Caught in the Act: A Metal-rich High-velocity Cloud in the Inner Galaxy

We characterize the chemical and physical conditions in an outflowing high-velocity cloud (HVC) in the inner Galaxy. We report a supersolar metallicity of [O/H] = +0.36 ± 0.12 for the HVC at v LSR = 125.6 km s−1 toward the star HD 156359 (l = 328.°7, b = −14.°5, d = 9 kpc, z = −2.3 kpc). Using archival observations from the Far-Ultraviolet Spectroscopic Explorer (FUSE), the Hubble Space Telescope Imaging Spectrograph, and the European Southern Observatory Fiber-fed Extended Range Optical Spectrograph we measure high-velocity absorption in H i, O i, C ii, N ii, Si ii, Ca ii, Si iii, Fe iii, C iv, Si iv, N v, and O vi. We measure a low H i column density of log N(H i) = 15.54 ± 0.05 in the HVC from multiple unsaturated H i Lyman series lines in the FUSE data. We determine a low dust depletion level in the HVC from the relative strength of silicon, iron, and calcium absorption relative to oxygen, with [Si/O] = −0.33 ± 0.14, [Fe/O] = −0.30 ± 0.20, and [Ca/O] = −0.56 ± 0.16. Analysis of the high-ion absorption using collisional ionization models indicates that the hot plasma is multiphase, with the C iv and Si iv tracing 104.9 K gas and N v and O vi tracing 105.4 K gas. The cloud’s metallicity, dust content, kinematics, and close proximity to the disk are all consistent with a Galactic wind origin. As the HD 156359 line of sight probes the inner Galaxy, the HVC appears to be a young cloud caught in the act of being entrained in a multiphase Galactic outflow and driven out into the halo.


INTRODUCTION
The supermassive black hole, Sagittarius A * , and surrounding regions of active star formation power an outflowing multiphase wind from the center of the Milky Way (MW).Outflows play a critical role in the baryon Corresponding author: Frances Cashman, Andrew Fox frcashman@stsci.edu,afox@stsci.educycle, the cycling of gas into and out of galaxies, which helps regulate the evolution of galaxies.The MW offers us a front-row seat view to study outflows over a range of wavelengths and phases and understand their impact on galaxy evolution.
Ultraviolet (UV) absorption-line studies have found gas outflowing at high-velocity in a number of sight lines passing through or near the Fermi Bubbles (see Keeney et al. 2006;Zech et al. 2008;Fox et al. 2015;Bordoloi et al. 2017;Savage et al. 2017;Karim et al. 2018;Ashley et al. 2020Ashley et al. , 2022)).These high-velocity clouds (HVCs) consist of neutral (e.g., C I, O I), singly ionized (S II, Si II, Fe II), and highly ionized (C IV, Si IV, O VI) gas.Cold gas has also been detected in the nuclear outflow via emission in neutral hydrogen and molecules.Several hundred H I 21 cm clouds have been detected at low latitudes in the Fermi Bubbles outflowing from the Galactic Center (GC) (McClure-Griffiths et al. 2013;Di Teodoro et al. 2018;Lockman et al. 2020).Carbon monoxide (CO) has been detected in submillimeter emission in two dense molecular clouds comoving within the H I 21 cm outflow (Di Teodoro et al. 2020).More recently, molecular hydrogen was discovered in the Galactic nuclear outflow ∼1 kpc below the GC (Cashman et al. 2021).Outflowing gas associated with the Fermi Bubbles has also been seen in optical emission (Krishnarao et al. 2020a).Together, these observations of HVCs near the Galactic Center provide observational constraints on the properties of the MW nuclear wind.
The outflowing wind interacts with gas in the disk in multiple ways.First, it can disrupt the position and velocity of the disk (Krishnarao et al. 2020b), leading to warps and perturbations.Second, it can accelerate and entrain cool gas into the halo.In some instances, this high-velocity gas may cool, lose buoyancy, and fall back onto the disk as a "Galactic fountain", supplying new fuel for further star formation (see Shapiro & Field 1976;Bregman 1980;Kahn 1981;de Avillez 1999).In other cases, the high-velocity gas may survive being accelerated into the halo, and even escape.Cloud survival is the subject of recent theoretical work focusing on how cool gas clouds develop and grow in the hot wind (see Gronke & Oh 2020;Sparre et al. 2020).
Metallicities of HVCs are important to determine because they provide direct evidence for the origin of the clouds.Fully UV-based metallicities (with the metal and H I column densities measured along the same UV sight lines) are ideal to ensure that the metal absorption and hydrogen absorption are tracing the same gas, but they are rare because of the difficulty in isolating unsaturated high-velocity H I components in UV absorption.These high-velocity H I components are often saturated or blended, even for high-order Lyman series lines, thus there is only a narrow range of N (H I) values for which precise measurements are possible (French et al. 2021).Therefore HVC metallicities are frequently determined from a combination of a UV measurement along an infinitesimal sight line with a H I measurement made using a much larger 21 cm beam.However, this combination introduces beam-smearing uncertainties on the metallicity.Currently, the only fully-UV-based metallicities are for two HVCs toward M5-ZNG1 (Zech et al. 2008), but this is a high-latitude direction (b ≈ 50 • ) far from the GC.
At lower latitudes, we can target UV-bright stars as background sources, but very few UV spectra of dis-tant GC stars (d 8 kpc) exist.Therefore, it is crucial to fully analyze the few spectra available, including LS 4825 (Savage et al. 2017;Cashman et al. 2021) and HD 156359 (this paper).It is precisely these low-latitude sight lines that are likely to harbor young clouds that have only recently been entrained in the Galactic outflow.Finding an outflowing cloud at low latitude means catching it close to its origin, thus offering a rare opportunity to study the cloud before it has undergone significant mixing as it begins its journey into the halo.Observing recently entrained clouds therefore gives a snapshot of outflowing gas when it first exits the disk.
In this paper we present a detailed spectroscopic analysis of an HVC detected toward an inner Galaxy sight line, HD 156359.By studying the properties of this cloud, we characterize the physical and chemical conditions of the gas in the inner Galaxy, in a region likely influenced by the nuclear wind., −14.52 • , 9 kpc) is one of the best-studied GC sight lines (Sembach et al. 1991(Sembach et al. , 1995)).The sight line lies in the inner Galaxy at the boundary of the southern Fermi Bubble and within the eROSITA X-ray bubble (Predehl et al. 2020), a region of enhanced X-ray emission (see Figure 1).This sight line also passes close to a complex of small HVCs dubbed "Complex WE" by Wakker & van Woerden (1991), less than 1 • from one of the densest cores.Sembach et al. (1991) classified HD 156359 as a O9.7 Ib-II star on the basis of stellar photospheric lines and wind profiles in its UV spectrum.The spectral type and apparent magnitude of the star yield a spectroscopic distance of 11.1±2.8kpc.However, recent Gaia EDR3 parallax measurements (Bailer-Jones et al. 2021) place HD 156359 at a distance of 9.0 +3.1 −2.1 kpc, which implies a z-distance below the plane of 2.3 kpc.We adopt the Gaia distance in our analysis.
The sight line toward HD 156359 intersects at least three spiral arms, the Sagittarius, Scutum, and Norma arms.Spiral-arm models from Vallée (2017) give the expected velocities of the Sagittarius, Scutum, and Norma arms at approximately −10, −55, and −103 km s −1 , respectively.Thus, absorption components observed at these velocities can be attributed to gas in (or associated with) the spiral arms.However, the spiral arms cannot explain the HVC observed at v LSR =125 km s −1 .
High-ion absorption toward HD 156359 was first observed with the International Ultraviolet Explorer (IUE; Sembach et al. 1991) and later by the Goddard High Resolution Spectrograph (GHRS; Sembach et al. 1995).
The high-ion information in the FUSE and HST STIS spectra of this sight line has not been previously published -the sight line is not included in the FUSE O VI survey of Galactic disk sight lines by Bowen et al. (2008), as that survey was restricted to |b| < 10 • .We also include analysis of a single archival optical spectrum of HD 156359 taken with FEROS at the European Southern Observatory (ESO) at La Silla.Details of all observations of spectra used in this paper are described below.

UV Observations
HD 156359 was observed by FUSE (Moos et al. 2000) under programs P101, S701, andU109 between 2000 April 12 and2006 April 25 (PIs Sembach, Andersson, andBlair, respectively).The raw spectra were downloaded from the MAST FUSE archive, and the CalFUSE pipeline (v3.2.2; Dixon & Kruk 2009) was used to reduce and extract the spectra.Data from the SiC, LiF1, and LiF2 channels were used for the analysis.A detailed explanation of the data reduction, as well as refinements to the CalFUSE data reduction procedures can be found in Wakker et al. (2003) and Wakker (2006).The spectra have a signal-to-noise ratio (S/N) ∼12-26 per resolution element and a velocity resolution of 20 km s −1 (FWHM).The data were binned by three pixels for the fitting analysis.The FUSE wavelength coverage is ∼912-1180 Å.
A single exposure of HD 156359 was obtained with HST/STIS on 2003 March 26 under program 9434 (PI Lauroesch) using the E140M grating.The data were downloaded from the MAST HST archive and reduced using calstis (v.3.4.2, Dressel et al. 2007).The echelle orders were combined to create a single continuous spectrum, and in regions of order-overlap spectral counts were combined to increase the S/N ratio.The data have S/N ∼10-38 per resolution element and a FWHM velocity resolution of 6.5 km s −1 , i.e. spectral resolution (λ/∆λ) ∼45,800.The STIS E140M wavelength coverage is ∼1160-1725 Å. Wavelengths and velocities for the absorption-line features are given in the local standard of rest (LSR) reference frame, where the correction factor v LSR − v Helio = −0.35km s −1 for the direction toward HD 156359.

Optical Observations
A spectrum of HD 156359 was captured using the Fibre-fed Extended Range Optical Spectrograph (FEROS) at the European Southern Observatory (ESO) La Silla 2.2 meter telescope on 2006 April 30 (PI: Bouret, PID: 077.D-0635).The retrieved archival product covers the wavelength range ∼3565-9214 Å with R = 48000 and FWHM = 6.25 km s −1 .The primary data product was reduced automatically using the FEROS Data Reduction Software (DRS) pipeline version fern/1.0.In the automated reduction process, the bias is subtracted using overscan regions and bad columns are replaced by the mean values of the neighboring columns.The orders are rectified and then extracted using the standard method.Next, the extracted spectra are flat-fielded and wavelength calibrated, then rebinned to a constant dispersion of 0.03 Å.Finally, the individual orders are combined into a single 1D spectrum.As the archival reduced data are not flux calibrated, the continuum of the stellar spectrum was normalized using the linetools software package.

MEASUREMENTS
In this section we describe our absorption-line measurement processes, including stellar continuum fitting and procedures for measuring lines of different ionization states.Our measurements are based on the VP-FIT software program (v12.2;Carswell & Webb 2014), which we used to conduct a set of Voigt-profile fits to the absorption-line profiles, with wavelengths and oscillator strengths taken from the compilations of Morton (2003) and Cashman et al. (2017).The measured lines are listed in Table 1.

Stellar continuum modeling
Given the spectral type of HD 156359 (O9.7 Ib-II; Sembach et al. 1991), the stellar continuum placement requires careful consideration.In addition to estimating the stellar continuum through comparison with FUSE spectra of stars of similar spectral type, we also constructed a TLUSTY model (Lanz & Hubeny 2003) of this type of star as a continuum placement guide.This was particularly useful in regions where the interstellar absorption was stronger.
The adopted TLUSTY model has T eff = 26000 K, log g = 3.0, v turb = 10.0 km s −1 and is rotationally broadened by 90 km s −1 in order to match the Si III 1300 Å triplets.We matched the observed UV wind signatures found in the observed FUSE and STIS spectra, where all observed spectra are adjusted for the stellar radial velocity v rad = −82 km s −1 (Gontcharov 2006).The model is reddened by an E(B − V ) = 0.09, using the mean Fitzpatrick & Massa (2007) extinction curve (extrapolated to 912 Å), and Lyman series absorption for a foreground H I column of N (H I) = 6.0 × 10 20 cm −2 is also included (Sembach et al. 1991).The model is then scaled by 1.25 × 10 −20 in order to match observations to ∼10%.Finally, all spectra were binned to 0.1 Å, ∼15 to 30 km s −1 , depending on wavelength, see Figure 2. When consulting the model as a guide for interstellar line continua, additional tweaks of about 5% were = 20.78(Sembach et al. 1991) is also applied.The flux was scaled by 1.25×10 −20 to provide an overall agreement with observations to ∼10%.
needed to match local continua.Our subsequent column density measurements account for the uncertainty inherent in the continuum-placement process.

H I absorption
The FUSE spectrum of HD 156359 covers almost the entire H I Lyman series, from Lyman-β at 1025 Å down to the Lyman limit at 912 Å.However, we limited our fitting to H I lines in regions where the stellar continuum was better-defined and showed the least contamination from stellar features and/or interstellar molecular lines, which are prolific in the FUSE bandpass.
We selected H I λλ923, 926, 930, and 937 (see Figure 3) to derive the H I column density.We fit a continuum to each of these lines locally.In order to account for sensitivity to continuum placement, we consider a high and low continuum placement for our column density measurement across the wavelength range λ917-944.H I λ923 lies in a noisy region of the upper Lyman lines frequently associated with the Inglis-Teller effect (Inglis & Teller 1939), in which overlapping stellar absorption lines have merged to produce the appearance of a depressed continuum.We perform a simultaneous Voigt profile fit to λλ926, 930, 937 with the positions of the low-velocity absorption initially based on weak ISM lines.We derive log N HI = 15.54 ± 0.02 ± 0.05 in the HVC at +125.6 km s −1 , where the first error is the statistical error due to photon noise and the second is the systematic continuum-placement error.The result of these fits is shown Table 2 and in Figure 3, where we show that the simultaneous fit is also consistent with the noisier H I λ923 profile.We include a detailed explanation of our continuum placement procedures and how we considered contamination from stellar features in Appendix A. An important step in measuring H I absorption in FUSE spectra is to decontaminate H 2 absorption.Fortunately, the FUSE LiF1 and LiF2 channels provide us with an opportunity to model isolated H 2 lines.This allows us to account for the Milky Way's molecular contribution to the interstellar absorption lines blended with H I in the SiC2A spectrum.We conducted an H 2 decontamination and describe and illustrate this process in detail in Appendix B.

O I absorption
After placing a smooth continuum through the O I line at 1302 Å in the STIS E140M spectrum, we noticed a small absorption feature near +125 km s −1 (see Figure 3).This feature spans five pixels and has a significance of 3.2σ (EW /σ EW =3.2).Since only one STIS exposure exists, we cannot confirm the O I detection in another dataset, and this feature is not seen in the much weaker O I line at 1039 Å in the FUSE spectrum.To explore whether this feature is real, we compared the absorption profile to another low-ion line.Figure 4 shows an apparent column density profile comparison of the O I 1302 Å line to the weak Si II 1526 Å line.We chose λ1526 because it is unsaturated, unblended, and seen in a region with high signal to noise.Although noisier, the profile of O I λ1302 is very similar to Si II λ1526 in the velocity region near +125 km s −1 , and the ratio of their apparent column densities is flat with velocity.To confirm this similarity, the inset plot in Figure 4 shows a linear fit to the ratio over five pixels (over 2 resolution elements) with a slope equal to zero within the margin of error, as expected for a genuine O I detection.Therefore, we proceed with treating the O I λ1302 feature as real on the basis that: 1) the line is detected at 3.2σ significance, 2) there is close kinematic consistency with Si II, and 3) the feature is centered at the same velocity as multiple other ions, e.g., H I, C II, N II, and Ca II.After applying a low and high continuum to this region, we measure log N OI = 12.43 ± 0.08 ± 0.07, which includes the statistical error and a continuum-placement (systematic) error.This measurement is the foundation of our metallicity measurement in the HVC, which we discuss in Section 4.

Low and intermediate ion absorption
We detected C II λ1334, Si II λλ1190, 1193, 1304, 1526, N II λ1083, Si III λ1206, and Fe III λ1122 in the HVC near +125 km s −1 in the FUSE and STIS spectra (see Figure 3).Our absorption-line measurements of these lines are given in Table 2. Achieving a simultaneous Voigt profile fit using all Si lines is hindered by difficulties in continuum placement for the stronger lines since large spans of the variable stellar continuum are absorbed.Instead, we adopt the Voigt profile measurement log N SiII = 12.94±0.07for the weakest unblended line available at 1526 Å.We show that it is a good fit for λλ1190, 1193, 1304 in Figure 3.Although detected, N II 1083 lies in a portion of the FUSE spectrum where the stellar continuum is highly variable over a short range in wavelength.Comparing the neighboring low-velocity H 2 J3 1084 Å line to other J3 lines in regions with smooth continua reveals that the continuum must drop significantly in this region, making it difficult to estimate log N NII .For this reason, our measurement for N II is an upper limit.Fe II 1144 has a low absorption profile in the vicinity of the HVC.We applied a high, middle, and low continuum across the absorbing region of λ1144 and calculate a significance of 2.2σ after including a continuum placement error.We find a 3σ limiting column density of log N FeII ≤ 12.99, however, given the Fe III detection in the HVC, we use the Fe III column density in metallicity calculations going forward.In addition to the high-velocity absorption we observe for Fe III λ1122, we detect an intermediate-velocity cloud (IVC) centered near +86 km s −1 with log N = 12.98±0.13.
The archival ESO FEROS optical spectrum covers the Ca II K and H and Na I D lines, see Figure 5. Telluric H 2 O absorption lines are present in the velocity range of the HVC, however, there is a non-detection of highvelocity Na I and we find a 3σ limiting column density limit of log N NaI < 9.45.We see absorption in Ca II across 6 pixels at 123.7 km s −1 and measure an AOD column density of log N CaII = 10.68±0.01 for Ca II 3934.We also performed a simultaneous Voigt profile fit to Ca II λ3934, 3969 and find log N = 10.63±0.10.The resulting b-value for this small component has a high error.However, since the log N value from Voigt profile fitting agrees with the AOD measurement within the margin of error, we adopt its value.

High ion absorption
We detected C IV λλ1548, 1550, Si IV λλ1393, 1402, and N V λλ1238, 1242 absorption near +130 km s −1 in our STIS E140M spectrum.The C IV and Si IV profiles are complex, showing multiple distinct low and intermediate-velocity components, in addition to the HVC.The high-ion absorption features lie on welldefined continua due to their corresponding broad and smooth stellar P Cygni wind profiles, which gradually elevates their stellar continua in the region from −100 to +200 km s −1 for this star, as shown in Figure 6.The high-ion continuum fitting was also guided by fits to the GHRS data previously published by Sembach et al. (1991Sembach et al. ( , 1995)).We performed Voigt-profile fitting to all high ions with the initial positions of the components and b-values determined from the C IV 1548 Å line.The position and b-values were allowed to vary freely and the results for the fits are given in Table 2.We note the detection of an IVC centered near +78 km s −1 in both C IV and Si IV with log N = 13.57± 0.08 and 13.09 ± 0.10, respectively.
We observe high-velocity absorption in O VI 1031 and 1037 in the LiF1 channel of the FUSE spectrum at v LSR = 141 km s −1 .However, we only include λ1031 in our analysis, because λ1037 lies on the steep blueward side of the O VI P Cygni profile and suffers from significant blending with C II * 1037 and H 2 J1 1038 Å.
The wavelength region around O VI 1031 Å contains several absorption lines which can serve as contaminants, most notably HD 6-0 R(0) λ1031.915,Cl I λ1031.507,and several lines of molecular hydrogen including H 2 6-0 P (3) λ1031.192and H 2 6-0 R(4) λ1032.351.To gauge the effect of contamination by HD 6-0 R(0) λ1031.915,we examined other HD lines of similar oscillator strength that are isolated from interstel- lar absorption, such as 5-0 R(0) λ1042.850,7-0 R(0) λ1021.460,and 8-0 R(0) λ1011.461.We detect no HD molecular absorption in these lines and conclude that no subtraction of the HD 6-0 R(0) line at λ1031.915 from the O VI profile is necessary.For Cl I, a small amount of absorption in Cl I λ1347 is present near 0 km s −1 in the STIS spectrum and we simultaneously fit Cl I λ1031 to derive its contribution to the O VI profile.For molecular hydrogen, H 2 6-0 R(4) λ1032.351 is the most relevant potential contaminant as it overlaps with the high-velocity component in O VI λ 1031.We modeled other isolated H 2 J4 lines in the FUSE spectrum, e.g.5-0 R(4) λ1044.543and 4-0 R(4) λ1057.381(see description in Appendix B).Through simultaneous fitting of those lines with H 2 6-0 R(4) λ1032 , we account for the small amount of H 2 present in the O VI high-velocity absorption.A similar procedure was followed to account for contamination from H 2 6-0 P (3) λ1031.192near v LSR = −200 km s −1 using the isolated H 2 J3 lines 6-0 R(3) λ1028.986and 5-0 P (3) λ1043.503.The resulting Voigt profile fit and O VI column density are shown in Figure 6 and Table 2.

RESULTS: THE COOL LOW-ION GAS
The neutral and low ionization species including H I, O I, C II, N II, Si II, Ca II, Fe II, Si III, and Fe III trace the cool photoionized phase of the HVC detected at v LSR = +125 km s −1 toward HD 156359.In this section we use the ratios of metal column densities to H I column densities to derive the ion abundances for each observed species in the HVC, as shown in Table 3.The elemental abundances are then derived from the ion abundances using ionization corrections derived from custom photoionization modeling.Finally a comparison of the relative abundances of different elements is used to derive the dust depletion pattern in the HVC.

Photoionization Modeling
To model the ionization breakdown in the HVC and characterize its physical conditions, we ran a multi-dimensional grid of Cloudy (v.17.02, Ferland et al. 2017) photoionization models for the low-and intermediate-ions, assuming they arise in the same gas phase as the H I and are photoionized by the same incident radiation field.Our model assumes a plane-parallel slab of uniform density exposed to the magnitude of the escaping UV ionizing flux of the Milky Way (Fox et al. 2005;Barger et al. 2013;Fox et al. 2014) at the location of HD 156359 and includes the extragalactic UV background radiation field from Khaire & Srianand (2019).We ran our models for a grid of metallicity values with [Z/H] varying from −0.5 to +1.5 in steps of 0.1 dex, each over a range of hydrogen number densities log (n H /cm −3 ) from −3 to 0, in order to explore possible metallicities and densities.We determined the best-fit log n H for each metallicity model by matching the observed column density ratio of Si III/Si II to the model value (see the top panel of Figure 7).The Si III/Si II ratio was chosen because both ions have unsaturated detections in the HVC, and using a ratio of adjacent ions from the same element (Si) minimizes metallicity or depletion effects from different metal ions.Using the pairs of metallicity and log n H determined from the Si III/Si II  4. The error listed for log N (O I) is the combination of the statistical error and a continuum placement error in quadrature, and is ± 0.08 ± 0.07 separately.b This measurement is a non-detection and the 3σ limiting column density is given.ratio, we narrow the range of metallicities by identifying which models are consistent with the observed log N OI = 12.43 ± 0.11, as illustrated in the bottom left panel of Figure 7.We then run an even finer grid of metallicity values in increments of 0.02 dex over the narrowed metallicity and number density region to determine the range of metallicities and densities allowed by the data.
We find that metallicities in the range 0.18 ≤ [O/H] ≤ 0.51 and densities from −1.69 ≤ log (n H /cm −3 ) ≤ −1.37 are consistent with the data, giving a model best-fit of [O/H] = 0.36 +0.15  −0.18 at log n H = −1.53±-0.16.We repeated this procedure for C II (see bottom-right panel of Figure 7), which is expected to be weakly depleted (Jenkins 2009), and find that the range 0.21 ≤ [C/H] ≤ 0.42 overlaps with the observed data, giving a model best-fit of [C/H] = 0.30 +0.12 −0.09 at log n H = −1.49+0.09 −0.10 .The agreement between the oxygen-based metallicity and the carbon-based metallicity lends support to our methodology and to the robustness of the super-solar metallicity we infer.
The metallicity for this cloud ranks among the highest UV-based metallicities of any HVC observed thus far, along with the HVC at −125 km s −1 toward M5-ZNG (l = 3. • 9, b = +47.• 7 at z = +5.3kpc) with [O/H] = +0.22 ± 0.10 (Zech et al. 2008).M5-ZNG was also observed with FUSE and STIS E140M spectra.Their observed metallicity is not corrected for ionization effects, but is likely higher than reported, given the measured low log N HI = 16.50 ± 0.06 in the cloud and that oxygen shows a positive ionization correction when log N (H I) < 18.5 (Bordoloi et al. 2017).The metallicity of the HD 156359 HVC is also on the high end of the range of <20% to 3 times solar reported by Ashley et al. (2022) in their survey of metallicities of Fermi Bubble HVCs.We note that a super-solar abundance in the inner Galaxy may not be unexpected, given the oxygen abundance gradients reported from emission line measurements in H II regions (see Wenger et al. 2019;Arellano-Córdova et al. 2021).
derived from the difference between the model elemental abundance and the observed ion abundance.c Ionization-corrected elemental abundance [X/H] = [X i /H I] + IC(X i ).Also referred to as gas-phase abundance.
is the depletion of element X relative to oxygen.

Ionization corrections
In   ionization state of hydrogen and oxygen together Field & Steigman (1971), and (2) oxygen is only lightly depleted onto dust grains (Jenkins 2009).However, the assumption that [O I/H I]=[O/H] breaks down when N (H I) is so low that the gas is optically thin or when the ionizing photon flux is extremely high (Viegas 1995).In 3.0 2.5 2.0 1.5 1.0 0.5 0.0 log (n H /cm 3 ) 67 the HVC at +125 km s −1 toward HD 156359, log N (H I) is only 15.54±0.05,so an ionization correction must be made to account for the column densities of all observed neutral and ion stages.
Our Cloudy models directly provide the ionization corrections for all our observed ion stages.We define the ionization correction as the difference between the model (true) elemental abundance and the observed ion abundance, i.e.,

Dust depletion effects
Following convention, we define the depletion δ O (X) of each refractory element X as the ionization-corrected abundances of that element compared to the ionizationcorrected oxygen abundance, i.e., where oxygen represents an undepleted (or lightly depleted) volatile element.In this formalism, a negative value of δ O (X) means that element X is depleted relative to oxygen.This method assumes that the total (gas+dust) abundances are solar, which is often assumed to apply to the Galactic ISM (Savage & Sembach 1996), though local ISM abundance variations cannot be ruled out in some sight lines (De Cia et al. 2021).We show δ O (X) for the low ions in the bottom panel of Figure 8 compared to F * , the line-of-sight depletion strength factor, for [X/O] determined for low-depletion (F * =0) and high-depletion (F * =1) gas from the comprehensive ISM gas-phase element depletions study of Jenkins (2009).
For the HVC toward HD 156359, we find a low value of δ O (C)=−0.09±0.17,i.e. carbon shows no significant depletion.This is consistent with the low values of [C/O] measured in Galactic ISM gas (Jenkins 2009).
For the other low ions, we find a 3σ upper limit for the nitrogen depletion δ O (N)≤ +0.57and low depletions for silicon, iron, and calcium of δ O (Si)=−0.33±0.14, δ O (Fe)=−0.30±0.20,and δ O (Ca)=−0.56±0.16,respectively.To summarize the final results of the CLOUDY modeling after the dust depletion levels have been derived, Figure 9 shows the model curves of the detected ions shifted by their respective depletions at [X/H] = +0.36 and log n H = −1.53.The ability of this model to match the observed column densities shows that all low-ion measurements can be explained by photoionization once dust depletion effects are accounted for.
A low level of dust depletion for the HVC is consistent with the well-known Routly-Spitzer effect (RS effect), in which the amount of dust observed in high-velocity clouds decreases significantly at higher LSR velocity (see Routly & Spitzer 1952;Siluk & Silk 1974;Vallerga et al. 1993;Smoker et al. 2011;Ben Bekhti et al. 2012;Murga et al. 2015).The RS effect is typically traced by the N (Na I)/N (Ca II) ratio, which has been found to significantly decrease with increasing LSR velocity.Although the RS effect was historically interpreted as a dust depletion effect, it may also be related to ionization effects, and likely differs depending on the environment of the cloud.We measure an 3σ upper limit of log N (Na I)/N (Ca II) ≤ −1.2 in the +124 km s −1 HVC toward HD 156359 from the optical FEROS data, which is on the low end of the observed ratios for HVCs (Smoker et al. 2011;Ben Bekhti et al. 2012;Murga et al. 2015), and is also lower than is typically observed at low velocities.In summary, the dust depletion pattern we derive in the HVC is consistent with other HVCs and with the Routly-Spitzer effect, even though the cloud has high metallicity.

RESULTS: THE HIGHLY IONIZED GAS
The HVC toward HD 156359 shows high-ion absorption in the resonance doublets of Si IV, C IV, N V, and O VI at a velocity range of 130-140 km s −1 (Figure 6; Table 2), slightly higher than the velocity range of the low ions, which are centered near 125 km s −1 .To compare the high-ion absorption between different ions, we show in Figure 10  cal depth profiles of each high ion.These plots provide a way to inter-compare the profiles of different ions in a linear manner (e.g., Fox et al. 2003).The HVC line profiles of C IV and Si IV are well aligned in velocity, but are offset from the centroid of O VI absorption by ≈8 km s −1 .This suggests that the O VI exists in a separate (likely hotter) gas phase.The STIS profile of N V is too noisy to draw any conclusions about line centroid.However, the GHRS N V profiles (Sembach et al. 1995) are less noisy and have a similar profile shape to the O VI seen in the FUSE data, supporting the placement of O VI and N V in the same phase.

Collisional Ionization Modeling
The observed high ions have column densities that are too high to be explained by the Cloudy photoionization models described in Section 4.1.Where the Si IV and C IV column densities differ by ∼0.7 and 1.0 dex respectively, the N V and O VI differ by orders of mag-  Gnat & Sternberg (2007).The observed ratios in the HD 156359 HVC are shown as dashed orange horizontal lines.The top and bottom panels show the C IV/Si IV and O VI/N V ratios, respectively.In both panels, the solid black curve is the collisional ionization equilibrium model ion fraction.The solid green and violet curves are the timedependent isochoric and isobaric collisional ionization models for a solar metallicity, respectively.Their dashed counterparts are the same but for a metallicity twice the solar value.The vertical dashed blue line is the temperature of the gas determined by the models.The vertical blue band in the top panel represents a second time-dependent solution for a range of temperatures from log T = 4.2-4.4K.
nitudes.These discrepancies imply that a separate ionization mechanism is required for the high ions.We do not account for energetic photons from the radiation field intrinsic to the Fermi Bubbles, however, we explore collisional ionization as a separate mechanism for gas with temperatures of T 10 5 K.We consult the collisional ionization models from Gnat & Sternberg (2007) to determine the exact range of temperatures allowed by the high-ion column densities and their ratios.While we considered the more recent models from Gnat 2017, which include photoionization effects from the extragalactic UV background, we do not adopt them because HD 156359 lies in close proximity to the radiation field of the Galactic plane which is much stronger.We considered both the collisional ionization equilibrium and non-equilibrium regimes.
A single-temperature solution explaining all four high ions (Si IV, C IV, N V, O VI) in the HVC is ruled out, as no such solution exists to the observed high-ion column densities.Instead, we find that a two-phase solution is needed, with one temperature explaining the N CIV /N SiIV ratio and another explaining N OVI /N NV .This two-phase model is consistent with the kinematic information in the UV spectra, where Si IV and C IV show very similar line profiles but O VI is offset in velocity centroid.The two-phase high-ion model is illustrated in Figure 11, which shows the observed high-ion ratios of N CIV /N SiIV and N OVI /N NV compared to model predictions from Gnat & Sternberg (2007) for solar ([Z/H]=0) and super-solar ([Z/H]=+0.3)metallicities.
For the C IV/Si IV phase, we find two possible solutions for the temperature.First, the non-equilibrium isochoric (constant volume) and isobaric (constant pressure) models give a low-temperature solution at T = 10 4.2−4.4K, where the lower end corresponds to the solar-metallicity isochoric model and the higher end with the super-solar isobaric model.Second, the collisional ionization equilibrium (CIE) model returns a higher temperature T =10 4.9 K.We are inclined to adopt the non-equilibrium (lower-temperature) solution, as plasma near T =10 5 K is at the peak of the cooling curve, where oxygen dominates the radiative cooling and the gas can cool faster than it recombines, reaching a nonequilibrium state.For the O VI/N V phase, a single temperature solution for the observed log N OVI /N NV =1.01 is found at T =10 5.4 K for all models (see bottom panel of Figure 11).
In conclusion, we can successfully model the high-ion plasma in the HVC toward HD 156359 as containing collisionally-ionized gas at two temperatures: a cooler phase seen in C IV and Si IV at T = 10 4.2−4.4K, and a hotter phase seen in N V and O VI at T =10 5.4 K.

SUMMARY
Using archival FUSE, HST STIS, and ESO FEROS spectra, we have analyzed the chemical composition of the HVC near +125 km s −1 toward HD 156359, a massive star lying 9 kpc away toward the Galactic Center.The sight line passes less than 1 • from one of the densest cores of a complex of small HVCs dubbed "Complex WE" by Wakker & van Woerden (1991), as shown in Figure 1.Furthermore, the sight line passes through a region of enhanced X-ray emission (the southern eROSITA Bubble; Predehl et al. 2020); this region indicates energetic feedback from the Galactic Center.Our main results are as follows.
1. We determined an H I column density measurement of log N (H I) = 15.54 ± 0.05 in the HVC using unsaturated Lyman series absorption lines.
2. 5. We detect high ion species C IV, Si IV, N V, and O VI near +130 km s −1 , at slightly higher velocities than the velocity range of the lower ions.We find that an ionization mechanism separate from photoionization, such as collisional ionization, is required to explain the column densities of the high ions.We determine that a two-phase temperature solution best explains the observed N CIV /N SiIV and N OVI /N NV ratios, with a cooler phase seen in C IV and Si IV at T = 10 4.2−4.4K, and a hotter phase seen in N V and O VI at T =10 5.4 K.
The high metallicity, low depletion, complex high-ion absorption, and high positive velocity of the HVC toward HD 156359 are all consistent with a wind origin, in which a swept-up clump of material is being carried out from the Galactic disk into the halo.As such, this HVC may represent a freshly entrained cool clump of gas caught in the act of being accelerated into the nuclear wind.While we cannot rule out a foreground origin for the HVC, in which the cloud exists at an anomalous velocity somewhere between the Sun and the Galactic Center, we can exclude a spiral-arm explanation for the HVC on kinematic grounds, because the cloud's +125 km s −1 velocity lies over 100 km s −1 away from the nearest spiral arm, the Sagittarius Arm at ∼ −10 km s −1 .Under a biconical outflow model (Fox et al. 2015;Bordoloi et al. 2017;Di Teodoro et al. 2018), the cloud's radial velocity and its location only 2.3 kpc below the disk imply a short timescale of ∼5 Myr for the HVC to have reached its current position, which is much shorter than the timescales on which chemical mixing is expected to be significant (10s to 100s of Myr; Gritton et al. 2014;Heitsch et al. 2022).Our observations therefore provide a snapshot into the chemical and physical conditions prevailing in this early stage of a nuclear outflow, before chemical mixing has diluted or enriched the gas from its initial state.
We gratefully acknowledge the invaluable contributions to early versions of this manuscript from the late Blair Savage, who passed away during the manuscript's preparation.This paper would not have been possible without Blair's foundational work on the HD 156359 sight line, the chemical abundances in the ISM, apparent optical depth analysis, and the inner Galaxy.We gratefully acknowledge support from the NASA Astrophysics Data Analysis Program (ADAP) under grant 80NSSC20K0435, 3D Structure of the ISM toward the Galactic Center.The FUSE data were obtained under program P101.FUSE was operated for NASA by the Department of Physics and Astronomy at the Johns Hopkins University.D.K. is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-2102490.We thank Isabel Rebollido for valuable conversations on the FEROS spectrograph.The FUSE and HST STIS data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST) at the Space Telescope Science Institute.The ESO FEROS spectrum was obtained from the ESO Archive Science Portal.
The FUSE and HST STIS data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST) at the Space Telescope Science Institute.The specific observations analyzed can be accessed via 10.17909/rrnk-3e58.Stellar absorption lines present a continuum-fitting challenge, particularly in the complex far-UV FUSE spectra used to derive the H I column density in the HVC under study.Our approach followed in the analysis is to fit the continuum locally around each H I Lyman series line of interest, since this allows us to account for stellar absorption lines, which are present in the continuum when the radiation field encounters the HVC.However, for completeness here we consider a global continuum fitting process, which fits the H I continuum over a larger range (916-944 Å) in the SiC2 channel.This approach neglects stellar absorption features but ensures the continua are continuous between adjacent lines in the Lyman series.Wavelength regions free from stellar emission and absorption lines defined the flux of the global continuum.We show a global fit to the stellar flux in the vicinity of the higher order H I Lyman series lines in Figure A1, in which areas of emission and absorption due to both stellar and interstellar features can be seen.We performed a simultaneous Voigt profile fit on the same H I lines modeled with the global continua, using the mid-, high-, and low-continuum fits shown in Figure A1.A separate local continuum was applied to the depressed region resulting from stellar line blanketing and which contains H I λ923.Using the global continuum fit (which has a slightly higher continuum placement than the local fits) results in log N HI = 15.69±0.03±0.03 for the HVC, where the first error is the statistical error due to photon noise and the second error is the systematic continuum-placement error.This corresponds to a moderate difference of ∆log N HI = 0.15 dex between the H I column densities derived by the global and local continuum methods.

B. DECONTAMINATION OF MOLECULAR HYDROGEN
The FUSE LiF1 and LiF2 channels provide us with an opportunity to model isolated H 2 lines.This effort is important because it enables us to account for the Milky Way's molecular contribution to the interstellar absorption Left panel: The location of HD 156359 with reference to the Fermi and eROSITA bubbles.The composite Fermi-eROSITA image from Predehl et al. (2020), where the softer X-ray emission (0.6-1 keV, in cyan) envelopes the harder component of the extended GeV emission of the Fermi bubbles (in red; adapted from Selig et al. 2015).Right panel: H I column density map from the 21 cm HI4PI survey showing the distribution of H I in the region from 100-150 km s −1 HI4PI Collaboration et al. ( The HD 156359 Sight line HD 156359(at l, b, d = 328.68

Figure 3 .
Figure3.Velocity profiles of the neutral, low-, and intermediate-ion absorption lines detected toward HD 156359.The normalized flux is shown in black, the continuum level is in red, and the 1σ error in the normalized flux is in blue.The vertical line at 0 km s −1 marks the region associated with the Milky Way.Each of these profiles shows a Voigt fit to the data except for O I for which we provide an AOD measurement, and for Fe II, which is a non-detection.The solid green curve is the overall Voigt profile fit to the absorption.The magenta curve is the Voigt profile fit to the HVC absorption feature near +125 km s −1 and the vertical dashed magenta line marks the velocity centroid of the fitted component.

Figure 4 .
Figure 4.A comparison of the normalized apparent column density plots between O I 1302 in violet and Si II 1526 in green.The profiles have similar shapes although the O I 1302 profile shows a lower signal to noise.The inset panel shows the apparent column density ratio of Si II 1526 to O I 1302.A linear fit to 5 pixels in the region from +110 v 125 km s −1 has a slope near zero within the margin of error, supporting the reality of the detection in both ions.

Figure 5 .
Figure5.Velocity profiles of the Ca II λλ3934, 3969 and Na I λλ5891, 5897 absorption lines detected toward HD 156359 in the archival FEROS spectrum.The normalized flux is shown in black, the continuum level is in red.The gray vertical line at 0 km s −1 marks the absorption region associated with the Milky Way.The green curves in the top two panels are the Voigt profile fits to the data, where the green vertical line at 123.7 km s −1 marks the location of the HVC.There is a non-detection of high-velocity Na I, although telluric H2O absorption lines are present in this region.The lower end of the y-axis begins at +0.2 in all panels to improve the visibility of the absorption features.

Figure 6 .
Figure 6.Velocity profiles of the high-ion absorption lines detected toward HD 156359.The flux is shown in black and the 1σ error in the flux is in blue.The dashed black vertical line marks 0 km s −1 .The solid green curve is the overall Voigt profile fit to the absorption.The solid magenta curve is the Voigt profile fit to the HVC absorption near +125 km s −1 and the vertical dashed magenta line marks the velocity centroid of the fitted component.The pink and yellow curves are fits to components at intermediate velocities.

Figure 7 .
Figure 7. Cloudy photoionization models of the HVC toward HD 156359 that explore a parameter space in metallicity and number density, −0.50 ≤ [Z/H] ≤ +1.50 and −3 ≤ log (nH/cm −3 ) ≤ 0, respectively.Top panel: Model results for log NSiIII/NSiII for the set metallicities are shown as solid colored curves.The dashed horizontal turquoise line is the observed ratio of −0.62±0.17.The best-fit log nH for each metallicity model is determined by the intersection of the model and observed ratio and is marked by vertical colored lines and circle markers.Bottom left panel: The model curves for log NOI for each metallicity are plotted against log nH, where the values derived from log NSiIII/NSiII are marked with vertical lines and markers.The horizontal turquoise line and bar show the observed log NOI = 12.43 ± 0.11.The model data corresponding to the metallicity region from +0.18 ≤[Z/H] ≤ +0.51 and −1.69 ≤ log nH ≤ −1.37 overlap with the observed data.The model best-fit values are [Z/H] = +0.36+0.15 −0.18 and log nH = −1.53± 0.16.Bottom right panel: Same as for the bottom left panel, but for C II.The horizontal turquoise line and bar show observed log NCII = 13.83 ± 0.10.The model data corresponding to the metallicity region from +0.21 ≤ [Z/H] ≤ +0.42 and −1.59 ≤ log nH ≤ −1.40 overlap with the observed data.The model best fit values are [Z/H]= +0.30 +0.12 −0.09 and log nH = −1.49+0.09 −0.10 .

Figure 8 .
Figure 8. Top panel: ionization corrections for the low ions O I, C II, N II, Si II, Fe III, and Ca II (colored curves) plotted against log nH from our Cloudy model to the HD 156359 HVC.The model uses our best-fit [Z/H]=+0.36.The vertical line marks the best-fit log (nH/cm −3 )=−1.53,determined from the Cloudy models for O I (see Figure 7).The ionization correction calculated at −1.53 is added to the observed ion abundance to determine the elemental abundance.Middle panel: comparison of the ionization-corrected abundances of the low ions detected in the HVC, where the solar abundance is plotted as a horizontal line.Bottom panel: Comparison of the dust depletion levels δ(X) = [X/O] in the HVC with the depletion pattern for [X/O] measured for sight lines with the lowest and highest collective depletions (F * =0 and F * =1, respectively) fromJenkins (2009).A slight offset is applied in the x-direction of each element for distinction.

Figure 9 .
Figure 9.Our final Cloudy photoionization model of the low-and intermediate ions detected in the HVC toward HD 156359 near +125 km s −1 .The model uses the best-fit [Z/H] = +0.36determined from our O I analysis.These model predictions for each ion (colored dotted curves) have been scaled by the dust depletions required to match the observed values, where the observed column densities are indicated by circle markers at the best-fit log (nH/cm −3 )=−1.53 (black vertical line).

Figure 10 .
Figure 10.Optical depth profiles of the high ions, where each profile is arbitrarily normalized to its peak value for the velocity region shown in order to facilitate comparison.Top panel: C IV λ1548 and Si IV λ1393 display similar line shapes in the velocity region of the HVC from 112 v 140 km s −1 , indicated with cyan shading.Middle panel: C IV λ1548 and O VI λ1031 are compared, where the center of the HVC absorption for O VI is shifted to a higher velocity at v ≈ 141 km s −1 .Bottom panel: N V λ1238 and O VI λ1031 are compared, where the range of velocities attribute to absorption by O VI are indicated in light pink.

Figure 11 .
Figure11.Comparison of observed high-ion columndensity ratios with predictions from the collisional ionization models ofGnat & Sternberg (2007).The observed ratios in the HD 156359 HVC are shown as dashed orange horizontal lines.The top and bottom panels show the C IV/Si IV and O VI/N V ratios, respectively.In both panels, the solid black curve is the collisional ionization equilibrium model ion fraction.The solid green and violet curves are the timedependent isochoric and isobaric collisional ionization models for a solar metallicity, respectively.Their dashed counterparts are the same but for a metallicity twice the solar value.The vertical dashed blue line is the temperature of the gas determined by the models.The vertical blue band in the top panel represents a second time-dependent solution for a range of temperatures from log T = 4.2-4.4K.

Facilities:
FUSE, HST(STIS), ESO(FEROS) Software: astropy (Astropy Collaboration et al. 2018), Cloudy (Ferland et al. 2017), linetools (Prochaska et al. 2017), VPFIT (Carswell & Webb 2014), CalFUSE (Dixon & Kruk 2009), calSTIS (Dressel et al. 2007), FEROS-DRS (Kaufer et al. 1999) APPENDIX A. GLOBAL CONTINUUM FITTING Figure A1.A portion of the FUSE SiC2 spectrum from 916-944 Å in which several higher order H I Lyman series lines are identified.The flux and 1σ error in the flux are shown in black and blue, respectively.A global fit to the stellar continuum is traced by a solid red curve, with a high and low stellar continuum placement marked by dashed red lines.The orange box shows a local fit to a depressed portion of the stellar continuum created by stellar line blanketing, from ∼920-925 Å.

Table 3 .
HVC elemental abundances and depletions a Ion abundance [X i /H I] = log (X i /H I)HVC − log (X/H) where X i is the observed ion of element X.These are not corrected for ionization effects.The error includes both systematic and continuum placement uncertainties.b Ionization correction IC(