Observer program, have a disproportionate number of LoBAL QSOs considering their 10% fraction within the group of BAL QSOs. To avoid this pitfall, we are in the process of obtaining rest-frame ultraviolet photometry and spectroscopy for a complete sample of LBQS BAL QSOs, as described in 6.1, to supplement the information in the literature as well as the Chandra data. The § imaging and spectroscopy, obtained with the Hobby-Eberly Telescope within one month of each Chandra observation, will enable concurrent determinations of both f2500 and f2 keV for measuring αox, as well as the absorption-line parameters. Finally, these data will provide the means to investigate long-term variability of the velocity structure of the BALs to test the radiation-driven wind models of Proga, Stone & Kallman (2000). In their hydrodynamical models, they predict variability on timescales of a few years. The most comprehensive BAL QSO variability study to date, that of Barlow (1993), was not of sufficient duration to test these models. With the data of WMFH and our Hobby- Eberly Telescope spectroscopy, we can approximately triple the timescale of Barlow (1993) to span more than a decade.

Acknowledgments: We thank K. Korista for generously providing the reduced spectra from WMFH and Korista et al. (1993). This work was made possible by Chandra X- ray Center grant GO1-2105X. SCG gratefully acknowledges support from NASA GSRP grant NGT5-50277 and from the Pennsylvania Space Grant Consortium. WNB thanks NASA LTSA grant NAG5-8107. 103

Fig. 6.2 (a) The BALnicity index versus αox. (b) The detachment index versus αox. The spectroscopic data are from WMFH; neither of these absorption-line parameters has a statistically significant correlation with weakness in X-rays. In both plots, the solid and dashed gray lines represent the mean and one standard deviation of αox for typical radio-quiet QSOs from (Brandt et al. 2001). 104

Fig. 6.3 (a) The broad-band continuum polarization, P , versus αox. There is no statis- tically significant correlation between these two properties. (b) The spectral index of α the continuum blueward of C iii], α for F ν− B , versus α . Note that the reddest B ν ∝ ox BAL QSOs (those with the largest values of αB) are also X-ray weak. Both P and αB are from Hutsem´ekers et al. (1998). The vertical gray lines in both plots are as defined in the caption to Figure 6.2. 105

Fig. 6.4 (a) Hardness ratio versus αox. The weakest X-ray sources (relative to their ultraviolet continuum flux) tend to have the hardest spectra. This is consistent with absorption as the primary cause of X-ray weakness. The vertical gray lines are as defined in the caption to Figure 6.2. (b) Hardness ratio versus z. There is clearly an anti- correlation of HR and z; this trend would be expected if the observed weakness in X-rays results from intrinsic absorption. As z increases, the absorbed continuum is redshifted out of the ACIS bandpass. The dashed curves represent the expected HR for a power- 20 2 law continuum with Γ = 2.0, Galactic N = 2.8 10 cm− (the median value for this H × sample), and an intrinsic, neutral column density with the log value labeled in the figure. These curves are meant for reference only as the actual absorption is likely to be complex (Gallagher et al. 2002a). 106

Chapter 7

Conclusions

Since the work of this thesis was begun in 1997, significant progress has been made in the understanding of the X-ray properties of highly absorbed QSOs. The employment of sensitive X-ray instruments with 2–10 keV bandpass such as ASCA and Chandra drove these advances. As of mid-2002, the following is now generally known:

A value of α 1.7 for a radio-quiet QSO is a strong indicator of the presence of • ox. − both ultraviolet and X-ray absorption (Brandt et al. 2000; Gallagher et al. 2001b). BAL QSOs exhibit the extremes of both. Though extremely faint in soft X-rays, BAL QSOs are not intrinsically X-ray weak. • In terms of both X-ray spectral shape and αox, they have properties consistent with the population of normal radio-quiet QSOs (Green et al. 2001; Gallagher et al. 2002a). BAL QSOs suffer from complex, intrinsic absorption of their X-ray spectra. For • those with moderate-quality X-ray data, the column densities can be constrained 23 2 to lie in the range from (0.1–0.5) 10 cm− (Gallagher et al. 2002a). These val- × ues are consistent with those inferred from X-ray weakness in exploratory surveys (Chapter 6). Despite these high column densities, BAL QSOs in general do not suffer from extreme dust extinction. This implies that the X-ray absorbers are either virtually dust-free or that the ultraviolet continuum does not pass through the X-ray absorbing gas. Empirically, LoBAL QSOs are observed to be weaker in X-rays than normal BAL QSOs • (Green et al. 1999; Chapter 6). This is conceivably due to larger column densi- ties of obscuring gas along the line-of-sight to the X-ray source and may result from orientation or a fundamental difference in some physical parameter related to LoBAL QSOs. If the observations of Mrk 231 and H 1413+117 are representa- tive, the only X-rays from LoBAL QSOs that can reach the observer are indirect, scattered or reflected from circumnuclear material (Gallagher et al. 2002a,b).

7.1 Plans for the Future

In general, the conclusions from this work about the properties of the X-ray ab- sorbing gas in BAL QSOs are consistent with the requirements for the shielding gas introduced by Murray et al. (1995) for their radiatively driven winds. Principally, the 23 24 2 values for the column densities (NH= 10 − cm− ), originally considered outrageously large, are being directly and indirectly observed. Furthermore, if the ultraviolet contin- uum passes through the X-ray absorber, the extremely low dust-to-gas ratio implied is 107 consistent with the high ionization state of the proposed shielding gas. However, this model requires more exacting observational tests. In particular, the ionization state of the gas has not been constrained, and the velocity structure of the X-ray absorber is completely unknown. These values are of fundamental importance for unifying X-ray and ultraviolet absorbers. Addressing these issues requires a deep, X-ray grating obser- vation of the brightest BAL QSO. PG 2112+059 (Gallagher et al. 2001b) remains the best target for this work, and we have proposed it for XMM-Newton Reflection Grating Spectrometer observations. This particular proposal offers the most promise for con- straining a fundamental paramenter of AGN physics, the mass outflow rate from the nucleus. Completing the Large Bright Quasar Survey BAL QSO sample also offers promise for further understanding the relationship between X-ray and ultraviolet absorption. Though no obvious correlations between the two are apparent at present, a sample of 11 detections is probably not sufficient for uncovering any connections. In addition, more than a decade has passed since the Weymann et al. (1991) spectroscopy was performed, and changes in absorption properties in the X-ray and/or ultraviolet since that time could obscure relationships. Our ongoing project to observe all of this sample concurrently with Chandra and the Low Resolution Spectrograph on the Hobby-Eberly Telescope will remove problems due to possible variability from our analysis. These sample data, combined with the Weymann et al. (1991) observations, will also permit searches for velocity structure variability over much longer timescales than probed by Barlow (1993). If velocity changes are not found, the current modeling of Proga et al. (2000) will need to revisit the physical scales of the launching radius of the wind. In terms of more ambitious projects, spectroscopy of this sample would ideally encompass the rest-frame optical and ultraviolet spectral regions, from the [O iii] through Lyα emission lines. Such studies will enable comparisons with lower redshift samples which have already shown interesting connections with X-ray properties. For example, the luminosity of [O iii] has been shown to correlate significantly with ultraviolet and X-ray absorption for less luminous AGN (Brandt et al. 2000). Extending this type of analysis to higher luminosities and larger samples of BAL QSOs could be profitable. Due to the redshift of most of our sample, these projects will require near-infrared spectroscopy with instruments such as JCAM on the Hobby-Eberly Telescope. The observational tools to understand the nature of the energetic winds from lumi- nous QSOs are available. With 10 m class telescopes, the flexibility of queue-scheduling, near-infrared spectrographs and sensitive X-ray observatories, these proposed projects are within reach. 108 Bibliography

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Zamorani, G. et al. (1981) X-ray studies of quasars with the Einstein Observatory. II. Astrophysical Journal 245:357. Vita Sarah Connoran Gallagher Education The Pennsylvania State University State College, Pennsylvania 1997–Present Ph.D. in Astronomy & Astrophysics, expected in May 2002 Area of Specialization: Extragalactic X-ray Astronomy Yale College New Haven, Connecticut 1991–1995 B.A. in Physics (graduated in seven semesters), cum laude with distinction in the major Awards and Honors NASA Graduate Student Research Program Fellowship 2000–Present Pennsylvania Space Grant Consortium Fellowship 1998–2000, 2000–Present Zaccheus Daniel Foundation for Astronomical Science Grant 1999, 2000 University Graduate Fellowship 1997–1998 Roberts Fellowship, Eberly College of Science 1997–1998 Summer Study Grant, Holderness School 1996 Research Experience Doctoral Research The Pennsylvania State University 1998–Present Thesis Advisor: Prof. William N. Brandt Conducted exploratory Chandra observations of a well-defined sample of Broad Ab- sorption Line (BAL) QSOs (coupled with optical spectroscopy with the Hobby-Eberly Telescope) to characterize the X-ray properties of the class as a whole. Analyzed spec- troscopic X-ray observations of brighter BAL QSOs and other heavily absorbed QSOs with Chandra and XMM-Newton to study absorption structure. Graduate Research The Pennsylvania State University 1998–1999 Research Advisor: Prof. Jane Charlton With HST multicolor images, identified and characterized the luminous star cluster populations in Stephan’s Quintet, a compact group of interacting galaxies. Connected the dynamical history of the galaxies to the ages of the star clusters in different regions. Undergraduate Research Yale University 1994–1995 Research Advisor: Prof. Paolo Coppi Modeled gravitational lensing systems in support of the QUEST project. Teaching Experience Guest Lecturer The Pennsylvania State University 1997–Present Prepared and taught college classes in Introductory Astronomy for Non-Majors (regular and honors sections), Freshman Seminar for Majors, and in-service teacher workshops. Teaching Assistant The Pennsylvania State University 1998 Prepared and taught independent astronomy laboratory sections for college students. Physics Teacher Holderness School, Plymouth, New Hampshire 1995–1997 Taught conceptual and honors physics to juniors and seniors in high school. As the sole physics teacher, developed curriculum to integrate computer technology into the physics classroom. Coached soccer and lacrosse, supervised a girls’ dormitory, co-led student groups for 10-day winter camping trips, and advised senior projects.