The Warm Hot Environment of the Milky Way

Transcription

1 The Warm Hot Environment of the Milky Way DISSERTATION Presented in Partial Fulfillment of the Requirements for the Degree Doctor of Philosophy in the Graduate School of The Ohio State University By Rik Jackson Williams ***** The Ohio State University 2006 Dissertation Committee: Approved by Professor Smita Mathur, Adviser Professor David H. Weinberg Professor Richard W. Pogge Adviser Astronomy Graduate Program

2 ABSTRACT I present an investigation into the local warm hot gaseous environment of the Milky Way as observed through highly ionized metal absorption lines in ultraviolet and X-ray spectra. These X-ray lines (primarily O vii) had been reported at redshifts consistent with zero in previous studies of background quasars; however, it has been unclear whether this gas exists close to the Galaxy (within a few tens of kpc) or extends far out into intergalactic space, thereby comprising most of the mass in the local universe. Additionally, highly ionized Ovi high velocity clouds (HVCs), some of which are associated with the ubiquitous extended neutral hydrogen HVCs seen around the Galaxy, had been extensively studied. However, the distance to the Ovi HVCs, and their relation to the X-ray lines, remained undetermined. With three of the highest quality Chandra grating spectra of extragalactic sources to date, a large number of z = 0 absorption lines are detected; the FUSE spectra of these same objects show low and high velocity O vi absorption. Using advanced curve of growth and ionization balance analysis, limits are placed on the velocity dispersion, temperature, and density of the warm hot gas along these lines of sight. In none of these cases can the absorption be placed conclusively at Galactic ii

3 or extragalactic distances. However, in two of the three cases (Mrk 421 and Mrk 279), the observed O vi UV absorption components are found to be inconsistent with the X-ray absorber, indicating that the X-ray absorption is either extragalactic or traces a previously undiscovered Galactic component. The third sightline (PKS ) exhibits absorption with properties more similar to Mrk 421 than Mrk 279; thus, there may be more than one physical process contributing to the observed absorption along any given sightline. While the X-ray components of this research exclusively employ Chandra data, the XMM Newton mission can in principle be used for the same purpose. XMM s effectiveness in observations of WHIM lines is quantitatively analyzed in the context of two recently detected intervening WHIM systems toward Mrk 421. The XMM grating spectrograph is found to be inferior to Chandra/LETG due to lower resolution and narrow detector features that hinder the detection of unresolved lines. iii

4 Dedicated to Walter J. Williams iv

5 ACKNOWLEDGMENTS I cannot thank my advisor, Smita Mathur, enough for the fantastic research opportunities, constant support, and for fending off the wolves when necessary (while simultaneously teaching me how to do it for myself). I look forward to many years of collaboration with her, writing last minute proposals for observations on unfamiliar instruments. Many thanks also to Rick Pogge, who (as my effective first year advisor) got me started on some excellent projects here at OSU and has been a continuous source of support and advice throughout. Likewise, I thank David Weinberg and Andy Gould for their consultation on a number of matters both political and scientific. None of this would have been possible without Fabrizio Nicastro and Martin Elvis giving me access to their one-of-a-kind data and teaching me how to analyze and interpret it. I get by with a little help from my friends (Lennon & McCartney 1967) in the literal sense, particularly Juna Kollmeier, Reni Ayachitula, Amy Stutz, and Iljie Kim Fitzgerald. And, of course, a little help from my family, with their unceasing (if bemused) encouragement and support of my foray into academia. v

6 Although they may not have been as directly involved in my research as the people specifically mentioned above, I am indebted to those other past and present members of the OSU Astronomy Department who have transformed it into the astronomy field s foremost venue for graduate research and scientific interaction. Generous financial support for this work was provided by an Ohio State University Presidential Fellowship, Chandra award AR5 6017X (issued by the Chandra X-ray Observatory Center, which is operated for and on behalf of NASA under contract NAS ), and the National Radio Astronomy Observatory. I salute the efforts of the Chandra, FUSE, and XMM scientific and support staff for making these excellent missions possible. vi

7 VITA January 29, Born Silverton, Oregon, USA B.S. Astronomy, California Institute of Technology M.S. Astronomy, The Ohio State University Graduate Fellow, The Ohio State University Graduate Research Associate, The Ohio State University Presidential Fellow, The Ohio State University PUBLICATIONS Research Publications 1. I. N. Reid, J. D. Kirkpatrick, J. E. Gizis, C. C. Dahn, D. G. Monet, R. J. Williams, J. Liebert, and A. J. Burgasser, Four Nearby L Dwarfs, AJ, 119, 369, (2000). 2. J. D. Kirkpatrick, I. N. Reid, J. Liebert, J. E. Gizis, A. J. Burgasser, D. G. Monet, C. C. Dahn, B. Nelson, and R. J. Williams, 67 Additional L Dwarfs Discovered by the Two Micron All Sky Survey, AJ, 120, 447, (2000). 3. J. E. Gizis, D. G. Monet, I. N. Reid, J. D. Kirkpatrick, J. Liebert, and R. J. Williams, New Neighbors from 2MASS: Activity and Kinematics at the Bottom of the Main Sequence, AJ, 120, 1085, (2000). vii

8 4. R. J. Williams, R. W. Pogge, and S. Mathur, Narrow-Line Seyfert 1 Galaxies from the Sloan Digital Sky Survey Early Data Release, AJ, 124, 3042, (2002). 5. S. Mathur, and R. J. Williams, Chandra Discovery of the Intracluster Medium Around UM 425 at Redshift 1.47, ApJ, 589, L1, (2003). 6. R. J. Williams, S. Mathur, and R. W. Pogge, Chandra Observations of X-ray Weak Narrow-Line Seyfert 1 Galaxies, ApJ, 610, 737, (2004). 7. F. Nicastro, S. Mathur, M. Elvis, J. Drake, T. Fang, A. Fruscione, Y. Krongold, H. Marshall, R. Williams, and A. Zezas, The mass of the missing baryons in the X-ray forest of the warm hot intergalactic medium, Nature, 433, 495, (3 February 2005). 8. F. Nicastro, S. Mathur, M. Elvis, J. Drake, F. Fiore, T. Fang, A. Fruscione, H. Marshall, and R. Williams, Chandra Detection of Two Warm Hot IGM Filaments along the Line of Sight to Mkn 421, ApJ, 629, 700, (2005). 9. R. J. Williams, S. Mathur, F. Nicastro, M. Elvis, J. J. Drake, T. Fang, F. Fiore, Y. Krongold, Q. D. Wang, and Y. Yao, Probing the Local Group Medium Toward Mkn 421 with Chandra and FUSE, ApJ, 631, 856, (2005). 10. Q. D. Wang, Y. Yao, T. M. Tripp, T. T. Fang, W. Cui, F. Nicastro, S. Mathur, R. J. Williams, L. Song, and R. Croft, Warm Hot Gas in and around the Milky Way: Detection and Implications of O VII Absorption Toward LMC X 3, ApJ, 635, 386, (2005). 11. R. J. Williams, S. Mathur, F. Nicastro, and M. Elvis, XMM Newton View of the z > 0 Warm Hot Intergalactic Medium Toward Markarian 421, ApJ, 642, L95, (2006). 12. R. J. Williams, S. Mathur, and F. Nicastro, Chandra Detection of Local Warm Hot Gas Toward Markarian 279, ApJ, 645, 179, (2006). viii

9 FIELDS OF STUDY Major Field: Astronomy ix

10 Table of Contents Abstract Dedication Acknowledgments Vita List of Tables List of Figures ii iv v vii xii xiii Chapter 1 Introduction Missing Baryons at Low Redshift Relation to Previous Work Scope of the Dissertation Chapter 2 The Markarian 421 Sightline Observations and Data Preparation Chandra FUSE Line Measurements Absorption Line Diagnostics Doppler Parameters x

11 2.3.2 Column Densities Temperature and Density Constraints Discussion Potential Caveats Where does the X-ray absorption originate? Comparisons to Other Studies Summary and Future Work Chapter 3 The Markarian 279 Sightline Data Reduction and Measurements Chandra FUSE Analysis Doppler Parameters and Column Densities Temperature and Density Diagnostics The AGN Warm Absorber Discussion Comparison to the Mrk 421 Sightline Origin of the Absorption Conclusions Chapter 4 The PKS Sightline Data Reduction and Measurements Chandra xi

12 4.1.2 FUSE Analysis Doppler Parameters and Column Densities Temperature and Density Diagnostics z = Absorption Reported by Fang et al Discussion Comparison to Other Lines of Sight Where is the Absorption? Comparison to Nicastro et al. (2002) Conclusions Chapter 5 Instrumental Considerations: Chandra or XMM Newton? Data Reduction and Measurements Discussion Disputed Results Conclusion Chapter 6 Summary and Future Work Individual X-ray Sightlines The Importance of Spectral Fidelity Future Prospects X-ray Observations Longer Wavelengths Bibliography xii

13 List of Tables 2.1 Observed z 0 lines Observed z 0 lines Observed z 0 absorption lines Chandra observation log Observed z 0 absorption lines Observed z 0 absorption lines XMM Newton observation log Absorption line equivalent width measurements xiii

14 List of Figures 2.1 Mkn 421 Chandra LETG spectrum Mrk 421 FUSE spectrum near the Ovi line Ovii curve of growth diagnostics Ovi curve of growth diagnostics Temperature and density diagnostics from oxygen lines Temperature and density diagnostics with solar abundances Temperature and density diagnostics with shifted abundances Ionic abundance models for the cooler (likely Galactic) ions Ionic abundances vs. temperature for possible extragalactic ions, low density case Ionic abundances vs. temperature for possible extragalactic ions, high density case Full Chandra spectrum of Mrk Å Chandra spectrum of Mrk Velocity plots of the local O vii and O vi absorption lines Curve of growth diagnostics for the O vii K series Curve of growth analysis for the O vi UV absorption Temperature and density constraints from Ovii and Ovi, b = 100 km s Temperature and density constraints, b = 200 km s xiv

15 4.1 Chandra ACIS S/LETG continuum fit Chandra HRC S/LETG continuum fit Detected z = 0 absorption lines (ACIS S/LETG) Detected z = 0 absorption lines (HRC S/LETG) Åregion of the FUSE spectrum Ovii curve of growth analysis Ovi curve of growth analysis Oxygen ion temperature and density constraints X-ray ion temperature and density constraints (low b) X-ray ion temperature and density constraints (high b) Chandra spectrum near the Oviii z = wavelength XMM Newton spectrum of Mrk RGS1 and RGS2 instrumental response functions Line spread functions for Chandra and XMM RGS xv

16 Chapter 1 Introduction Look on my works, ye mighty, and despair! Ozymandias, Percy Bysshe Shelley Gonna get my PhD I m a teenage lobotomy Teenage Lobotomy, The Ramones 1.1. Missing Baryons at Low Redshift Over the past fourteen billion years, the baryonic mass found in the intergalactic medium (IGM) a tenuous web of gas bridging the gaps between galaxies and clusters is thought to outweigh the baryons found in all other sources, stars, galaxies, and the hot gas that dominates the mass of clusters of galaxies. Indeed, at high redshifts (z > 2) the forest of Lyα absorption lines seen in spectra of distant quasars reveals a vast network of cool, photoionized hydrogen that is consistent with the expected baryon density at those redshifts (Weinberg et al. 1997). At more recent epochs, 1

17 however, the process of structure formation has shock heated this intergalactic gas to produce the warm hot IGM (WHIM; Cen & Ostriker 1999; Davé et al. 2001) with temperatures of T K densities cm 3, or cosmic overdensities of δ This WHIM gas has proved extremely difficult to detect, resulting in a discrepancy between the observed baryon census and predictions from the cosmic microwave background (Bennett et al. 2003). At such low densities and high temperatures, the combination of collisional and photoionization renders most of the gas too highly ionized to be detected through its Lyα absorption, though some broad Lyα systems at low redshift have been reported (Sembach et al. 2004; Richter et al. 2004). Moreover, its extremely low density prevents thermal and/or line emission from the WHIM from being detected even with the most sensitive current instruments. Heavier elements such as oxygen, nitrogen, and neon would be highly (but not fully) ionized in such a medium, and these metals are predicted to provide a unique view of the WHIM through their higher energy UV and X-ray resonance absorption lines (Perna & Loeb 1998; Hellsten et al. 1998; Fang, Bryan, & Canizares 2002). Though they are quite weak, these WHIM lines are now detectable in principle with the advent of such facilities as the Chandra X-ray Observatory, XMM Newton, the Hubble Space Telescope, and the Far Ultraviolet Spectroscopic Explorer (FUSE). Detections of such intervening WHIM filaments, with varying levels of confidence, 2

18 have been reported along several lines of sight (Nicastro et al. 2005a,b; Mathur et al. 2003; Fang et al. 2002). The total baryonic mass reported by Nicastro et al. (2005a) is indeed consistent (within the admittedly large errors) with the aforementioned baryon deficit at low redshifts. Since most galaxies are expected to trace the same cosmic overdensities as the web of WHIM filaments, it would be no surprise if the Galaxy itself resided in such a filament. Indeed, X-ray spectra of several quasars have shown likely z = 0 Ovii absorption, but it is unclear whether this absorption is actually due to the nearby WHIM or is instead a component of the Galaxy itself, such as a hot halo or corona (or some combination of the two). Some O vii absorption has indeed been found within 50 kpc of the Galaxy (Wang et al. 2005), but this is unlikely to be uniformly distributed. Simulations of the Local Group strongly indicate that a large amount of warm hot gas is expected near zero redshift (Kravtsov et al. 2002). Thus, in reality the X-ray absorption is likely to be caused by a combination of Galactic and extragalactic components, perhaps with one dominating the other in certain directions. Further complicating the issue is the presence of other gaseous components of unknown origin. H I high velocity clouds (HVCs) have a velocity distribution inconsistent with Galactic rotation and therefore are thought to be either neutral gas high in the Galactic halo or cooling, infalling gas from the surrounding IGM. Along many lines of sight studied with FUSE, high velocity Ovi absorption lines at 3

19 velocities coincident with the Hi HVCs are seen, while in some other directions the high velocity Ovi is present even in the absence of Hi emission at that velocity (Sembach et al. 2003). Some of these latter, unassociated Ovi HVCs were found to be at rest in the Local Group rest frame (as a population), indicating that they may be extragalactic in origin (Nicastro et al. 2003). On the other hand, many of these O vi HVCs also show absorption from lower ionization states that are unlikely to arise in a low density, warm hot IGM (Sembach 2003). While the evidence appears to point to both Galactic and extragalactic characteristics for the O vi HVCs, their connection to the highly ionized gas seen in X-rays (if any) is unknown. Part of the problem has been the tremendous amount of Chandra observing time that is required to obtain a high quality grating spectrum of an extragalactic source: since X-ray telescopes are essentially photon counting devices, thousands of scarce (compared to optical telescopes), high energy X-ray photons must be collected for each resolution element in order to detect even relatively strong WHIM lines ( 20 må, corresponding to N OVII cm 2 ). In the past few years, however, there have been several opportunities to overcome this difficulty: observing AGN only during extremely bright flares (as with Mrk 421), re analyzing long exposures that were originally performed for other purposes (Mrk 279), and co adding many short calibration observations of the same object taken over the past seven years (PKS ). 4

20 The three aforementioned AGN have thus been observed with these techniques to unprecedented levels of sensitivity with the Chandra gratings. In this dissertation I present measurements and analyses of the z 0 X-ray absorption lines seen in these spectra with the goal of determining the location and physical properties of the absorbing material, and its connection to the Ovi absorption seen along each line of sight in FUSE. Such high quality X-ray data have not been previously available, and new techniques are developed to obtain the greatest amount of information on the physical state of the absorbing medium (while overcoming some unique aspects of the data, such as a complicated instrumental response function and far lower spectral resolution than is available in the FUSE data). The efficacy of the XMM Newton observatory for these studies is also investigated Relation to Previous Work Prior to this work, a few studies had reported the presence of z = 0 Ovii absorption, namely toward 3C 273 (Fang et al. 2003) and PKS (Nicastro et al. 2002). Since this latter detection was published, far more X-ray data became available in the form of calibration observations; the analysis in Chapter 4 takes all this available data into account and compares my new results to those of Nicastro et al. (2002). New observations of 3C 273 are also available; however, hot gas from a supernova remnant is likely to lie along this line of sight. Due to the low 5

21 velocity resolution of Chandra ( 700 km s 1 at the Ovii wavelength), absorption from the remnant would be fully blended with any WHIM or Galactic corona Ovii, making this sightline of limited use. Constrained simulations of the Local Group by Kravtsov et al. (2002) provide a strong theoretical basis for the existence of z = 0 WHIM while predicting roughly where the strongest absorption can be expected. In fact, far more literature on the z > 0 WHIM has been published. Fang et al. (2002), Mathur et al. (2003), and McKernan et al. (2003) reported early tentative detections of the WHIM toward several sources, and I address the Fang et al. (2002) detection of z = Oviii toward PKS in Chapter 4. By far the most confident detection of the z > 0 WHIM to date is in the Mrk 421 spectrum (which I analyze in Chapter 2) by Nicastro et al. (2005a). Two WHIM filaments were unambiguously detected in this spectrum, allowing the temperature density, and relative metal abundances of the WHIM to be estimated. These provide a valuable set of parameters to compare with those derived from the z = 0 absorption. The detections of these WHIM filaments are also revisited with XMM Newton in Chapter Scope of the Dissertation The data, analysis, and interpretation of the X-ray and UV data along three extragalactic sightlines are the focus of this dissertation. Owing to the individual 6

22 peculiarities of the observations and the variety of analysis techniques required for the different data, the three following chapters are each devoted to one line of sight: Markarian 421 (Chapter 2), Markarian 279 (Chapter 3), and PKS (Chapter 4). Chapter 5 presents a comparison of the XMM Newton data of Mrk 421 with the previously reported Chandra detection of two WHIM filaments along this line of sight, and quantitatively describes why XMM is unable to detect these filaments. Finally, I summarize the results of this dissertation and comment on present and future avenues for this research. A major fraction of the research presented in this dissertation has been published in the scientific literature. Chapter 2 is largely taken from Williams et al. 2005, ApJ, 631, 856; chapter 3 has appeared as Williams et al. 2006, ApJ, 645, 179; and most of chapter 5 also appears in Williams et al. 2006, ApJ, 642, L95. 7

23 Chapter 2 The Markarian 421 Sightline Through a program of observing nearby blazars in outburst phases, we have obtained high quality Chandra and FUSE spectra of Mkn 421, sufficient to study in detail the local WHIM (and Galactic halo/thick disk) absorption. Here I report on these observations, and the inferred properties of the local absorption Observations and Data Preparation Chandra A full description of the Chandra observations, data reduction, and continuum fitting can be found in Nicastro et al. (2005a); a brief summary follows. Mkn 421 was observed during two exceptionally high outburst phases for 100 ks each as part of our Chandra AO4 observing program: one at f 0.5 2keV = erg s 1 cm 2 with the Low Energy Transmission Grating (LETG) combined with the Advanced CCD Imaging Spectrometer Spectroscopic (ACIS-S; Garmire et al. 2003) array, and another at f 0.5 2keV = erg s 1 cm 2 with the High Resolution 8

24 Camera Spectroscopic (HRC-S; Murray & Chappell 1985) array and LETG. Each of these observations contains 2500 counts per resolution element at 21.6 Å. Additionally, another short observation of Mkn 421 was taken with HRC/LETG (29 May 2004), providing another 170 counts per resolution element. These three spectra were combined over the Å range to improve the signal to noise ratio (S/N 55 at 21Å with Å binning). The final coadded spectrum of Mkn 421 is one of the best ever taken with Chandra: it contains over 10 6 total counts with 6000 counts per resolution element at 21.6 Å, providing a 3σ detection threshold of W λ 2 må (N OVII = cm 2 for an unsaturated line). Effective area files (ARFs) for each observation were built using CIAO 1 v3.0.2 and CALDB 2 v Those pertaining to the ACIS/LETG observations were corrected for the ACIS quantum efficiency degradation 3 (Marshall et al. 2003). For the HRC/LETG observations, the standard ARFs were used. Each ARF was then convolved with the relevant standard redistribution matrix file (RMF), and the convolved RMFs were weighted by exposure time, rebinned to the same energy scale, and averaged to provide a response file for the coadded spectrum. Using the CIAO fitting package Sherpa 4, we initially modeled the continuum of Mkn 421 as a simple power law and a Galactic absorbing column density of 1 cxc.harvard.edu/ciao/ 2 cxc.harvard.edu/caldb/ 3 See also cxc.harvard.edu/ciao/why/acisqedeg.html 4 cxc.harvard.edu/sherpa/ 9

25 N H = cm 2 (Dickey & Lockman 1990), excluding the 48 57Å HRC chip gap region. Metal abundances for the Galactic gas were then artificially adjusted to provide a better fit around the O I and C I K edges near 23Å and 43Å respectively. This is not intended to represent actual changes to the absorber composition, but rather to correct uncertainties in the instrument calibration. These adjustments affect the continuum mostly near the carbon, oxygen, and neon edges, but individual narrow absorption lines are unaffected. After this fit there were still some systematic uncertainties in the best fit continuum model; these were corrected with broad (FWHM = Å) Gaussian emission and absorption components until the modeled continuum appeared to match the data upon inspection. Indeed, the residuals of the spectrum to the final continuum model have a nearly Gaussian distribution, with a negative tail indicating the presence of narrow absorption lines (see Nicastro et al. 2005a, Figure 8) FUSE Mkn 421 was also observed with FUSE as part of our Director s Discretionary Time observing program on January 2003 for a total of 62.8 ks. An additional 21.8 ks observation from 1 December 2000 was also available in the archive. We used the time tagged, calibrated data from only the LiF1A detector, since inclusion of the LiF2B data provides a small ( 20%) increase in S/N but degrades the overall 10

26 spectral resolution. 5 These two observing programs consist of four observations, which in turn contain a total of 29 individual exposures. The wavelength scales of each observation s constituent exposures were cross correlated and shifted (typically by 1 2 pixels) to account for slight uncertainties in the wavelength calibration. The exposures for each observation were checked for consistency and coadded, weighted by exposure time. The resulting four spectra were then cross-correlated against each other, coadded (with a 10% downward shift in flux applied to the 2000 observation due to source variability), and binned by 5 pixels (0.034 Å, or one half of the nominal 20km s 1 resolution) providing a S/N of 17 near 1032Å. To check the absolute wavelength calibration we followed the method of Wakker et al. (2003), using their 4 component fit to the Murphy et al. (1996) Green Bank H I 21 cm data as a velocity reference. They find four main components of H I emission with an N H weighted average velocity of 31.7km s 1. In the FUSE spectrum, the Si II λ å and Ar I λ å lines are expected to trace the same gas as the H I emission. Each UV line was fit with two Gaussian components in Sherpa, giving average velocity offsets of 30.9 and 34.9km s 1 respectively. These agree well with the H I data, though the slight difference between the Ar I and Si II measurements suggest at least a 4km s 1 intrinsic wavelength uncertainty. 5 See the FUSE Data Analysis Cookbook v1.0, fuse.pha.jhu.edu/analysis/analysis.html 11

27 2.2. Line Measurements To find and identify narrow absorption lines in the Chandra spectrum of Mkn 421, we visually inspected small (2 5 Å) regions of the spectrum, beginning around the rest wavelength of Ovii Kα (21.602Å) since this tends to be the strongest z = 0 X-ray absorption line (e.g. Nicastro et al. 2002; Fang et al. 2003; Chen et al. 2003). Three Ovii Kα (Figure 2.1) absorption features were found: one at z = 0, one at z = 0.011, and one at z = (with typical redshift errors of 0.001). There is also a strong feature which is 3σ from the Ov Kα rest wavelength, but is more likely Ovii Kα at v +900km s 1 relative to the blazar. A close pair of lines consistent with Lyα at this velocity has been observed (Shull et al. 1996; Penton et al. 2000), so this may be indicative of an inflow to Mkn 421 or uncertainty in the blazar redshift (based on rather old spectrophotometric measurements by Margon et al. 1978). A weak Ovi Kα line is seen at the rest wavelength of 22.02Å. Other regions of the spectrum were then searched for lines corresponding to these systems, with particular emphasis paid to strong transitions of the most abundant elements (C, N, O, and Ne). All in all there were 13 lines marginally or strongly detected at z 0 (including the Nvii, Ov, and Arxv upper limits), 3 at z = 0.011, and 7 at z = The latter two systems are the subject of other papers (Nicastro et al. 2004a,b) and thus will not be discussed further here. 12

28 These 13 z = 0 X-ray lines were fitted in Sherpa with narrow Gaussian features superposed on the fitted continuum described in (see Figure 2.1). We are excluding the strong O I (23.51Å) line since it arises in the neutral ISM and is not of interest here, as well as the O 2 (23.34Å) absorption since it coincides with a strong instrumental feature and cannot be accurately measured. Due to the FWHM = 0.04Å ( 600km s 1 ) LETG resolution the lines are all unresolved, so only the position and equivalent width of each line are measured. Errors are calculated using the projection command in Sherpa, allowing the overall continuum normalization to vary along with all parameters for each line. The resulting line parameter estimates are presented in Table 2.1. The 0.02 Å systematic wavelength uncertainty of the LETG 6 is in most cases larger than the statistical uncertainty of the line centroid; thus, Table 2.1 lists whichever is greater. Additionally, a meaningful upper error bar on the C vi equivalent width could not be calculated with Sherpa. In this case, the FWHM was frozen at the instrumental resolution and the error was recalculated; a visual inspection confirms the new limit to be more reasonable. Upper limits for the Ov, Nvii, Arxv, and Nex lines were calculated with both the position and FWHM frozen. The FUSE spectrum (Figure 2.2) shows a strong, broad low velocity O vi 1032Å absorption line at z 0 due to gas in the Galactic thick disk and halo (Savage et al. 2003). An asymmetric wing on the red side of this line is evident, 6 cxc.harvard.edu/cal/ 13

29 possibly a kinematically distinct HVC. We fitted the Ovi 1032Å line in Sherpa using a constant local continuum (in a ±2Å window) and two Gaussian absorption components: one for the v 0 Ovi LV line, and one at v 100km s 1 for the HVC. No H 2 contamination is seen at the Ovi 1032Å wavelength when absorption templates are fit to the other H 2 lines seen in the spectrum. The 1037Å line is somewhat blended with a single H 2 absorption line; this is taken into account with another narrow Gaussian. From this fit, we find equivalent widths of 18.6 ± 5.6 må for the 1032Å HVC and ± 7.9 må for the Galactic component. The best fit model for the HVC is fairly robust and not sensitive to variations in the initial parameters; however, the derived equivalent width of 18.5 ± 5.6 må is lower than the 37 ± 11 ± 29 må (errors are statistical and systematic, respectively) measured by Wakker et al. (2003) in the initial 21.8 ks observation. They employed a direct integration method which may not have taken into account the substantial blending of the Galactic Ovi LV with the HVC. Our total Ovi equivalent width (LV+HVC = 279 ± 10 må) is in good agreement with their value of 285 ± 20 må. Deblending the Ovi 1037Å line is less certain due to the presence of adjacent Galactic C II and H 2 absorption. A flat continuum was again employed from Å and single Gaussian components were used to fit the C II, Ovi, and H 2 absorption lines. The HVC on the 1032Å Ovi line should also appear at 1038Å with W λ = 0.50 W λ (1032Å). Although this component is too weak to be detected directly, it could cause the measurement of the Galactic LV O vi 1037 Å 14

30 line to be systematically high. Another absorption Gaussian with one half of the 1032Å HVC equivalent width (and with the same FWHM and velocity offset) was included in the 1037Å Ovi line fit to account for this. Table 2.1 lists the measured properties of the Ovi LV and HVC absorption lines Absorption Line Diagnostics Doppler Parameters To convert the measured equivalent widths to ionic column densities, we calculated curves of growth for each absorption line over a grid of Doppler parameters (b = km s 1 ) and column densities (log N H /cm 2 = ), assuming a Voigt line profile. Since the X-ray lines are unresolved, b cannot be measured directly. It can, however, be inferred from the relative strengths of the three measured O vii K series lines. These lines are produced by the same ionic species, so in an unsaturated medium W λ f lu λ 2 where f lu is the oscillator strength. The expected equivalent width ratio of Ovii Kβ to Kα is then W λ (Kβ)/W λ (Kα) = 0.156, so the measured value of 0.49 ± 0.09 indicates that the Kα line is saturated. On the other hand, the measured Ovii Kγ/Kβ ratio is 0.43 ± 0.16, in agreement with the predicted (unsaturated) value of

31 These line ratios by themselves are insufficient to determine the physical state of the Ovii absorbing medium since b and N OVII are degenerate: the Kα line saturation could be due to high column density, low b, or a combination of both. However, given an absorption line with a measured equivalent width and known f lu λ 2 value, the inferred column density as a function of the Doppler parameter can be calculated. The measured equivalent width (and errors) for each transition thus defines a region in the N OVII b plane. Since the actual value of N OVII is fixed, b and N OVII can be determined by the region over which the contours overlap; i.e. the range of Doppler parameters for which the different transitions provide consistent N OVII measurements. Figure 2.3 shows such 1σ contours for the three measured O vii transitions. As expected, the inferred N OVII is nearly constant in the unsaturated regime (large b), and rises sharply at low b as the lines begin to saturate. At each value of b, the differences (log N αβ ) = log(n Kα ) log(n Kβ ) and (log N αγ ) = log(n Kα ) log(n Kγ ) were calculated, along with the errors on each (log N). The quantity (log N αβ ) is consistent with zero at the 1σ and 2σ levels for 15 < b < 46km s 1 and 13 < b < 55km s 1 respectively, while (log N αγ ) provides limits of 31 < b < 50km s 1 and 24 < b < 76km s 1 respectively. Since (log N αγ ) provides a more stringent lower limit on b while (log N αβ ) better constrains the upper limit, we thus assume a 1σ range of 31 < b < 46km s 1, and a 2σ range of 24 < b < 55km s 1. It should be noted that Figure 2.3 also shows some overlap between the K α and K γ 16

32 at b < 12km s 1 ; however, this solution is unlikely given the lower limit provided by the K β line. Moreover, b = 12km s 1 implies a maximum temperature (assuming purely thermal motion) of T max = K; such a low temperature is unlikely to produce the observed strong high ionization lines. A similar analysis is not as effective when applied to the strong Ovi LV UV doublet (from the thick disk), since these lines are only slightly saturated: the measured W λ ratio is 0.61 ± 0.04, compared to the expected unsaturated value of When the inferred N OVI is calculated as a function of b for both lines of the Ovi LV doublet, the predicted N OVI values are consistent over b = km s 1 (at the 2σ level; see Figure 2.4). Since the Ovi LV 1032Å line is fully resolved by FUSE ( 15 resolution elements across the line profile) and relatively unblended, its Doppler parameter can be estimated much more accurately using the measured line width and strength. In an unsaturated absorption line, FWHM = 2(ln 2) 1/2 b; however, the measured FWHM increases if the line is saturated. We compensated for this by calculating Voigt profile FWHMs on a grid of N OVI and b, and determining the region consistent with the Ovi LV 1032Å FWHM measurement of 152 ± 7km s 1 (or b = 91 ± 4km s 1 assuming no saturation). When the FWHM derived contour is overlaid on the N OVI b contour inferred from the equivalent width measurement of the LV O vi 1032 Å line, the two regions overlap nearly orthogonally (Figure 2.4) leading to a constraint of b(ovi LV )= 80.6 ± 4.2km s 1. This is 2σ lower than the unsaturated FWHM, once 17

33 again confirming that the Ovi LV is only weakly saturated. At this b the inferred column densities from the two lines of the Ovi LV doublet differ by 1.3σ but this is only a minor discrepancy and likely due to errors introduced by the blending of the 1037Å line; thus, we will only consider results from the more reliable 1032Å line measurement. However, at no value of the Doppler parameter do the 1032Å, 1037Å, and Ovi Kα lines all produce a consistent N OVI measurement; in fact, the Ovi Kα column density is a factor of 4 higher than that inferred from the UV data. This discrepancy is discussed further in Column Densities The Doppler parameters measured for the Ovii (31km s 1 < b < 46km s 1 ) and Ovi LV (b = 80.6 ± 4.2km s 1 ) absorption are inconsistent at the 3σ level, indicating the presence of at least two distinct components: the Galactic thick disk gas traced by broad v 0 Ovi LV absorption, and another lower b phase, possibly of extragalactic origin, traced by the O vii absorption lines. It cannot be assumed a priori that any given line (other than those used to determine b) originates in one phase or another; moreover, the uncertainty in the calculated column density depends not only on the equivalent width error but also the error in b. To take this into account, for each ion the derived column density log N i and its ±1σ limits were averaged over the ±1σ ranges of both measured Doppler parameters. As it turns out, the choice of b does not make a significant difference since all other 18

34 lines (besides the Ovii and Ovi LV absorption) are essentially unsaturated; i.e., the difference in N i calculated with the Ovi LV and Ovii Doppler parameters is small compared to the 1σ error on the equivalent width measurements. Even so, to avoid possible systematic errors, we assumed b = 80.6 ± 4.2km s 1 for those lines likely to originate in the Galactic thick disk (Ovi LV, Ov, and Cv), and b = 31 46km s 1 for all other species. The derived ionic column densities are listed in Table Temperature and Density Constraints At densities such as those found in the Galactic interstellar medium (ISM; n e 1 cm 3 ), photoionization is unimportant because thermal collisions are by far the dominant ionization source. This is also the case for very high temperatures (T > 10 7 K) even at low densities, since the collisional rate is greater than the photoionization rate. However, at the low densities typically found in the intergalactic medium (n e = cm 3 ), photoionization from the diffuse UV/X-ray background begins to play a greater role by enriching the abundances of highly ionized elements at typical WHIM temperatures (log T(K) 5 7) relative to those expected from pure collisional ionization (Nicastro et al. 2002; Mathur et al. 2003). It is thus imperative that the ionizing background be taken into account in order to accurately predict ionic abundances in the WHIM. 19

35 Version of the ionization balance code Cloudy (Ferland 1996) was used to compute collisional plus photoionization hybrid models for the absorbing medium. Relative ionic abundances were computed over a grid of log T(K) = and log n e (cm 3 ) = 7 0, with a step size of 0.1 dex in both log n e and log T. Initially, a rigid scaling of [Z/H] = 1 for all metals was assumed. For the ionizing background we employed the Sternberg et al. (2002) fit to the metagalactic radiation field: J ν = J ν0 ( ν ν 0 ) < ν ν 0 < J ν0 ( ν ν 0 ) 0.46 ν ν 0 > 4 (2.1) where here J ν0 = ergs s 1 cm 2 Hz 1 sr 1 and ν 0 = 13.6 ev. The total flux of ionizing photons is then given by f γ = 4π ν 0 (J ν /hν)dν = photons s 1 cm 2, and the ionization parameter is log U = log(f γ /c) log n e = 6.36 log(n e ) where n e is the electron density in cm 3. Using the ionic abundances calculated with Cloudy, we derived expected abundance ratios for all observed ions at each point in the log n e log T plane. Since any given density and temperature uniquely determines a set of abundance ratios (N a /N b for all ions a and b), the problem can be inverted: any value of N a /N b defines a curve in the log n e log T plane, i.e. a set of temperatures and densities which can produce the measured ratio. When the errors on N a /N b are taken into account, the curves become contours, and the overlap between two or more contours defines the temperatures and densities for which the measured ratios are consistent. 20

36 This is analogous to the method used in to determine Doppler parameters for Ovi LV and Ovii. The most powerful diagnostics are those using ratios between different ions of the same element, since these ratios are independent of the relative metal abundances. Unfortunately the Nvii/Nvi and Nex/Neix upper limits are not stringent enough to place meaningful constraints on the temperature and density. Figure 2.5 shows the log n e log T contours for ratios between the X-ray Ovi Kα, Ovii, and Oviii lines as well as the Ovi HVC /Ovii ratio. The X-ray line ratios are inconsistent with a high density (n e > 10 3 cm 3 ), high temperature (log T > 6.2) medium, and instead converge on a partially photoionized plasma with n e = cm 3 (from the overlap between the Ovi Kα /Ovii and Oviii/Ovii contours) and T = K (from the limits provided by Ovi Kα /Ovii in this density range). These ranges of temperatures and densities are in line with those expected from WHIM gas (Davé et al. 2001). Of course, this is all contingent on the Ovi Kα line being a reliable tracer of N OVI ; this caveat is discussed in detail in 2.4. On the other hand, the Ovi HVC /Ovii ratio overpredicts the temperature by at least an order of magnitude for all values of log n e in order to be consistent with the Oviii/Ovii ratio, the Ovi HVC /Ovii ratio would need to be stronger by a factor of 2.5 (or 3.5σ). It is possible that the HVC is not a physically distinct component, but is instead the result of some systematic error (such as fixed pattern noise or an unexpected anomaly in the Galactic Ovi LV velocity distribution). In this 21

37 case, the Ovi associated with the Ovii and Oviii may be completely blended with the thick disk Ovi LV and thus unmeasurable. Consistency with the Oviii/Ovii ratio (in the collisional ionization regime) requires log(n OVI /N OVII ) 2.5, or roughly 20% of the Galactic UV Ovi absorption. Although it appears that the Ovi HVC as measured cannot originate in the same medium as the Ovii absorption, we suspect that additional atomic physics may be at work here and could in principle reconcile this disagreement (see 2.4.1). While the Ovi Kα /Ovii and Oviii/Ovii ratios provide strong constraints, it is important to consider other ion ratios as well (particularly since the Ovi Kα and 1032Å Ovi column densities disagree). Figure 2.6 shows the log n e log T contours for several different ion ratios, all calculated relative to O vii since the error on N OVII is small. With a rigid metallicity shift relative to solar, the Neix/Ovii Oviii/Ovii, and Ovi Kα /Ovii ratios are not all consistent with each other for any combination of temperature and density; however, the consistency can be improved with adjustments to the [Ne/O] ratio (see 2.4.1). Both the C vi/o vii and N vii/o vii measurements are consistent with a high or low density medium at solar abundances. Limits on the temperature of the Galactic thick disk absorption can be derived in a similar fashion, although it is not the primary focus of this work and there are far fewer measured lines to work with. The most accurately measured line is the Ovi LV ; additionally, Cv and Nvi X-ray lines are measured, and upper limits have 22

38 been determined for O v and N vii. Figure 2.8 shows the temperature constraints derived for this Galactic absorption, assuming pure collisional ionization. The Ov/Ovi LV and Nvii/Nvi upper limits provide metallicity independent constraints of log T > 5.39 and log T < 6.64 respectively. A more stringent upper limit on temperature of log T < 6.03 is provided by the Nvii/Ovi LV ratio, but this is somewhat dependent on [N/O]. Within this range the Cv/Ovi LV ratio provides an even stricter limit of 5.3 < log T < 5.7, but again this depends on [C/O]. At these temperatures the expected O vii column density is at most an order of magnitude less than measured; thus, the Ovi LV, Ovii, and Oviii cannot all originate in the same phase assuming pure collisional ionization (see also Mathur et al. 2003) Discussion Our Chandra and FUSE observations have provided a wealth of data on absorption near the Galaxy, constraining the temperature and density tightly (log T(K) = and n e = cm 3 when the Ovi Kα measurement is included), which are conditions suggestive of the local group intergalactic medium and require supersolar [Ne/O]. Here we first examine the assumptions that have led us to these results ( 2.4.1), and then we discuss their implications for the location of the absorbing gas ( 2.4.2), subject to these caveats. 23

39 Potential Caveats The Ovi Discrepancy The interpretation of the UV and X-ray data are particularly important, since (as Figure 2.5 shows) the combined Oviii/Ovii and Ovi/Ovii ratios can provide tight constraints on the absorber temperature and density simultaneously (see also Figure 5 in Mathur et al. 2003). However, in this case the Ovi column density inferred from the Ovi Kα is a factor of 4 higher than the combined 1032Å low and high velocity components. Since both the X-ray and UV transitions trace the same atomic state, the inferred column densities should match. A similar disagreement has been seen in intrinsic AGN absorption systems (see Krongold et al. 2003; Arav et al. 2003); however, in these cases it is typically attributed to saturation or a velocity dependent covering factor, neither of which is relevant to this z 0 absorption. On the other hand, our Ovi Kα measurement provides a test for these attributions; the local absorption, after all, is likely a dramatically different physical system than an AGN outflow, yet the same conflict arises. A macroscopic explanation does not adequately describe how this discrepancy is seen in both physical systems, so the actual reason may lie in the atomic physics of highly ionized plasmas. For example, some fraction of the Ovi may be excited through collisions or recombination from Ovii, and thus unable to produce 1032Å absorption while still 24

40 absorbing Ovi Kα photons. While a scenario that produces significant depopulation of the Ovi ground state is difficult to envisage in such a low density plasma, we are investigating further the statistical equilibrium of Ovi including photoexcitation and recombination in order to study such effects in more detail. However, it should be emphasized that this is not an isolated case so there must be a physical explanation for the Ovi discrepancy, and the resolution of this paradox is crucial to our understanding of Ovi UV and X-ray absorption and how it relates to the Ovii. There is also the possibility that the line was misidentified as Ovi Kα, and is actually another intervening O vii absorption line at z = This latter explanation is unlikely since no other absorption lines at this redshift are seen in the FUSE or Chandra spectra; additionally, this would require the line to fall exactly on the O vi rest wavelength, which seems like an improbable coincidence. Another possibility is that the theoretical oscillator strength of the Ovi Kα transition is incorrect, but the value given in Pradhan (2000) would need to be low by a factor of 2 4, in sharp contrast to the successful calculations of f lu for inner shell transitions in other ions in the same paper. Nevertheless, due to the discrepancy between the UV and X-ray O vi column density measurements, we present both possibilities: either (a) the Ovi Kα line measures N OVI, or (b) it does not and is thus ignored. 25

41 Absorption Components The Doppler parameter measurements indicate the existence of two distinct components along the Mkn 421 line of sight: one seen in the thick disk Ovi LV 1032Å absorption with b LV = 80.6 ± 4.2km s 1, and the Ovii absorber with b OVII = 31 46km s 1 (1σ limits). The Ovi HVC may represent a third phase (if case (b) above is correct) with b HVC = km s 1 (from the FWHM measurement). This agrees surprisingly well with the O vii b measurement, and is consistent with numerical simulations of the nearby IGM (Kravtsov et al. 2002). However, the extremely low Ovi HVC /Ovii ratio requires a temperature much higher than the upper limit provided by b OVII. In order for the HVC to trace the same gas as O vii (case a), then, the aforementioned atomic physics effects would need to be suppressing Ovi HVC absorption and not the Ovi LV line. Sembach et al. (2003) list mean Doppler parameters for a variety of HVCs, both Galactic and probable Local Group; unfortunately, the dispersion in these values and the errors on b OVII and b HVC measured here are both too large to associate the components presented here to one of their classifications. It is also important to note that our analysis assumes a single phase origin for the included X-ray lines. This assumption is consistent with the data, given the good agreement between the three Ovii lines in the calculated ranges of b and N OVI (Figure 2.3). Even so, if any of the ionic species arises in more than one 26