# Caught in the Act: A Metal-Rich High-Velocity Cloud in the Inner Galaxy

Frances H. Cashman Space Telescope Science Institute, 3700 San Martin Drive,
Baltimore, MD 21218, USA  Andrew J. Fox AURA for ESA, Space Telescope Science
Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  Bart P. Wakker
Department of Astronomy, University of Wisconsin-Madison, 475 North Charter
Street, Madison, WI 53706, USA  Trisha Ashley Space Telescope Science
Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  Derck Massa Space
Science Institute, 4750 Walnut Street Suite 205 Boulder, Colorado 80301, USA
Edward B. Jenkins Department of Astrophysical Sciences, Princeton University,
Princeton, NJ 08544-1001, USA  Dhanesh Krishnarao NSF Astronomy & Astrophysics
Postdoctoral Fellow, Johns Hopkins University, 3400 N. Charles Street,
Baltimore, MD 21218, USA  Department of Physics, Colorado College, 14 East
Cache La Poudre Street, Colorado Springs, CO 80903, USA  Space Telescope
Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  Felix J.
Lockman Green Bank Observatory, P.O. Box 2, Rt. 28/92, Green Bank, WV 24944,
USA  Robert A. Benjamin Department of Physics, University of Wisconsin-
Whitewater, 800 West Main Street, Whitewater, WI 53190, USA  Rongmon Bordoloi
Department of Physics, North Carolina State University, 421 Riddick Hall,
Raleigh, NC 27695-8202, USA  Tae-Sun Kim  Department of Astronomy, University
of Wisconsin-Madison, 475 North Charter Street, Madison, WI 53706, USA

(Received September 28, 2022)

###### Abstract

We characterize the chemical and physical conditions in an outflowing high-
velocity cloud in the inner Galaxy. We report a super-solar metallicity of
[O/H] = $+0.36\pm 0.12$ for the high-velocity cloud at $v_{\mathrm{LSR}}$ =
125.6 km s-1 toward the star HD 156359 (catalog ) ($l$ = 328.$\degree$7, $b$ =
$-$14.$\degree$5, $d$ = 9 kpc, $z$ = $-$2.3 kpc). Using archival observations
from FUSE, HST STIS, and ESO FEROS 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) = $15.54\pm 0.05$ 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]=$-0.33\pm 0.14$, [Fe/O]=$-0.30\pm 0.20$, and [Ca/O] =$-0.56\pm 0.16$.
Analysis of the high-ion absorption using collisional ionization models
indicates that the hot plasma is multi-phase, 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 (catalog ) line of sight probes
the inner Galaxy, the HVC appears to be a young cloud caught in the act of
being entrained in a multi-phase Galactic outflow and driven out into the
halo.

Galactic Center (565) — Ultraviolet astronomy (1736) — High-velocity clouds
(735) – Chemical abundances (224) — Galactic winds (572)

††journal: ApJ††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)

##  1 Introduction

Figure 1: Left panel: The location of HD 156359 (catalog ) 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. (2016) in a section of “Complex WE”
by Wakker & van Woerden (1991). The location of HD 156359 (catalog ) is marked
with a white cross. The contours correspond to log $N_{\mathrm{HI}}$ = 2.08,
6.23, 10.4, 14.5, 18.7 $\times$ 1018 cm-2.

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 cycle, 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.

Presently, the farthest-reaching known outflows in the MW are the Fermi (Su et
al., 2010; Ackermann et al., 2014) and eROSITA (Bland-Hawthorn & Cohen, 2003;
Predehl et al., 2020) Bubbles, which are giant gamma- and X-ray emission
extending $\sim$10 kpc and 14 kpc above and below the Galactic disk,
respectively (see Figure 1). The origin of the Fermi Bubbles is believed to be
an energetic GC Seyfert flare event $\sim$3.5 Myr ago (Bland-Hawthorn et al.,
2019).

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. 2020, 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 $\sim$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 ($b\approx 50\degree$) far from the GC.

At lower latitudes, we can target UV-bright stars as background sources, but
very few UV spectra of distant GC stars ($d\gtrsim 8$ 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 (catalog ). 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.

##  2 Data

###  2.1 The HD 156359 Sight line

HD 156359 (at $l,b,d$ = 328.68°, $-$14.52°, 9 kpc) is one of the best-studied
GC sight lines (Sembach et al., 1991, 1995). The sight line lies in the inner
Galaxy at the boundary of the southern Fermi Bubble and within the eROSITA
X-ray bubble (Predehl et al., 2020), a region of enhanced X-ray emission (see
Figure 1). This sight line also passes close to a complex of small HVCs dubbed
“Complex WE” by Wakker & van Woerden (1991), less than 1$\degree$ from one of
the densest cores. Sembach et al. (1991) classified HD 156359 (catalog ) 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$\pm$2.8 kpc. However, recent Gaia EDR3
parallax measurements (Bailer-Jones et al., 2021) place HD 156359 (catalog )
at a distance of 9.0${}^{+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 (catalog ) 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 $v_{\rm LSR}$=125 km s-1.

High-ion absorption toward HD 156359 (catalog ) 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 $|b|<10\degree$. We also include analysis of a single
archival optical spectrum of HD 156359 (catalog ) 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.

###  2.2 UV Observations

HD 156359 (catalog ) was observed by FUSE (Moos et al., 2000) under programs
P101, S701, and U109 between 2000 April 12 and 2006 April 25 (PIs Sembach,
Andersson, and Blair, 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) $\sim$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 $\sim$912–1180Å.

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 $\sim$10–38 per
resolution element and a FWHM velocity resolution of 6.5 km s-1, i.e. spectral
resolution ($\lambda$/$\Delta\lambda$) $\sim$45,800. The STIS E140M wavelength
coverage is $\sim$1160–1725Å. Wavelengths and velocities for the absorption-
line features are given in the local standard of rest (LSR) reference frame,
where the correction factor $v_{\mathrm{LSR}}$ $-$ $v_{\mathrm{Helio}}$ =
$-0.35$ km s-1 for the direction toward HD 156359 (catalog ).

###  2.3 Optical Observations

A spectrum of HD 156359 (catalog ) 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 $\sim$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.

##  3 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 VPFIT 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.

###  3.1 Stellar continuum modeling

Given the spectral type of HD 156359 (catalog ) (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 $T_{\mathrm{eff}}$ = 26000 K, log $g$ = 3.0,
$v_{\mathrm{turb}}$ = 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 $v_{\mathrm{rad}}=-82$ 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 $\times$
1020 cm-2 is also included (Sembach et al., 1991). The model is then scaled by
1.25 $\times$ 10-20 in order to match observations to $\sim$10%. Finally, all
spectra were binned to 0.1 Å, $\sim$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 needed to match local continua.
Our subsequent column density measurements account for the uncertainty
inherent in the continuum-placement process.

Figure 2: A comparison of a section of the FUSE FUV spectrum with a TLUSTY
stellar model for the same spectral type, O9.7, which aided in the continuum
placement. The FUSE SiC2a spectrum from 914–944 Å is shown in black with a
stellar TLUSTY model shown in red. The TLUSTY model has $T_{\mathrm{eff}}$ =
26000 K, log $g$ = 3.0, $v_{\mathrm{turb}}$ = 10.0 km s-1, and is rotationally
broadened by 90 km s-1 in order to match the observed Si III 1300 Å triplets.
The model is reddened by an $E(B-V)$ = 0.09 using the mean Fitzpatrick & Massa
(2007) extinction curve and Lyman line absorption for log $N$(H I) = 20.78
(Sembach et al., 1991) is also applied. The flux was scaled by
1.25$\times$10-20 to provide an overall agreement with observations to
$\sim$10%.

###  3.2 H I absorption

Figure 3: 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 1$\sigma$ 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: 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 $+110\lesssim v\lesssim 125$ km s-1 has a slope near zero
within the margin of error, supporting the reality of the detection in both
ions.  Figure 5: Velocity profiles of the Ca II $\lambda\lambda$3934, 3969 and
Na I $\lambda\lambda$5891, 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.

The FUSE spectrum of HD 156359 (catalog ) covers almost the entire H I Lyman
series, from Lyman-$\beta$ 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 $\lambda\lambda$923, 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 $\lambda$917–944. H I $\lambda$923 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 $\lambda\lambda$926, 930, 937 with the
positions of the low-velocity absorption initially based on weak ISM lines. We
derive log $N_{\mathrm{HI}}$ = 15.54 $\pm$ 0.02 $\pm$ 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 $\lambda$923 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 H2 absorption. Fortunately, the FUSE LiF1 and LiF2 channels
provide us with an opportunity to model isolated H2 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 H2
decontamination and describe and illustrate this process in detail in Appendix
B.

###  3.3 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
3.2$\sigma$ ($EW$/$\sigma_{EW}$=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 $\lambda$1526
because it is unsaturated, unblended, and seen in a region with high signal to
noise. Although noisier, the profile of O I $\lambda$1302 is very similar to
Si II $\lambda$1526 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 $\lambda$1302 feature as real on the basis that:
1) the line is detected at 3.2$\sigma$ 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
$N_{\mathrm{OI}}$ = 12.43 $\pm$ 0.08 $\pm$ 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.

###  3.4 Low and intermediate ion absorption

We detected C II $\lambda$1334, Si II $\lambda\lambda$1190, 1193, 1304, 1526,
N II $\lambda$1083, Si III $\lambda$1206, and Fe III $\lambda$1122 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 $N_{\mathrm{SiII}}$ = 12.94$\pm$0.07 for the weakest unblended
line available at 1526 Å. We show that it is a good fit for
$\lambda\lambda$1190, 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 H2 $J$3 1084 Å line to other $J$3 lines
in regions with smooth continua reveals that the continuum must drop
significantly in this region, making it difficult to estimate log
$N_{\mathrm{NII}}$. 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
$\lambda$1144 and calculate a significance of 2.2$\sigma$ after including a
continuum placement error. We find a 3$\sigma$ limiting column density of log
$N_{\mathrm{FeII}}$ $\leq$ 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
$\lambda$1122, we detect an intermediate-velocity cloud (IVC) centered near
$+$86 km s-1 with log $N$ = 12.98$\pm$0.13.

The archival ESO FEROS optical spectrum covers the Ca II K and H and Na I D
lines, see Figure 5. Telluric H2O absorption lines are present in the velocity
range of the HVC, however, there is a non-detection of high-velocity Na I and
we find a 3$\sigma$ limiting column density limit of log
$N_{\mathrm{NaI}}<9.45$. We see absorption in Ca II across 6 pixels at 123.7
km s-1 and measure an AOD column density of log $N_{\mathrm{CaII}}$ =
10.68$\pm$0.01 for Ca II 3934. We also performed a simultaneous Voigt profile
fit to Ca II $\lambda$3934, 3969 and find log $N$ = 10.63$\pm$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.

###  3.5 High ion absorption

Figure 6: Velocity profiles of the high-ion absorption lines detected toward
HD 156359. The flux is shown in black and the 1$\sigma$ 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.

We detected C IV $\lambda\lambda$1548, 1550, Si IV $\lambda\lambda$1393, 1402,
and N V $\lambda\lambda$1238, 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 well-defined 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. (1991,
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 $N$ = $13.57\pm 0.08$ and
$13.09\pm 0.10$, respectively.

We observe high-velocity absorption in O VI 1031 and 1037 in the LiF1 channel
of the FUSE spectrum at $v_{\mathrm{LSR}}$ = 141 km s-1. However, we only
include $\lambda$1031 in our analysis, because $\lambda$1037 lies on the steep
blueward side of the O VI P Cygni profile and suffers from significant
blending with C II∗ 1037 and H2 $J$1 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) $\lambda$1031.915,
Cl I $\lambda$1031.507, and several lines of molecular hydrogen including H2
6–0 $P$(3) $\lambda$1031.192 and H2 6–0 $R$(4) $\lambda$1032.351. To gauge the
effect of contamination by HD 6–0 $R$(0) $\lambda$1031.915, we examined other
HD lines of similar oscillator strength that are isolated from interstellar
absorption, such as 5–0 $R$(0) $\lambda$1042.850, 7–0 $R$(0)
$\lambda$1021.460, and 8–0 $R$(0) $\lambda$1011.461. We detect no HD molecular
absorption in these lines and conclude that no subtraction of the HD 6–0
$R$(0) line at $\lambda$1031.915 from the O VI profile is necessary. For Cl I,
a small amount of absorption in Cl I $\lambda$1347 is present near 0 km s-1 in
the STIS spectrum and we simultaneously fit Cl I $\lambda$1031 to derive its
contribution to the O VI profile. For molecular hydrogen, H2 6–0 $R$(4)
$\lambda$1032.351 is the most relevant potential contaminant as it overlaps
with the high-velocity component in O VI $\lambda$ 1031\. We modeled other
isolated H2 $J4$ lines in the FUSE spectrum, e.g. 5–0 $R$(4) $\lambda$1044.543
and 4–0 $R$(4) $\lambda$1057.381 (see description in Appendix B). Through
simultaneous fitting of those lines with H2 6–0 $R$(4) $\lambda$1032 , we
account for the small amount of H2 present in the O VI high-velocity
absorption. A similar procedure was followed to account for contamination from
H2 6–0 $P$(3) $\lambda$1031.192 near $v_{\mathrm{LSR}}$ = $-$200 km s-1 using
the isolated H2 $J3$ lines 6–0 $R$(3) $\lambda$1028.986 and 5–0 $P$(3)
$\lambda$1043.503. The resulting Voigt profile fit and O VI column density are
shown in Figure 6 and Table 2.

##  4 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 $v_{\mathrm{LSR}}$ = +125 km s-1 toward HD 156359 (catalog ). 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.

Table 1: Measured Absorption Lines

Ion | $\lambda^{a}$ | $f$ | Dataset
---|---|---|---
| (Å) |  |
H I | 923.150 | 0.0022b | FUSE SiC2A
| 926.226 | 0.0032b | FUSE SiC2A
| 930.748 | 0.0048b | FUSE SiC2A
| 937.804 | 0.0078b | FUSE SiC2A
C II | 1334.532 | 0.1290c | STIS E140M
C IV | 1548.202 | 0.1900d | STIS E140M
| 1550.774 | 0.0948d | STIS E140M
N II | 1083.990 | 0.1110c | FUSE SiC2B
N V | 1238.821 | 0.1560e | STIS E140M
| 1242.804 | 0.0777e | STIS E140M
O I | 1302.168 | 0.0480c | STIS E140M
O VI | 1031.912 | 0.1330e | FUSE LiF1
Na I | 5891.583 | 0.6500g | FEROS
Si II | 1190.420 | 0.2560f | STIS E140M
| 1193.280 | 0.5440f | STIS E140M
| 1304.370 | 0.0910f | STIS E140M
| 1526.720 | 0.1440f | STIS E140M
Si III | 1206.510 | 1.6100g | STIS E140M
Si IV | 1393.760 | 0.5130g | STIS E140M
| 1402.770 | 0.2540g | STIS E140M
Ca II | 3934.770 | 0.6490h | FEROS
| 3969.590 | 0.3210h | FEROS
Fe II | 1144.926 | 0.0830i | FUSE LiF2
Fe III | 1122.524 | 0.0642j | FUSE LiF2

Note. — a All wavelengths are taken from Morton (2003). b Pal’chikov (1998), c
Froese Fischer & Tachiev (2004), d Yan et al. (1998), e Peach et al. (1988), f
Bautista et al. (2009), g Froese Fischer et al. (2006), h Safronova &
Safronova (2011) i Donnelly & Hibbert (2001), j Deb & Hibbert (2009)

Table 2: HD 156359 High-velocity absorption

Ion | $v_{\mathrm{LSR}}$ | $b$ | log $N$
---|---|---|---
| (km s-1) | (km s-1) | ($N$ in cm-2)
H I | 125.6 $\pm$ 0.9 | 18.9 $\pm$ 1.2 | 15.54 $\pm$ 0.05
O I | 125.0 $\pm$ 4.8 | … | 12.43 $\pm$ 0.11a
Na I | … | … |  $<$ 9.45b
C II | 120.9 $\pm$ 0.6 | 9.8 $\pm$ 1.4 | 13.83 $\pm$ 0.10
N II | … | … |  $<$ 13.98b
Si II | 118.5 $\pm$ 1.0 | 7.0 $\pm$ 1.6 | 12.94 $\pm$ 0.07
Ca II | 123.7 $\pm$ 1.0 | 2.3 $\pm$ 2.7 | 10.63 $\pm$ 0.10
Fe II | … | … |  $<$ 12.99b
Si III | 117.6 $\pm$ 1.4 | 6.9 $\pm$ 2.8 | 12.32 $\pm$ 0.16
Fe III | 129.3 $\pm$ 2.3 | 11.0 $\pm$ 3.6 | 12.82 $\pm$ 0.14
C IV | 133.3 $\pm$ 2.3 | 8.6 $\pm$ 3.6 | 12.60 $\pm$ 0.13
Si IV | 129.5 $\pm$ 1.5 | 5.7 $\pm$ 2.9 | 11.93 $\pm$ 0.15
N V | 130.2 $\pm$ 2.7 | 12.2 $\pm$ 4.0 | 12.64 $\pm$ 0.11
O VI | 141.3 $\pm$ 1.2 | 24.0 $\pm$ 1.8 | 13.65 $\pm$ 0.03

Note. — a Except where noted, all measurements presented in this table are
derived from Voigt-profile fitting, with the exception of O I, for which we
derive an apparent column density from the profile in the range $+110\lesssim
v\lesssim+125$ km s-1, see Figure 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 $\pm$ 0.08 $\pm$ 0.07 separately.

b This measurement is a non-detection and the 3$\sigma$ limiting column
density is given.

###  4.1 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 (catalog ) 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 ($n_{\mathrm{H}}$/cm-3) from $-3$ to 0, in order to explore
possible metallicities and densities. We determined the best-fit log
$n_{\mathrm{H}}$ 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 $n_{\mathrm{H}}$ determined from
the Si III/Si II ratio, we narrow the range of metallicities by identifying
which models are consistent with the observed log $N_{\mathrm{OI}}$ = 12.43
$\pm$ 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 $0.18\leq\mathrm{[O/H]}\leq 0.51$ and
densities from $-1.69\leq$ log ($n_{\mathrm{H}}$/cm-3) $\leq-1.37$ are
consistent with the data, giving a model best-fit of [O/H] =
0.36${}^{+0.15}_{-0.18}$ at log $n_{\mathrm{H}}$ = $-$1.53$\pm$-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 $\leq$ [C/H] $\leq$ 0.42 overlaps with the observed data, giving a model
best-fit of [C/H] = 0.30${}^{+0.12}_{-0.09}$ at log $n_{\mathrm{H}}$ =
$-$1.49${}^{+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.$\degree$9, $b$ = $+$47.$\degree$7 at $z$ = $+$5.3 kpc) with
[O/H] = $+$0.22 $\pm$ 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 $N_{\mathrm{HI}}$ = 16.50 $\pm$ 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 (catalog ) 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).

Table 3: HVC elemental abundances and depletions

Ion | [Xi/H I]a | IC(Xi)b | [X/H]c |  $\delta_{\mathrm{O}}$(X)d
---|---|---|---|---
O I | 0.20$\pm$0.10 |  $+$0.16${}^{+0.04}_{-0.07}$ | 0.36${}^{+0.11}_{-0.12}$ | 0
C II | 1.86$\pm$0.12 |  $-$1.59$\pm$0.04 | 0.27$\pm$0.13 |  $-$0.09$\pm$0.17
N II |  $<$ 2.60 |  $<-$1.67 |  $<$ 0.93 |  $<$ 0.57
Si II | 1.89$\pm$0.08 |  $-$1.86$\pm$0.03 | 0.03$\pm$0.09 |  $-$0.33$\pm$0.14
Fe II |  $<$ 1.95 |  $-$1.17${}^{+0.20}_{-0.12}$ |  $<$ 0.79 |  $<$ 0.43
Fe III | 1.78$\pm$0.15 |  $-$1.73$\pm$0.08 | 0.06$\pm$0.17 |  $-$0.30$\pm$0.20
Ca II | 0.75$\pm$0.11 |  $-$0.95$\pm$0.01 |  $-$0.20$\pm$0.11 |  $-$0.56$\pm$0.16

Note. —

a Ion abundance [Xi/H I] = log (Xi/H I)HVC $-$ log (X/H)⊙ where Xi 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(Xi) = [X/H]model $-$ [Xi/H I], derived from the
difference between the model elemental abundance and the observed ion
abundance.

c Ionization-corrected elemental abundance [X/H] = [Xi/H I] + IC(Xi). Also
referred to as gas-phase abundance.

d $\delta_{\mathrm{O}}$(X) $\equiv$ [X/O] = [X/H] $-$ [O/H] is the depletion
of element X relative to oxygen.

Figure 7: Cloudy photoionization models of the HVC toward HD 156359 (catalog )
that explore a parameter space in metallicity and number density, $-0.50\leq$
[Z/H] $\leq+1.50$ and $-3\leq$ log ($n_{\mathrm{H}}$/cm-3) $\leq 0$,
respectively. Top panel: Model results for log
$N_{\mathrm{SiIII}}$/$N_{\mathrm{SiII}}$ for the set metallicities are shown
as solid colored curves. The dashed horizontal turquoise line is the observed
ratio of $-$0.62$\pm$0.17. The best-fit log $n_{\mathrm{H}}$ 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 $N_{\mathrm{OI}}$ for each metallicity are
plotted against log $n_{\mathrm{H}}$, where the values derived from log
$N_{\mathrm{SiIII}}$/$N_{\mathrm{SiII}}$ are marked with vertical lines and
markers. The horizontal turquoise line and bar show the observed log
$N_{\mathrm{OI}}$ = 12.43 $\pm$ 0.11. The model data corresponding to the
metallicity region from $+0.18\leq$[Z/H] $\leq+0.51$ and $-1.69\leq$ log
$n_{\mathrm{H}}$ $\leq-1.37$ overlap with the observed data. The model best-
fit values are [Z/H] = $+$0.36${}^{+0.15}_{-0.18}$ and log $n_{\mathrm{H}}$ =
$-1.53\pm 0.16$. Bottom right panel: Same as for the bottom left panel, but
for C II. The horizontal turquoise line and bar show observed log
$N_{\mathrm{CII}}$ = 13.83 $\pm$ 0.10. The model data corresponding to the
metallicity region from $+0.21\leq$ [Z/H] $\leq+0.42$ and $-1.59\leq$ log
$n_{\mathrm{H}}$ $\leq-1.40$ overlap with the observed data. The model best
fit values are [Z/H]= $+$0.30${}^{+0.12}_{-0.09}$ and log $n_{\mathrm{H}}$ =
$-1.49^{+0.09}_{-0.10}$. 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 $n_{\mathrm{H}}$ from our Cloudy model to the HD 156359 (catalog )
HVC. The model uses our best-fit [Z/H]=$+$0.36. The vertical line marks the
best-fit log ($n_{\rm H}$/cm-3)=$-$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 $\delta$(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) from Jenkins (2009). A slight offset is applied in the
$x$-direction of each element for distinction. 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] =
$+$0.36 determined from our O I analysis. These model predictions for each ion
(colored dotted curves) have been scaled by the dust depletions required to
match the observed values, where the observed column densities are indicated
by circle markers at the best-fit log ($n_{\mathrm{H}}$/cm-3)=$-$1.53 (black
vertical line).

###  4.2 Ionization corrections

In ISM abundance work, [O I/H I] is often considered a robust indicator of
elemental abundance [O/H], since (1) charge-exchange reactions closely couple
the 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 the HVC at +125 km s-1 toward HD 156359
(catalog ), log $N$(H I) is only 15.54$\pm$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.,

$IC(X^{i})=[\text{X/H}]-[\text{X}^{i}/\textsc{H i}].$ (1)

The resulting ion abundances, ionization corrections, and ionization-corrected
elemental abundances are given in Table 3. The relationships between the
ionization corrections and the hydrogen number density for [Z/H] = 0.36 are
shown in the top panel of Figure 8 and the ionization-corrected abundances are
shown in the middle panel of Figure 8.

The positive ionization correction of $IC$(O I) = $0.16^{+0.04}_{-0.07}$ from
the photoionization modeling results in a high gas-phase oxygen abundance of
$+0.36$. We calculate gas-phase abundances of [C/H] = 0.27$\pm$0.13, [N/H] =
$\leq$ 0.93, [Si/H] = 0.03$\pm$0.09, [Fe/H] = 0.06$\pm$0.17, and [Ca/H] =
$-$0.20$\pm$0.11 (see Table 3).

###  4.3 Dust depletion effects

Following convention, we define the depletion $\delta_{\mathrm{O}}$(X) of each
refractory element X as the ionization-corrected abundances of that element
compared to the ionization-corrected oxygen abundance, i.e.,

$\delta_{\mathrm{O}}(\text{X})\equiv[\text{X/O}]=[\text{X/H}]-[\text{O/H}],$
(2)

where oxygen represents an undepleted (or lightly depleted) volatile element.
In this formalism, a negative value of $\delta_{\mathrm{O}}$(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
$\delta_{\mathrm{O}}$(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 (catalog ), we find a low value of
$\delta_{\mathrm{O}}$(C)=$-$0.09$\pm$0.17, i.e. carbon shows no significant
depletion. This is consistent with the low values of [C/O] measured in
Galactic ISM gas (Jenkins, 2009). For the other low ions, we find a 3$\sigma$
upper limit for the nitrogen depletion $\delta_{\mathrm{O}}$(N)$\leq+0.57$ and
low depletions for silicon, iron, and calcium of
$\delta_{\mathrm{O}}$(Si)=$-$0.33$\pm$0.14$,\delta_{\mathrm{O}}$(Fe)=$-$0.30$\pm$0.20,
and $\delta_{\mathrm{O}}$(Ca)=$-$0.56$\pm$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 $n_{\mathrm{H}}$= $-$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 3$\sigma$ upper limit of log $N$(Na I)/$N$(Ca II)
$\leq$ $-$1.2 in the $+124$ km s-1 HVC toward HD 156359 (catalog ) 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.

##  5 Results: the Highly Ionized Gas

The HVC toward HD 156359 (catalog ) 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 a comparison of the normalized optical
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 $\approx$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.

###  5.1 Collisional Ionization Modeling

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 $\lambda$1548 and Si IV
$\lambda$1393 display similar line shapes in the velocity region of the HVC
from $112\lesssim v\lesssim 140$ km s-1, indicated with cyan shading. Middle
panel: C IV $\lambda$1548 and O VI $\lambda$1031 are compared, where the
center of the HVC absorption for O VI is shifted to a higher velocity at
$v\approx 141$ km s-1. Bottom panel: N V $\lambda$1238 and O VI $\lambda$1031
are compared, where the range of velocities attribute to absorption by O VI
are indicated in light pink. Figure 11: Comparison of observed high-ion
column-density ratios with predictions from the collisional ionization models
of Gnat & Sternberg (2007). The observed ratios in the HD 156359 (catalog )
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 time-dependent 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 $T$ = 4.2–4.4 K.

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 $\sim$0.7 and 1.0 dex respectively, the N
V and O VI differ by orders of magnitudes. 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 $T\gtrsim$ 105 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
$N_{\mathrm{CIV}}$/$N_{\mathrm{SiIV}}$ ratio and another explaining
$N_{\mathrm{OVI}}$/$N_{\mathrm{NV}}$. 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 $N_{\mathrm{CIV}}$/$N_{\mathrm{SiIV}}$ and
$N_{\mathrm{OVI}}$/$N_{\mathrm{NV}}$ 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 $T$ = $10^{4.2-4.4}$ K,
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 $T$=104.9 K.
We are inclined to adopt the non-equilibrium (lower-temperature) solution, as
plasma near $T$=105 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 non-equilibrium state. For the O VI/N V phase, a single
temperature solution for the observed log
$N_{\mathrm{OVI}}$/$N_{\mathrm{NV}}$=1.01 is found at $T$=105.4 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 (catalog ) as containing collisionally-ionized gas at two
temperatures: a cooler phase seen in C IV and Si IV at $T$ = $10^{4.2-4.4}$ K,
and a hotter phase seen in N V and O VI at $T$=105.4 K.

##  6 Summary

Using archival FUSE, HST STIS, and ESO FEROS spectra, we have analyzed the
chemical composition of the HVC near $+$125 km s-1 toward HD 156359 (catalog
), a massive star lying 9 kpc away toward the Galactic Center. The sight line
passes less than 1$\degree$ 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. 1.

We determined an H I column density measurement of log $N$(H I) = $15.54\pm
0.05$ in the HVC using unsaturated Lyman series absorption lines.

  2. 2.

We measured a super-solar O I abundance of [O I/H I] = 0.20$\pm$0.11 in the
HVC. After applying an ionization correction, we derive an oxygen abundance of
[O/H] = 0.36$\pm$0.12, indicating the cloud is enriched to 2.3 times the solar
level. This abundance is consistent with super solar O abundances of H II
regions measured in the inner Galaxy (Wenger et al., 2019; Arellano-Córdova et
al., 2021).

  3. 3.

The ionization-corrected carbon abundance [C/H] = 0.27$\pm$0.13 is consistent
with the oxygen abundance. This indicates that carbon shows low dust depletion
relative to oxygen, consistent with the pattern seen by Jenkins (2009), as
shown in Figure 8.

  4. 4.

A low level of depletion is inferred from the silicon, iron, and calcium
lines, with [Si/O] = $-$0.33 $\pm$ 0.14, [Fe/O] = $-$0.30 $\pm$ 0.20, and
[Ca/O] = $-$0.56 $\pm$ 0.16. The Na I/Ca II ratio in the HVC is $\leq-1.2$,
which is lower than is typically observed at low velocities, consistent with
the trend known as the Routly-Spitzer effect.

  5. 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 $N_{\mathrm{CIV}}$/$N_{\mathrm{SiIV}}$ and
$N_{\mathrm{OVI}}$/$N_{\mathrm{NV}}$ ratios, with a cooler phase seen in C IV
and Si IV at $T$ = $10^{4.2-4.4}$ K, and a hotter phase seen in N V and O VI
at $T$=105.4 K.

The high metallicity, low depletion, complex high-ion absorption, and high
positive velocity of the HVC toward HD 156359 (catalog ) 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 $\sim-10$
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 $\sim$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 (catalog ) 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 http://dx.doi.org/10.17909/rrnk-3e58
(catalog 10.17909/rrnk-3e58).

##  Appendix A Global Continuum Fitting

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 12, 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 12. A separate local continuum was applied to the depressed region
resulting from stellar line blanketing and which contains H I $\lambda$923\.
Using the global continuum fit (which has a slightly higher continuum
placement than the local fits) results in log $N_{\mathrm{HI}}$ =
15.69$\pm$0.03$\pm$ 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 $\Delta$log
$N_{\mathrm{HI}}$ = 0.15 dex between the H I column densities derived by the
global and local continuum methods.

Figure 12: 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 1$\sigma$
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 $\sim$920–925 Å.

##  Appendix B Decontamination of Molecular Hydrogen

The FUSE LiF1 and LiF2 channels provide us with an opportunity to model
isolated H2 lines. This effort is important because it enables us to account
for the Milky Way’s molecular contribution to the interstellar absorption
lines. In particular, contamination from Milky Way H2 affected the higher
order Lyman series H I in the SiC2A channel and O VI in the LiF1 channel. We
performed simultaneous single-component Voigt profile fits to multiple H2
transitions from rotational levels $J$ = 0–6, allowing both the velocity and
$b$-value to vary freely for the stronger $J$ = 0–4 transitions. The
$b$-values of the significantly weaker $J$ = 5,6 transitions were fixed to the
$b$-value of the $J$ = 4 transition in order to deter unrealistic column
densities for these fainter lines. These results of these fittings are shown
in Figure 13.

Figure 13: We performed Voigt profile fits to vibrationally and rotationally
excited H2 lines from the FUSE LiF1 and LiF2 channels in order to account for
molecular blending with H I. In each panel, the flux and 1$\sigma$ error in
the flux are shown in black and blue, respectively. Voigt profile fits to the
data are shown in green. Green vertical lines mark the positions of the H2
absorption lines. Top panel: LiF1 channel H2 transitions from upper
vibrational levels 4 and 5, spanning rotational levels $J=0-6$. Fits to Ar I
$\lambda$1048 and Fe II $\lambda$1055 are included. Bottom left panel: LiF1
channel H2 transitions from upper vibrational levels 2 and 3, including
rotational levels $J=0,1,2,5,6$. Bottom right panel: LiF2 channel H2
transitions from upper vibrational levels 1 and 2, including rotational levels
$J=0,1,2,5$.

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