A Search for Associated O VI Absorption in Intermediate Redshift Galaxies with the Cosmic Origins Spectrograph

Transcription

1 A Search for Associated O VI Absorption in Intermediate Redshift Galaxies with the Cosmic Origins Spectrograph Kathryn Grasha ABSTRACT We present the results of a search for ultraviolet (UV) absorption line spectra toward 69 quasars at < z < detected with the O vi λλ1032, 1038 Å doublet associated with the host galaxy, drawn from the COS-Halos and COS-Dwarfs surveys. We attempt to properly assess the ionization state and metallicity to characterize the basic properties of the cool and photoionized gas residing within the intergalactic medium, as traced with the O vi doublet transition and the feedback processes between galaxy/black hole growth and outflows. We search within ±5000 km s 1 of the systemic redshift of each quasar for associated absorption line systems, finding a total of 88 absorbing systems, where a total of 80 are detected in H i absorption and 63 are detected with the O vi doublet lines (56 absorbing systems exhibit both H i and O vi absorption, 25 absorbing systems exhibit H i but not O vi absorption, and seven absorbing systems exhibit O vi but not H i absorption), where 50 of the associated absorbing systems are also detected in at least one metal line. Using the set ratio of 2:1 for optical depth, we find that 31 out of 61 of our oxygen absorbers (two systems do not show clear 1038 Å lines, and hence we are unable to evaluate the covering factor for those systems) show evidence for partial covering, suggesting that many of our absorbers are located very close to the central quasar embedded within their galaxy. We also find 16 N v λλ1239, 1243 Å doublet in our data, an ion uncommonly associated with intervening absorbers. We employ CLOUDY ionization code to model the parameters of our associated absorbing systems, finding a median metallicity of [Z/H] = 0.3 and a mean ionization parameter of U = 1.2. Our mean metallicity is very high when compared to prior studies investigating the metallicity of intervening O vi absorption systems, however, this is still consistent with what is expected in associated absorbing systems undergoing enhanced radiation fields near the proximity of an AGN. Subject headings: intergalactic medium quasars: absorption lines galaxies: halos ultraviolet: galaxies

2 2 1. Introduction In the void between the stars within a galaxy is the interstellar medium (ISM), a crucial component of a galaxy composed primarily of gas, dust, and metals. The ISM predominantly consists of three phases: the cold neutral medium (CNM) component (T 10 4 K), wellunderstood from galactic studies; a warm, ionized medium (WIM) component, consisting of photoionized, diffuse gas (T K); and a shocked, hot ionized medium (HIM) component (T K). All phases have the capability to extend several hundred kpc from the galaxy (Lanzetta et al. 2005). This repository of the cold and hot phases of the ISM gas, when located at great distances from the galactic center, is usually referred to as the halo gas or the circumgalactic medium (CGM). Believed to be a major source of fuel for star formation, the hot phase of the ISM (T 10 5 K) may be the dominate reservoir of baryonic material at low-redshift (Cen & Ostriker 1999) in combination with the gas that resides between galaxies, known as the intergalactic medium (IGM). Additionally, the CGM moderates galaxy feedback via outflows (Di Matteo et al. 2005; Hopkins & Elvis 2010); studies will reveal connections between the CGM and galaxy/black hole formation as outflows are inevitably associated with star formation processes(pettini et al. 2001; Heckman et al. 2002; Weiner et al. 2009). Due to the relatively low-density of the IGM and CGM, the diffuse, ionized gases that compose the WIM/HIM cannot be detected in emission (e.g., Furlanetto et al. 2004; Bertone et al. 2010) and are primarily detected with ultraviolet (UV) and x-ray absorption line spectroscopy. At low-redshift, highly ionized phases are mostly detected with the O vi λλ , doublet (e.g., Tripp et al. 2008; Thom & Chen 2008; Wakker & Savage 2009; Tilton et al. 2012), the most cosmically abundant metal, extensively studied since the arrival of the Hubble Space Telescope (HST) and the Cosmic Origin Spectrograph (COS; Green et al. 2012). Requiring an ionization energy of 114 ev for production, O vi is the best tracer of hot gas owing to its large oscillator strength, the relatively common abundance of the O 5+ ion, peaking in collisional ionization equilibrium (CIE) at T K (Sutherland & Dopita 1993), and the transition occurring long-ward of the Lyman limit ( 912 Å). The incidence of O vi absorption depends on the physical conditions (e.g., temperature, metallicity, ionization), the location, and the dispersion of the metals within their environment (Oppenheimer & Dáve 2009). Despite the extensive work that has recently been dedicated to the study of the ionized gas within galaxies and the IGM with the O vi doublet (Burles & Tytler 1996; Tumlinson et al. 2005; Stocke et al. 2006; Tripp et al. 2008; Thom & Chen 2008,a; Danforth & Shull 2005; Danforth et al. 2006; Danforth & Shull 2008; Fox et al. 2008; Fox 2011; Howk et al. 2009; Prochaska et al. 2011), where these studies focused on intervening gas systems with

3 3 quasar-galaxy pairs, the location and mechanism giving rise to O vi and the physical origin of the gas is still contested (i.e., whether the gas is bound in dark matter halos, participating in outflows, or being ejected via galactic winds). Research suggests that O vi absorption arises predominately from photoionized systems (Tripp et al. 2008; Thom & Chen 2008a; Howk et al. 2009; Prochaska et al. 2011) due to frequently observed well-aligned profiles of O vi and H i, the general narrowness of accompanying Lyα profiles, and common detection of C iii λ977 lines with the O vi absorber. Other research supports O vi arising from collisionally ionized environments (Fox 2011). This has lead to models suggesting that O vi arises from a combination of both collisionally ionized and photoionized gas (Sembach et al. 2003; Collins et al. 2004) as the observed ionization conditions cannot be described with a single origin. Although UV absorption line spectroscopy is an efficient means of studying the properties and distribution of the CGM, the origin and relevance of the CGM material to galaxy evolution over cosmic time is not fully understood. The distribution of O vi may provide insight to how metals are dispersed and depleted in galaxies, in addition to advancing knowledge on the nature and physical conditions of the hot, diffuse ISM. Because it is common to detect both O vi and H i in absorption in the same kinematic region, it is natural to expect that these ions exist in a multiphase medium (i.e., the hot and warm gas reside together), arising from shocks and/or turbulent mixing of the different phases. Indeed, recent work by Rupke & Veilleux (2013) has found that it is common for ionized and neutral gas phases of an AGN-driven outflow to be cospatial and share similar kinematics. Studies of the CGM will provide insight to the feedback interactions occurring in galaxies and how important the role of feedback has on galactic evolution. Unfortunately, due to observational constrains, outflow processes and the interactions that occur between a galaxy and the IGM, such as gas accretion or galactic winds, are usually limited to regions very close to the galaxy. Fortunately, quasar absorption line studies are excellent probes of the gas located near an active galactic nucleus (AGN) and also have the ability to explore the physical conditions of IGM gas located far from a galaxy, thus allowing us to build upon our knowledge of galactic evolution and processes that add and/or remove gas from galaxies. In this respect, the connection between quasar absorption line systems and their environment provide a crucial understanding of the properties and evolution of galaxies. The classification of AGN absorbers are normally categorized according to their H i Lyα column densities: damped Lyα absorbers (DLA; Wolfe et al. 2005) show H i columns of N HI cm 2, Lyman limit systems (LLS; Tytler 1982) exhibit N HI < cm 2, and the Lyα forest (Rauch 1998) exhibits columns of N HI < cm 2. The LLS likely represent the interface between the highly ionized Lyα forest the most-common H i absorbing system we find in our study and the dense, neutral DLAs (Lehner et al. 2009).

4 4 The origins and distribution of absorbing systems within the galactic environments around the host quasar, in addition to furthering knowledge of the properties of the ionized medium, motivates this study. Studies like this help shed light on the properties of the environments near quasars, where it is generally believed that bright quasars reside in luminous galaxies (Jahnke & Wisotzki 2003), embedded in over-dense environments with massive dark matter halos (Serber et al. 2006). Galactic feedback can be regulated with black hole accretion and ejection processes (Heckman et al. 1991a,b; Haiman & Rees 2001), capable of driving material from the central region of a galaxy with outflows reaching velocities that exceed km s 1, enriching the chemical composition of the IGM (Hopkins & Elvis 2010; Barai et al. 2011; Cavaliere et al. 2002). These systems, characterized as broad absorption line (BAL) systems, are common in the universe, however, knowledge of these far-reaching outflows on galactic evolution still needs further investigation. Absorption line spectroscopy is the best way to study the environment of the host galaxy in the vicinity of the quasar, however, it is often challenging to distinguish if the gas associated with the host quasar originates from the accretion disk on sub-parsec scales (Elvis 2000; Krongold et al. 2007), or from the halo gas, out to kpc-scales (Kinkhabwala et al. 2002). In this paper, we focus on O vi absorbers associated with their host galaxy to constrain how the near-proximity of an AGN influences the halo gas within the galaxy; we expect the physical condition that gives rise to these associated absorption systems will greatly differ in properties from intervening absorption systems that do not interact with an AGN. What is the nature, origin, and physical conditions of these associated O vi systems? Are these collisionally ionized clouds associated with separate cooler, photoionized clouds of H i absorption, or arising in systems with a combination of both processes? We set out to further our understanding on the nature of highly ionized absorbers in our redshift range of < z < Despite the limitations of the pencil-beam measurements along the line of sight to the background source of absorption line surveys, we believe that all of our absorbers lie within their host galaxy, and quite possibly, have a close proximity to the central quasar. Because of this reason, we do not calculate any cosmological quantities from our data owing to the effects of enhanced radiation fields on associated absorption systems. Throughout this paper, we adopt a ΛCMD concordance cosmology model of Ω m = 0.27, Ω Λ = 0.73, and H = 70 km s 1 Mpc 1 (Komatsu et al. 2011). All numbers taken from the literature are re-calculated (if necessary) to this cosmological model. 2. Sample Selection The galaxies in this study are drawn from the COS-Halos (Tumlinson et al., in prep) and COS-Dwarfs (Bordoloi et al., in prep) surveys, with the primary motivation of studying

5 5 the CGM and its contribution to galactic formation. We are specifically focused on the properties of gas in the immediate vicinity of the target quasars, e.g., gas ejected by the quasar, gas falling in to feed the black hole, or the ambient gas in the host galaxy. TheCOS-Halossurveyiscomposedof39quasarsintheredshiftrange0.229 < z < 0.873, while the COS-Dwarfs survey has 41 quasars in the redshift range < z < The sources in both samples were selected based on quasar-absorber pairs, where each quasar has a known foreground absorber with an impact parameter of kpc. While the COS- Halos and COS-Dwarfs surveys are designed to investigate how both galaxies acquire their gas and how the gas is returned to the IGM within the foreground galaxies, we utilize these surveys to investigate the same science goals, but with IGM gas that interacts with the quasars themselves. Our survey has the benefit of being a blind survey for O vi absorption as the quasar sight-lines were selected based on the properties of foreground galaxies, not on the properties of the quasars themselves. The only criterion placed on the selection of the quasars was that they are brighter than magnitude V We select a total of 69 sight lines (37 from the COS-Dwarfs survey, 32 from COS-Halos) in the redshift range < z < 0.746, meeting our requirement of observing O vi at the systemicredshiftofthequasarinourwavelengthwindowof Å.Weexcludeeleven sight lines as the redshift of the quasar places the OVI doublet outside of our wavelength coverage. 3. Results 3.1. Line Identification We employ two methods of identifying absorption line systems: Identifying O vi systems irregardless of any other absorber present and identifying H i lines without placing constraints on the presence of O vi absorption. For our first method, each sight line is inspected in the velocity range ±5000 km s 1 from the systemic redshift of the source for the O vi λλ1032, 1038 doublet absorption lines, present at the fiducial line strength ratio of 2:1 expected from the oscillator strengths of the two transitions. If the doublet feature is found, we label these as O vi absorption provided that the λ1032 absorption line has an equivalent width W 0 > 3σ. Once we have determined the presence of the O vi doublet, we search at the same velocity offset for each sight line for affiliated H i Lyman series and other metals within our wavelength coverage at the same range of ±5000 km s 1. Table 1 lists all the metal lines we found with our line-identification procedure. Observationally, if the velocity separation between two systems in a given sight line are large, this as an indication that the absorbers are not bound to each other. While it is true

6 6 that two systems can be unbound but have a small velocity difference along the line of sight (or vice-verse), if the normalized flux level rises to unity between two components (i.e., there is no overlap in velocity space), we categorize these absorbing components as unique systems. Throughout this paper, we define a system such that they arise from unique absorption systems. WedetectassociatedO viabsorptionin49ofoursourceswithatotalof63kinematicallyunique O vi absorption systems in our survey. Table 2 lists all of our sources along with central velocity of each absorption system detected in O vi. For complex, multi-component absorption systems, the velocity is defined to be the central velocity for the strongest individual component. Of our 63 O vi absorption systems, seven are not detected in H i absorption at the same velocity. We believe that these O vi absorbers lacking evidence for H i lines originate in relatively over-dense regions situated close to the central quasar. The possibility of weak Lyα H i lines in O vi absorption systems is why we do not first require detection of Lyα, using that redshift to then identify the O vi doublet. If we detect moderately strong Lyα but Lyβ is not present, we assume that the Lyα line measurements are not badly saturated. The majority of our systems are only detected in O vi and H i absorption. We recognize that due to the nature of determining the presence of an O vi absorber in our sample, we may miss detections of systems that either have undetected λ1038 profiles due to low S/N and/or systems where either the λ1032 or λ1038 absorption lines are unobservable due to blending with unassociated features. Fortunately, the redshift regime of our survey occurs in a sweet spot, where there is very little contamination due to Lyα forest lines. We note that there are two systems in our sample that are not detected with both O vi transitions due to blending at the same wavelength of the O vi λ1038 line. For these systems, we are confident in the O vi λ1032 detection as there is associated H i absorption at the same wavelength. For our second line identification technique, we identify all possible H i absorbing systems without accompanying O vi absorption. For sources where Lyα is outside of our observing window (z > 0.48), we require at least two H i lines to declare the presence of an absorbing system. All in all, we find 88 kinematically-unique absorption systems 80 detected in H i absorption and 63 detected in O vi absorption in 49 sight lines. The breakdown is as follows: 56 absorbing systems exhibit both H i and O vi absorption, 25 absorbing systems exhibit H i but no O vi absorption, and seven absorbing systems exhibit O vi but no corresponding H i absorption. Additionally, 50 of the 88 absorption systems are also detected in at least one other metal ion transition. Figures 1 22 show all absorbing lines detected in our survey, listed by name where multiple absorption systems detected in one sight line are listed in parenthesis, for ease of comparison to Tables 2 and 3. Figure 23 shows

7 7 the sight line of the richest object in our survey with nine O vi systems detected. Figure 24 shows the distribution of the absorbers in our survey as a function of redshift, broken down into the number of absorbers detected per sight line. Excluding two absorbers, J at a velocity displacement with respect to the quasar of v abs,ovi = 4016 km s 1 and J at v abs,ovi = 1843 km s 1 all of our O vi absorbers are less than v abs,ovi = 586 km s 1 and 73% (46/63) have velocity outflows of v abs < v quasar. The large number of blueshifted absorption features indicate that we are looking at gas outflows, where the two highly redshifted systems of J and J are most likely inflows onto the accretion disk. It is important to note that absorbers detected in velocity space near the quasar redshift does not correspond to a spatial location with respect to the quasar. We search a maximum displacement velocity δv = v quasar v abs, where we make the cut-off that associated (intrinsic to the galaxy) absorbers are located at δv < 5000 km s 1 and all intervening absorbers (unassociated with the galaxy) are located at δv > 5000 km s 1. The cut-off in velocity space supports the different physical conditions observed in intrinsic compared to intervening absorbers as proximate absorbers can have time-varability of absorption (Narayanan et al. 2004), partial coverage of the background continuum source (Hamann & Ferland 1999), super-solar metallicities (Ganguly et al. 2006), and detection of highly excited ionic states (Petitjean & Srianand 1999; Fechner & Richter 2009). Employing a velocity cut-off to differentiate between intervening and intrinsic absorbers can introduce two problems: quasarejected intrinsic systems can appear at higher velocity outflows (Nestor et al. 2008) and some intervening systems can appear at lower velocity outflows (Sembach et al. 2004). We make the assumption that all absorbers detected within δv < 5000 km s 1 are intrinsic absorbers and we do not focus on studying the transition between intrinsic and intervening absorbers in this paper Column Densities We measure column densities for all absorption lines within 5000 km s 1 of z quasar using both the apparent optical depth (Sembach & Savage 1992) and Voigt profile fitting with the measured COS line-spread function. Multiple transitions of the same species (e.g., O vi λλ1032, 1038) are fit simultaneously such that multi-component profiles yield N, b, and v values that are identical for each species within a given system. We allow the components of different species to vary in their central velocity and b-parameter with respect to other transitions in a system as it is the measured column densities, line widths, and relative velocities that diagnose the physical conditions exhibited by the observed absorption systems. For fully resolved, optically thin lines that reside on the linear portion of the curve of

8 8 growth (COG) plot (see the Lyα line in J (1) of Figure 1), the optical depth per unit velocity is given as τ(v) = πe2 m e c fλ 0N(v), (1) where f is the oscillator strength of the transition, λ 0 is the transition wavelength (Å), and N(v) is the column density per unit velocity (cm 2 ). This method makes no a priori assumption about the functional form of the velocity distribution. However, the true optical depth measured is convolved with the instrumental spread function of the telescope φ I. This instrumental blurring of the true optical depth is referred to as the apparent optical depth, calculated as ( ) Iobs (v) τ a (v) = ln = ln[e τ(v) φ I (v)], (2) I 0 (v) where I 0 is the continuum intensity and I obs is the observed intensity of the line. The apparent column density (cm 2 (km s 1 ) 1 ) N a (v) profile is calculated from the apparent optical depth profile as N a (v) = m ec πe 2 τ a (v) fλ 0 = τ a(v) λ 0 f, (3) where the instrumentally-smeared total apparent column density is the integral over the entire velocity range, N a = N a (v)dv, dependent only on the resolution of the instrument and the shape of the line. When comparing the N a profiles of a doublet feature, a resolved system will exhibit similar N a (v) profiles. In contrast, when unresolved saturation features occur, the N a (v) profile of a stronger doublet line will be smaller than the N a (v) profile of a weaker line (see Figure 25 for an example of a resolved and unresolved system). For systems that are optically thin (τ(v) << 1) and the functional spread φ I is fully resolved [i.e., FWHM(line) >> FWHM(φ I )], the apparent integrated column density equals the total column density from component profile fitting. For lines that have moderate saturation, thus residing on the saturated portion of the COGplot(seetheLyαlineinJ (2)ofFigure1),findingthecolumndensityismore difficult as small uncertainties in the Doppler parameter b result in significant uncertainties in the estimated column density. For these moderately saturated lines, the central optical depth is given as τ = π1/2 e 2 fλ 0 m e c N b, (4) where b is the velocity dispersion (km s 1 ). In cases of strong saturation, the column density is better described by Voigt profile fitting. However, if all lines being fitted are badly sat-

9 9 urated, even profile fits will result in huge uncertainties. In such situation, we report lower limits to the derived column density. For H i column densities, the highest Lyman transition that is not saturated is used for the total H i column density measurement. Nearly half of our systems with Lyman series exhibit moderate saturation, meaning that direct integration of the line profiles will underestimate the total column density. For these cases, the column densities are derived from profile fitting as long as a couple of the lines are not strongly saturated. To perform this, we simultaneously fit multiple Lyman series profiles. By using the component information in the higher transition lines, the underlying component structure is more accurately determined, as they are usually only seen in these weaker lines. We do realize that there is an inherit uncertainty to the number of components necessary to adequately describe a line. Therefore, for all profile fitting, we fit with the minimum number of components necessary to obtain a satisfactory fit. For cases of non-detections where there is no absorption detected in the chosen velocity range (see the O vi λ1032 line in J (1) of Figure 1), we estimate the 3σ upper limit to the column density using the uncertainty in the equivalent width, acquired with direct integration techniques integrated over the same velocity range as the corresponding Lyα profile. If we assume that the non-detection line resides on the linear part of the curve of growth, the column density (cm 2 ) upper limit is given as, N 3σ < W λ,3σ. (5) λ 2 0f Table 3 lists the total column density acquired by direct integration, where no individual component is separated. Table 4 lists the individual component properties from Voigt profile fitting Covering Factor C(v) A valuable indicator for determining the location of gas in associated absorbers is the covering factor, C(v). Defined such that 1 C(v) indicates the fraction of the photons from the quasar that are not intercepted by the absorbing cloud. However, estimating C(v) for a given ion transition is non-trivial, and in practice, can only be determined with unblended doublet systems. When the strength of a doublet system is expected to have an optical depth ratio of 2 to 1 (e.g., O vi), we can derive (Hamann & Ferland 1999; Yuan et al. 2002) the partial

10 10 covering factor as C(v) = IR 2 2I R+1, I B 2I R +1 I R > I B I 2 R 1, I B < IR 2 1 I R, I B I R where I R and I B are the normalized intensities of the red (here, O vi λ1038) and blue (here, O vi λ1032) doublet lines, respectively, and C(v) is the total cloud area covering the emission source. For sources that are suspect of partial covering effects, the optical depth of the red doublet line is corrected for partial covering as ( ) IR (v) I B (v) τ corr,r (v) = ln. (6) 1 I R (v) We use the O vi doublet to compute the covering factor and we fix the ratio of the optical depths of the blue and red lines at 2:1 with the assumption that the lines are fully resolved and can be described with Gaussian profiles. When dealing with cases of partial coverage, we do not use the integrated column density in the reported tables, determined in the previous section; instead, we compute the corrected integrated apparent column density N a,corr (cm 2 ) of the ion (assuming full resolution and Gaussian distributions of optical depth) by integrating the corrected optical depth across the absorption profile as N a,corr = m ec πe 2 λ 0 f ( ) τ corr (v) dv = τ0,corr, (7) λ 0 f where τ 0,corr is the central optical depth of the absorption line, corrected for partial covering. Overestimating the covering factor will give rise to a measured apparent column density that is lower than the true column density. Therefore, applying the correction factor will increase our total N a values. For cases that have full coverage C(v) = 1, the corrected optical depth (eq. 6) reduces to eq. 2 and N a,corr N a. Figure 25 compares the apparent column density profiles for both the O vi doublet in a system not affected by partial covering and in a system that does show partial covering. Unfortunately, saturation and partial covering are degenerate and it is hard to distinguish between the two effects outside of Voigt profile fitting. Partial covering will result in missing flux, therefore if a system is subject to partial covering, the flux in the stronger line will be underestimated while the weaker line will be overestimated, which results in poor profile fits. Figure 26 shows an example case of how we determine when a line has partial covering from profile fitting.

11 Metallicity The absolute metallicity abundances for each source is measured using solar reference abundances ([X/H] ) from Asplund et al. (2009) as [X/H] = log[n X /N HI ]+log[f HI /f X ] [X/H], (8) where f is the ionization fraction for a given ion, which depends on the strength of the surrounding radiation field and we have assumed solar relative elemental abundances. Two ionization states of the same species is enough to pin down the metallicity and ionization parameter of a system. However, when two ions of a single element are not available, it is possible to use other elements to arrive at an estimation of the metallicity. However, it will add uncertainty because relative elemental abundances have to be considered, which we do not know. In general, we avoid using nitrogen for our metallicity measurements, as N is often be under-abundant due to nucleosynthesis effects in low-metallicity gas (Henry et al. 2000). All of our ionization fractions were constrained with CLOUDY modeling (Section 4.2). 4. Analysis 4.1. Gas Temperature of the Absorbers If we assume that two well-aligned species arise from the gas phase, it is possible to derive an estimate for the gas temperature from the Doppler parameters of the absorption profiles from Voigt profile fitting (Chen & Prochaska 2000; Tripp & Savage 2000). Using the assumption that line broadening results entirely from thermal motions, the Doppler parameter b (km s 1 ) is related to the gas temperature via b 2 = 2kT/m, (9) where T is the temperature (K), and m is the mass. However, we typically expect there to be gas turbulence and/or unresolved components that also contribute to the line width. Including these non-thermal components will only further to lower the estimated temperature, essentially making our temperature estimates upper limits. We can express this temperature relation using the atomic mass number (A) of both elements as b 2 = b 2 nt +2kT/m = b 2 nt +(0.129) 2 T/A, (10) with the assumption that any broadening owing to non-thermal motions (turbulence) can be described with Gaussian profiles and requiring that both species have different A values.

12 12 As an example, we demonstrate these effects for the sight line of J The temperature upper limits from equation 9 and the Doppler parameters of b HI = 34(7) km s 1 and b OVI = 16(2) km s 1, from which we derive T HI < 6.9(1.4) 10 4 K and T OVI < 2.4(3) 10 5 K. SincethevelocitycentroidsoftheO viandh iprofilesarequiteclose( v < 10km s 1 ), itis reasonable to believe that the O vi and H i absorption lines originate in the same gas cloud. Using this assumption, we can derive the non-thermal components and temperature from equation 10 using the values of b HI and b OVI : T = 5.7(1.3) 10 4 K and b nt = 14(3) km s 1. If we compare this with the temperature from our best-fit CLOUDY model to this absorption system (T cloudy = K), we find that our answer is very reasonable for photoionized conditions (see next section for detailed information about CLOUDY). Our answers are not identical owing to the large amount of simplicity and uncertainty within CLOUDY models, which excludes any physical conditions outside of photoionization equilibrium. Such low-temperatures favor photoionization conditions (T 10 4 K), however, it is important to note that there is a huge amount of uncertainty in the b-parameter. Changing the number of fitted components to an absorption profile can lead to substantially different profile fits through a different mix of broad and narrow components. Additionally, equation 10 assumes equilibrium and does not take into consideration the possibility that the absorption arises in non-equilibrium collisionally ionized gas (T > K), which cools faster than recombination. The Doppler parameters for O vi and H i are shown in Figure 27 along with a limiting case where b nt = Physical Conditions of the Absorbers The structure of photoionized gas is primarily determined by the shape of the incident continuum flux and the ionization parameter. We investigate the physical conditions of each absorber with photoionization models from the CLOUDY (v10.0; Ferland et al. 1998) ionization code, where we find both the metallicity of each absorption system and the ionizing parameter, U, from fitting the models to our observations with the models. In all of our models, we use a standard AGN continuum flux model, solar relative abundances, and the galactic gas is considered to be a uniform slab that is illuminated from one side by photoionizing radiation. The standard AGN model produces a multi-component continuum similar to that observed in typical AGN, similar to the continuum from Mathews & Madau (1987). The models do not heavily depend on the value of the hydrogen volume density n H. Using the observed column densities of each source for a specific kinematic component (i.e., each individual absorbing component), we use the total hydrogen column density N HI

13 13 from Voigt profile fitting and determine the ionization parameter (U) for each source as U = Q 4πr 2 n H c, (11) where r is distance from the ionizing source center to the illuminated face of our gas cloud, Q is the number of H-ionizing photons per second, and n H is the total number density of hydrogen (10 9 cm 3 ). For our whole sample, a grid of photoionization models range from 4 < logu < 2 with steps of logu = 0.1 dex and metallicity values cover the range of 2 < [Z/H] < 1 at relative solar abundances with step size logz = 0.1 dex. The best-fit CLOUDY model to a set of observations is done using the χ 2 statistic for combinations of metallicity Z and ionization parameter U as ( ) 2 Ni,obs N i,model (Z,U), (12) χ 2 (Z,U) = i σ(n i,obs ) where N i,obs are the measured column densities for each specific kinematic component for a source, σ(n i,obs ) is the uncertainty in the measured column density, N i,model are the column densities for a specific ion at different metallicity and ionization values, and we sum over the total number of ions i for each kinematic component. The estimated 1σ error bar corresponds to a 1σ confidence interval for each parameter, obtained by finding the optimized parameter values that satisfy χ 2 = 1 for each best-fit. For our CLOUDY fitting procedure, we assume that for a given absorption system, all of the absorbers arise from the same gas cloud. We investigate the systematic errors from our modeling resulting from this assumption by also placing more conservative CLOUDY estimates on our data. This is done through an assumption that all similar ionized states arise from a single gas cloud that is disjoint from a different gas cloudy at lower ionization state, where both clouds lie at the same velocity (i.e., O vi and N v absorption profiles are distinct from C iii absorption observed at the same velocity). For these more conservative estimates, we find that the median metallicity changes from [Z/H] = 0.3 to [Z/H] = 0 (conservative), where the spread is the same. Figure 28 shows some example CLOUDY fits to select absorption systems. Figures 29 shows the distribution of [Z/H] and U for our detections, upper limits, and lower limits for both the conservative case (not assuming all absorption features arise from a single origin) and the non-conservative case (all absorption features have single origin). Figures 30 shows the same, except broken down by systems exhibiting partial covering versus those that are coving factors of unity. For O vi systems not detected in H i, we place limits on the CLOUDY [Z/H] and U measurements. The proper statistical analysis of detections mixed with limits is dealt with in Section 4.7.

14 N v Occurrance ThepresenceoftheN vλλ1239, 1243Åinfoundinsixteensightlineswithinoursample, providing an independent diagnostic on both the physical conditions and the location of the absorbers. N v requires a large ionization potential of 77.5 ev, created only by the hard radiation of a close-by quasar (Fechner & Richter 2009), therefore, observations of N v are rarely observed in non-associated absorbers. Intervening absorbers located in the IGM are typically linked a softer ionizing background and will lack N v features. The strong N v detections, coupled with the fact that N v is rarely observed in intervening absorbers, seen in 16 of our sight lines, support the location of these absorbers being in the vicinity of a foreground quasar, undergoing a higher level of ionization. Figure 31 shows the metallicity and U parameter distributions of the absorption systems, both those detected and not, in N v absorption. No statistically significant differences are found between the systems detected and those not in N v absorption Absorber Classification: Single-Phase and Multiphase Absorbers In order to investigate the physical conditions of the relatively cool (T 10 4 K), photoionized gas arising from the O vi absorption systems, we classify the absorbing systems in our sample according to the velocity offset between the O vi and H i. This is done by segregating simple (single-phase) absorption systems from those exhibiting a more complex (multiphase) structure. Many of our O vi absorption systems are complex, and quite possibly, multiphase absorbers. Each unique kinematic velocity system is designated a central velocity; for multicomponent systems, this is the velocity of the strongest O vi component of the system. We determine the velocity alignment of the system as v = v HI v OVI with an uncertainty in the offset given as σ( v) = σ 2 OVI +σ2 HI +σ2 wave, where σ OVI and σ HI are the uncertainties in the central velocity of the O vi and H i absorption profiles and σ wave is the uncertainty in the wavelength calibration across the observed wavelength range. We take σ wave = 8 km s 1, the average value of the uncertainty in the wavelength calibration for COS, and define wellaligned profiles when v < 2σ( v). The similar velocity structure for systems that are well-aligned suggest that the O vi and H i gases are indeed, cospatial and arising from a single-phase gas cloud. Our assumption a single gas phase that we use in our CLOUDY models works well for these cases. For systems that have different velocity structure for the O vi and H i features suggests these gases may not be arising from a single origin and that our assumption of a single gas phase is likely invalid. It is this reason that we performed our CLOUDY modeling twice, once assuming

15 15 that all absorption features in a system arise from a single cloud, and again, where we are more conservative and do not force all absorption features to originate from a singlephase gas cloud. Figure 32 shows the velocity offsets between H i and O vi for our 56 systems detected in both H i and O vi absorption. 33 out of our total of 56 absorbers have well-aligned O vi and H i absorption profiles, according to our definition above. The similar velocity structure suggests a physical origin of the O vi absorbers, where the low-temperatures are consistent with the absorbing gas arising from photoionized processes. This is consistent with our temperature results from Section 4.1, where we showed that using the b-parameters from H i and O vi absorption profiles result in temperatures that are consistent with the gas arising from photoionization. Figure 32 shows the histogram for the absorber pair as a function of velocity difference between O vi and H i, where it can be easily seen that a great deal of our cases do exhibit well-aligned absorption profiles The N HI /N OVI Ratio The ratio of the N HI /N OVI, sensitive to the ionization conditions within an absorbing system, can be used to quantify the different physical conditions. Changes in the temperature, metallicity, ionizing radiation field, and gas density can lead to substantial changes in the N HI /N OVI ratio. In multiphase absorption systems, lower ionization phases can increase the strength of the H i absorption profile without an increase in the O vi strength if the input energy is less than the 114 ev necessary to create O 5+. In general, a low ratio would indicate a high-density region, resulting from high-metallicity gas and the occurrence of strong ionization. Figures 33, 34, 35, 36, 37, show the N HI /N OVI ratio as a function of the velocity offset of the absorbing system from the central quasar, velocity difference between the O vi and H i component, redshift, metallicity, and width of the O vi λ1032 absorption feature Final Analysis Information The nature of a project this large means that there is a lot of information to absorb and digest. The rest of Figures all show important information about the distribution of [Z/H], N OVI, U, and N HI as a function the velocity offset of the absorber from the quasar (Figure 33), velocity difference between the O vi and H i absorbing line (Figure 34), redshift (Figure 35), metallicity [Z/H] (Figure 36), and FWHM of the O vi λ1032 line (Figure 37). Each figure provides a new and unique analysis of our absorbers, where we interpret the statistical significance of any correlations present in the next Section.

16 Survival Analysis In order to transform our line measurements into statistical descriptions and distributions, we must properly account for all measurements, detections and non-detections. To best deal with the mixture of measurements and upper/lower limits, we apply methods of survival analysis to our data, statistics designed to properly correct for censored data sets. In our study, we use the Kaplan-Meier product limit, a single-variate survival statistic that provides a non-parametric maximum likelihood estimate of a distribution directly from an observed data set. For our data set, which has a combination of detections, upper limits, and lower limits, we will use the following notation, following that of Simcoe et al.(2004): the term measurement describes the combined data set of detections, upper, and lower limits, Z true represents the actual measured metallicity for a detected line such that Z true = Z i, while for upper limits, Z true < Z i, and for lower limits, Z true > Z i. For a collection of N data points with a given metallicity Z (either a detection or a upper/lower limit), we sort all the metallicity values such that Z i Z i+1. A cumulative probability distribution P(Z > Z i ) is built, describing the fraction of absorbers that have a metallicity above a given threshold (see Figure 38). At the maximum of the distribution (Z + ), every measurement is at a lower metallicity (i.e., P(Z < Z + ) = 1). Using standard notation, this is written as P(Z Z + ) = 1 P(Z < Z + ) = 0. (13) All subsequent values of P(Z Z i ) are calculated at each value i by stepping down from the N th measurement. We can calculate the conditional probability as P(Z Z N ) = 1 P(Z < Z N ) = 1 P [Z<ZN Z<Z + ]P(Z < Z + ) = 1 P [Z<ZN Z<Z + ], (14) where P [Z<ZN Z<Z + ] is the conditional probability that Z < Z N given that Z < Z +. P(Z Z N 1 ) becomes P(Z Z N 1 ) = 1 P [Z<ZN 1 Z<Z N ]P [Z<ZN Z<Z + ]P(Z < Z + ). (15) For an individual metallicity value i, this conditional probability becomes, P(Z Z i ) = 1 N P [Z<Zj Z<Z j+1 ]. (16) For a sample composed entirely of detections, the conditional probability at each value i can be written as P [Z<Zj Z<Z j+1 ] = number of detections with Z < Z j (17) number of detections with Z < Z j+1 j=i

17 17 = n (Z<Z j+1 ) n (Z=Zj ). n (Z<Zj+1 ) The combination of both upper and lower limits in our data makes the quantities n (Z<Zj+1 ) and n (Z=Zj ) not uniquely known. Because limits do not provide any information on the relation between Z true, Z j, and Z j+1, an ambiguity is introduced into the sample. The Kaplan-Meier product circumvents this ambiguity by retaining knowledge of all the limit values and effectively ignoring those values in the construction of the conditional probability P [Z<Zj Z<Z j+1 ]. The conditional probability, taking into account non-detections, can be calculated as P [Z<Zj Z<Z j+1 ] = number of measurements that must have Z true < Z j. (18) j We can combine the results for detections and non-detections, writing the conditional probability as { 1 Zj = limit P [Z<Zj Z<Z j+1 ] = j n (Z=Zj ) (19) Z j = detection. j In the final product, the Kaplan-Meier product estimate is a piecewise function, remaining constant for censored data (upper/lower limits), and only jumping at the metallicity values of detections. We assume that all limits are independent of each other and that the censoring is random. While it is true that very high-metallicity values are less likely to be censored compared to their low-metallicity counterparts (thus causing censoring of metal abundances to not be truly random), the systems in our sample were not selected with a metallicity criterion, and thus are expected to be unbiased. Figure 38 shows the Kaplan-Meier distribution for our sample of metallicities, shown as the entire sample, and divided into systems that do (not) exhibit partial covering. We take care to separate the sample by the covering factor as we believe systems fully covering the background emission source are further from the central black hole compared to those systems that do not fully cover their background source. These full-covering systems likely interact less with their black hole and may exhibit different properties than those systems with partial covering factors. For the case where we have contamination from false O vi detections, the KM-distribution represents an upper bound to the true metallicity distribution. As can be seen in Figure 38, the censored data are mixed in with the measurements over nearly the entire range covered by our sample. We can apply the results of this Kaplan-Meier product estimate to a Kolmogorov- Smirnov test (KS-test), where the KS-test is a statistical description whether or not two distributions are drawn from the same parent population. Using our metallicity distribution

18 18 of systems exhibiting partial covering compared to those that do not, we conclude that the two samples are not drawn from the same population at a significance value less than 1σ. In other words, the metallicity distribution for absorption systems that exhibit full covering is nearly identical to the metallicity distribution for the systems that show partial covering. One can conclude from this that all of our absorption systems are indeed intrinsic absorption systems, lying close to the vicinity of the quasar. Future work is needed to apply the KS-test to all the variables in this paper. 5. Discussion Given the information about our derived temperatures from the Doppler width b for our absorption systems, can we confidently claim whether not our O vi absorbers favor photoionization or collisional ionization? In extragalactic O vi absorbers, it is generally believed that intrinsic absorption systems are undergoing an UV ionizing radiation field that is substantially harder due to the close proximity of the central quasar compared to intervening absorption systems. Additionally, the absorption may also arise in very lowdensity gas with long path lengths. Both of these cases indicate photoionization as a more viable solution for the occurrence of O 5+ in our associated absorption systems and this is supportedwithourderivedtemperaturesfromlinefitting, typicallyt < 10 5 K.Thisassumes that the H i and O vi absorption arise from the same gas cloud, with the line width arising from thermal broadening. To complement our CLOUDY modeling, we compare our results with the theoretical modeling of time-dependent metal-enrichment of halo gas from a nearby AGN of Oppenheimer & Schaye (2013). In their models, they find that diffuse halo gas in thermal equilibrium, when exposed to near-by AGN radiation (i.e., the AGN turns on ), will quickly photoionize, reducing the observed N HI almost instantaneously (H i becomes ionized). When the AGN radiation is turned off, the hydrogen quickly recombines (1 20 Myr) while the renormalization of metal abundances occurs over a much longer timescale. Referred to as a fossil zone, the region immediately near the proximity of an AGN will exhibit enhanced high-ionization states with a reduction in low-ionization states. We can use this fossil zone model to explain associated absorption systems that show unusually strong, high-ion associated metal systems, as such systems may either be located within these fossil zones or near an AGN that has very recently turned on. Fossil proximity zones are especially important for high-redshift systems (z 2), as it is possible that a majority of metal-enriched systems are non-equilibrium fossil zones; using equilibrium coding to model such systems would result in incorrect inferred physical conditions of the absorber systems. Our findings of O vi absorption with little-to-no H i absorption supports the fossil proximity zone model as we observe enhanced O vi lines and weakened H i lines and other low-ion metals, consistent

19 19 with the the idea of metals taking longer to reach equilibrium after an AGN has turned off. In studies of absorption systems, one would expect to find a change in N HI, N OVI as the velocity of the absorption approaches the redshift of the quasar as the ionizing field would increase as one approaches the quasar (the proximity effect). We find no dependence of N OVI with proximity, which may cause one to question if photoionization from the background quasar is the main driver of the observed O vi. However, since we are only studying intrinsic systemswithin5000km s 1 fromthequasar,wenaturallywouldnotexpecttoseethischange as (ideally) all of our gas systems are located very close to the AGN and no absorber in our sample is free from the ionizing effects of the AGN. This is further supported by the large number of absorption systems exhibiting partial covering (31 out of 61 O vi absorbers) and the large number of sight lines observed in N v absorption, a species uncommonly observed in intervening systems. Along the same line, it is reasonable to expect Lyα absorbers to exhibit lower metal abundances at greater distances from their host galaxies as the UV ionizing radiation field decreases with distance. We have found a small number of associated O vi absorbers without corresponding H i absorption, which suggests an overdense environment, devoid of cold, neutral hydrogen (super ionized) accompanied with high-metallicity. While there are only seven of these systems, they are all located at velocity outflows greater than 1000 km s 1 from the redshift of the central quasar. We are unable to draw any correlations due to the small number of these systems in our sample. 6. Summary and Conclusion We present the results of 63 associated O vi absorption systems in 49 sight lines, selected from high-resolution HST COS-Halos and COS-Dwarfs surveys in the redshift range < z < We adopt a blind search technique for identification for O vi λλ1032, 1038 doublet lines, relying only on the presence of the O vi doublet and the ratio of doublet line strength, independent of the presence of other transitions. In the redshift range of each detected O vi absorber, we have also searched for the presence of H i absorption and other ionized species, where we find 80 H i absorption systems (25 have no accompanying O vi absorption), for a total of 88 absorption line systems. 50 absorbing systems also exhibit at least one other metal ion. Our study focuses on O vi absorption as it plays an important role in the baryon and metal budgets of gas, provides observational windows on intergalactic metal enrichment, and traces energetic galaxy and IGM interactions. For each absorption system, we distinguish between absorption features that are welldescribed as single-phase absorbers and absorbers that arise from multiphase absorption. We interpret multiphase absorption profiles as physically distinct gas structures within an