# A Reexamination of Phosphorus and Chlorine Depletions in the Diffuse
Interstellar Medium111Based on observations made with the NASA/ESA _Hubble
Space Telescope_ and the _Far Ultraviolet Spectroscopic Explorer_ mission,
obtained from the MAST data archive at the Space Telescope Science Institute.
STScI is operated by the Association of Universities for Research in
Astronomy, Inc., under NASA contract NAS5-26555.

Adam M. Ritchey  Eureka Scientific, 2452 Delmer Street, Suite 100, Oakland, CA
96402, USA<EMAIL_ADDRESS>J. M. Brown  Department of Physics and
Astronomy, University of Toledo, Toledo, OH 43606, USA  S. R. Federman
Department of Physics and Astronomy, University of Toledo, Toledo, OH 43606,
USA  Paule Sonnentrucker  Space Telescope Science Institute, 3700 San Martin
Dr., Baltimore, MD 21218, USA

###### Abstract

We present a comprehensive examination of interstellar P and Cl abundances
based on an analysis of archival spectra acquired with the Space Telescope
Imaging Spectrograph of the Hubble Space Telescope and the Far Ultraviolet
Spectroscopic Explorer. Column densities of P ii, Cl i, and Cl ii are
determined for a combined sample of 107 sight lines probing diffuse atomic and
molecular gas in the local Galactic interstellar medium (ISM). We reevaluate
the nearly linear relationship between the column densities of Cl i and H2,
which arises from the rapid conversion of Cl+ to Cl0 in regions where H2 is
abundant. Using the observed total gas-phase P and Cl abundances, we derive
depletion parameters for these elements, adopting the methodology of Jenkins.
We find that both P and Cl are essentially undepleted along sight lines
showing the lowest overall depletions. Increasingly severe depletions of P are
seen along molecule-rich sight lines. In contrast, gas-phase Cl abundances
show no systematic variation with molecular hydrogen fraction. However,
enhanced Cl (and P) depletion rates are found for a subset of sight lines
showing elevated levels of Cl ionization. An analysis of neutral chlorine
fractions yields estimates for the amount of atomic hydrogen associated with
the H2-bearing gas in each direction. These results indicate that the
molecular fraction in the H2-bearing gas is at least 10% for all sight lines
with $\log N({\rm H}_{2})\gtrsim 18$ and that the gas is essentially fully
molecular at $\log N({\rm H}_{2})\approx 21$.

interstellar medium — interstellar abundances — diffuse interstellar clouds —
interstellar dust

††facilities: HST(STIS), FUSE, Copernicus††software: ISMOD (Sheffer et al.,
2008), STSDAS, H2GUI (Tumlinson et al., 2002), LTOOLS

##  1 INTRODUCTION

Dust plays an important role in the evolution of galaxies due to the impact it
has on the physics and chemistry of the interstellar medium (ISM). However,
questions remain regarding the dominant dust production mechanisms, both in
the present day Milky Way and in galaxies at high redshift. Efficient dust
production occurs in the extended outer atmospheres of evolved stars (e.g.,
Gobrecht et al., 2016) and in the cooling ejecta of supernovae (SNe; e.g.,
Sarangi & Cherchneff, 2015). However, the timescale for dust production by
stars is longer than the timescale for destruction of the grains by SN shocks
(Dwek & Scalo, 1980; Jones et al., 1994; Slavin et al., 2015). Moreover, the
large observed depletions of refractory elements in cold, diffuse clouds
(where $\sim$95% of the Mg and Si atoms and $\sim$99% of the Fe atoms are
locked up in grains; e.g., Jenkins, 2009) are inconsistent with the levels of
condensation expected for material returned to the ISM by evolved stars and
SNe (e.g., Dwek, 2016). These considerations have generally been interpreted
as indications that dust grains must grow through the accretion of refractory
elements in cold diffuse interstellar clouds.

Direct constraints on the growth of dust grains in the diffuse ISM come from
detailed knowledge of the depletions of different elements from the gas-phase
and how those depletions change with changing environmental conditions.
Jenkins (2009) compiled information on interstellar depletions for 17 elements
along 243 sight lines sampling diffuse atomic and molecular gas in the solar
vicinity. Jenkins (2009) developed a unified framework for examining gas-phase
element depletions predicated on the empirical observation that, while
different elements exhibit different degrees of depletion, the depletions of
most elements tend to increase in a systematic way as the overall strength of
depletions increases from one sight line to the next. Ritchey et al. (2018)
extended that framework to include several rare neutron-capture elements (and
the light element B) yielding a more complete picture of gas-phase depletions
for elements with low-to-moderate condensation temperatures.

To quantify the overall strength of the depletions seen in a given direction,
Jenkins (2009) defined the sight line depletion strength factor $F_{*}$. Sight
lines showing strong depletions (such as those that characterize relatively
dense and/or cold H2-rich gas in the Galactic disk where grain growth occurs
via cold accretion) have depletion factors near $F_{*}=1$, while sight lines
showing only very modest depletions (characteristic of warm and/or low-density
H2-poor gas where substantial grain growth has not yet occurred or where the
grains have been destroyed or disrupted by shocks) have values of $F_{*}$
closer to zero. The observed depletions at $F_{*}=0$ are probably the best
representation that we have of the composition of grain cores that emerge from
evolved low and high-mass stars and core-collapse SNe.

Most elements with low-to-moderate condensation temperatures ($T_{C}\lesssim
800$ K) have depletions at $F_{*}=0$ that are consistent with zero (see Figure
21 in Ritchey et al., 2018). For elements with higher condensation
temperatures, a fairly well-defined trend emerges of increasing depletion with
increasing $T_{C}$, particularly for Ga, Ge, Cu, Mn, Cr, Fe, Ni, and Ti.
Ritchey et al. (2018) suggest that this trend constitutes a dust condensation
sequence and that the temperature where the sequence seems to terminate
($\sim$800 K) is related to the dust formation temperature. This idea is
corroborated by observations of dust shells around asymptotic giant branch
(AGB) stars, which are found to have temperatures in the range 800 to 1100 K
(Gail et al., 2013). However, while the trend of increasing depletion with
increasing $T_{C}$ seems convincing for the elements mentioned above, there
are also glaring exceptions to the trend for the elements Si and Mg but
especially for Cl, As, and P. These latter three elements appear to have gas-
phase abundances at $F_{*}=0$ that are significantly higher than the
corresponding solar values, which seems particularly unusual since the
elements have moderate-to-high condensation temperatures.

One possible factor that could be influencing the unusual Cl and P depletions
is that the column density measurements for the dominant singly-ionized
species may not be very accurate. All of the Cl ii measurements (and over half
of the P ii results) compiled by Jenkins (2009) are from observations obtained
by the Copernicus satellite (e.g., Jenkins et al., 1986). More precise
measurements of interstellar column densities may be derived from spectra
acquired using the higher resolution and/or higher sensitivity instruments of
the Hubble Space Telescope (HST) and Far Ultraviolet Spectroscopic Explorer
(FUSE).

An even more important factor in the case of P depletions is that the
transition oscillator strengths ($f$-values) used by Jenkins (2009) and others
to derive P ii column densities are in need of revision. Jenkins (2009)
adopted theoretical P ii $f$-values originally obtained by Hibbert (1988).
However, more precise experimentally-determined $f$-values for the important
$\lambda 1152$ and $\lambda 1301$ transitions are now available (Federman et
al., 2007; Brown et al., 2018). These new experimental results indicate that
the P ii column densities obtained from the $\lambda 1152$ and $\lambda 1301$
transitions should be revised downward by 0.045 dex and 0.188 dex,
respectively, in agreement with recent theoretical calculations (Tayal, 2003;
Froese Fischer et al., 2006). Furthermore, these more recent theoretical
efforts suggest that a downward revision of $\sim$0.4 dex is required for
column densities derived using the older Hibbert (1988) value for P ii
$\lambda 1532$ (e.g., Lebouteiller et al., 2005; Jenkins, 2009).

A major contributing factor to the unusual Cl depletions reported by Jenkins
(2009) is that he considered only Cl ii column densities when deriving total
Cl abundances, ignoring any contribution from Cl i. However, because Cl+
reacts rapidly with H2 to form HCl+, which leads (eventually) to Cl0 and H0,
Cl will be predominantly neutral in regions where H2 is abundant (Jura, 1974;
Jura & York, 1978). Indeed, there are a number of sight lines where previous
studies indicate that $N({\rm Cl~{}\textsc{i}})>N({\rm Cl~{}\textsc{ii}})$
(e.g., Harris & Bromage, 1984; Jenkins et al., 1986; Moomey et al., 2012) and
these sight lines tend to have high molecular hydrogen fractions. Since these
same sight lines will likely have high values of $F_{*}$ as well, the
depletion trend for Cl calculated by Jenkins (2009) is too steep and the
extrapolated value of [Cl/H] at $F_{*}=0$ is too high.

In this investigation, we seek to present a definitive analysis of
interstellar P and Cl depletions based on high-quality archival HST and FUSE
data. The selection of sight lines and the steps involved in processing the
archival data are described in Section 2. The procedures used to obtain column
densities of P ii, Cl i, and Cl ii for the sight lines in our survey are
explained in Section 3. In Section 4.1, we reexamine the relationships that
exist between various Cl and H species. In Section 4.2, the column densities
derived in Section 3 are used to evaluate the depletion trends for P and Cl,
with some additional insights discussed in Section 4.3. In Section 4.4, we use
the observed neutral chlorine fractions to estimate the amount of atomic
hydrogen associated with the molecular gas in each direction. The implications
of our results are discussed in Section 5. Our summary and conclusions are
presented in Section 6. Two appendices compile a set of chlorine measurements
from Copernicus observations and present an analysis of fluorine depletions
from the literature.

##  2 OBSERVATIONS AND DATA PROCESSING

Since the ionization potential of P i (10.5 eV) is below that of neutral
hydrogen (13.6 eV), most P in neutral, diffuse clouds is singly ionized. The
same would be true of Cl, since its first ionization potential is 13.0 eV,
except that Cl+ reacts exothermically with H2, which leads to the conversion
of Cl+ to Cl0 in regions where H2 is optically thick (e.g., Neufeld & Wolfire,
2009). Thus, while P ii column densities are sufficient for studies of
interstellar P abundances, column densities of both Cl i and Cl ii are
required for an evaluation of total interstellar Cl abundances.

Most previous studies of interstellar P have relied on either the P ii
$\lambda 1152$ transition or the $\lambda 1301$ transition. However, the
former is often strongly saturated and the latter may be blended with the wing
of the nearby strong O i $\lambda 1302$ absorption feature. Fewer studies make
use of the weak P ii $\lambda 1532$ absorption line, which is unblended and
typically unaffected by saturation, yet still strong enough to yield reliable
measurements for most sight lines. Similarly, while the Cl i $\lambda 1347$
line is almost always present in stellar spectra probing diffuse molecular
gas, this feature is very often affected by saturation in the line core,
making reliable column density determinations difficult. The nearby Cl i line
at 1379.5 Å is much weaker and yet (again) is often strong enough to be
reliably measured.

Since the goal of our survey is to accurately deduce the depletion
characteristics of Cl and P, minimizing the associated uncertainties, we use
measurements of the P ii $\lambda 1532$ and Cl i $\lambda 1379$ absorption
features, from high or medium resolution STIS spectra, wherever possible. We
include measurements of the P ii $\lambda 1301$ line only in cases where there
is no apparent blending with O i $\lambda 1302$. The strong Cl i $\lambda
1347$ line is included in our survey. However, in most cases, a weaker Cl i
line (typically $\lambda 1379$) is used to help constrain the total column
density. For some sight lines, we use observations of the Cl i $\lambda 1004$,
$\lambda 1094$, or $\lambda 1097$ transition, available from FUSE data, to
constrain the total Cl i column density (see Section 3.2). The column density
of Cl ii is derived from observations of the Cl ii line at 1071.0 Å, available
from FUSE spectra.

Table 1: Stellar Data for the Combined Sample Star | Name | Sp. Type | $V$ | $E(\bv)$ | $l$ | $b$ |  $d$aaDistances are based on Gaia EDR3 parallax measurements (Bailer-Jones et al., 2021). | $z$
---|---|---|---|---|---|---|---|---
|  |  | (mag) | (mag) | (deg) | (deg) | (kpc) | (kpc)
HD 108 | … | O8fp | 7.40 | 0.43 | 117.93 | $+1.25$ | $1.93^{+0.08}_{-0.09}$ | $+0.042$
HD 1383 | … | B1II | 7.63 | 0.47 | 119.02 | $-0.89$ | $2.50^{+0.18}_{-0.13}$ | $-0.039$
HD 3827 | … | B0.7V | 7.95 | 0.02 | 120.79 | $-23.23$ | $1.67^{+0.14}_{-0.15}$ | $-0.658$
HD 12323 | … | ON9.2V | 8.92 | 0.23 | 132.91 | $-5.87$ | $2.44^{+0.23}_{-0.22}$ | $-0.250$
HD 13268 | … | ON8.5IIIn | 8.18 | 0.36 | 133.96 | $-4.99$ | $1.77^{+0.09}_{-0.07}$ | $-0.154$
HD 13745 | V354 Per | O9.7IIn | 7.90 | 0.46 | 134.58 | $-4.96$ | $2.28^{+0.15}_{-0.09}$ | $-0.197$
HD 14434 | … | O5.5Vnfp | 8.49 | 0.48 | 135.08 | $-3.82$ | $2.24^{+0.10}_{-0.09}$ | $-0.150$
HD 15137 | … | O9.5II-IIIn | 7.86 | 0.35 | 137.46 | $-7.58$ | $2.05^{+0.17}_{-0.13}$ | $-0.271$
HD 23478 | … | B3IV | 6.67 | 0.27 | 160.76 | $-17.42$ | $0.249^{+0.004}_{-0.003}$ | $-0.074$
HD 24190 | … | B2Vn | 7.45 | 0.30 | 160.39 | $-15.18$ | $0.375^{+0.005}_{-0.005}$ | $-0.098$
HD 24534 | X Per | O9.5III | 6.72 | 0.59 | 163.08 | $-17.14$ | $0.596^{+0.017}_{-0.014}$ | $-0.176$
HD 25443 | … | B0.5III | 6.77 | 0.54 | 143.68 | $+7.35$ | $1.13^{+0.07}_{-0.06}$ | $+0.145$
HD 37903 | … | B1.5V | 7.83 | 0.35 | 206.85 | $-16.54$ | $0.394^{+0.003}_{-0.004}$ | $-0.112$
HD 41161 | … | O8Vn | 6.76 | 0.21 | 164.97 | $+12.89$ | $1.43^{+0.12}_{-0.10}$ | $+0.319$
HD 46223 | … | O4Vf | 7.28 | 0.54 | 206.44 | $-2.07$ | $1.39^{+0.05}_{-0.06}$ | $-0.050$
HD 52266 | … | O9.5IIIn | 7.23 | 0.29 | 219.13 | $-0.68$ | $1.35^{+0.06}_{-0.07}$ | $-0.016$
HD 53975 | … | O7.5Vz | 6.50 | 0.21 | 225.68 | $-2.32$ | $1.12^{+0.07}_{-0.06}$ | $-0.045$
HD 63005 | … | O7Vf | 9.13 | 0.27 | 242.47 | $-0.93$ | $3.72^{+0.60}_{-0.50}$ | $-0.060$
HD 66788 | … | O8V | 9.83 | 0.22 | 245.43 | $+2.05$ | $3.31^{+0.32}_{-0.30}$ | $+0.119$
HD 72754 | FY Vel | B2I:pe | 8.88 | 0.36 | 266.83 | $-5.82$ | $1.62^{+0.05}_{-0.05}$ | $-0.165$
HD 73882 | NX Vel | O8.5IV | 7.19 | 0.70 | 260.18 | $+0.64$ | $0.737^{+0.031}_{-0.030}$ | $+0.008$
HD 75309 | … | B1IIp | 7.84 | 0.29 | 265.86 | $-1.90$ | $1.82^{+0.12}_{-0.13}$ | $-0.060$
HD 79186 | GX Vel | B5Ia | 5.00 | 0.40 | 267.36 | $+2.25$ | $1.81^{+0.34}_{-0.23}$ | $+0.071$
HD 88115 | … | B1.5Iin | 9.36 | 0.16 | 285.32 | $-5.53$ | $2.53^{+0.16}_{-0.17}$ | $-0.243$
HD 89137 | … | ON9.7IIn | 7.97 | 0.27 | 279.69 | $+4.45$ | $2.44^{+0.20}_{-0.14}$ | $+0.189$
HD 90087 | … | O9.2III | 8.92 | 0.28 | 285.16 | $-2.13$ | $2.19^{+0.11}_{-0.12}$ | $-0.082$
HD 91597 | … | B1IIIne | 9.92 | 0.30 | 286.86 | $-2.37$ | $3.90^{+0.26}_{-0.24}$ | $-0.161$
HD 91651 | … | ON9.5IIIn | 9.52 | 0.28 | 286.55 | $-1.72$ | $1.85^{+0.07}_{-0.07}$ | $-0.055$
HD 91824 | … | O7Vfz | 8.14 | 0.24 | 285.70 | $+0.07$ | $1.83^{+0.08}_{-0.08}$ | $+0.002$
HD 91983 | … | B1III | 8.55 | 0.29 | 285.88 | $+0.05$ | $2.40^{+0.15}_{-0.14}$ | $+0.002$
HD 92554 | … | O9.5IIn | 10.15 | 0.39 | 287.60 | $-2.02$ | $4.04^{+0.31}_{-0.28}$ | $-0.142$
HD 93129 | … | O2If* | 6.90 | 0.48 | 287.41 | $-0.57$ | $2.41^{+0.10}_{-0.10}$ | $-0.024$
HD 93205 | V560 Car | O3.5V | 7.75 | 0.38 | 287.57 | $-0.71$ | $2.25^{+0.12}_{-0.11}$ | $-0.028$
HD 93222 | … | O7IIIf | 8.10 | 0.36 | 287.74 | $-1.02$ | $2.41^{+0.14}_{-0.15}$ | $-0.043$
HD 93843 | … | O5IIIf | 7.33 | 0.27 | 288.24 | $-0.90$ | $2.37^{+0.19}_{-0.14}$ | $-0.037$
HD 94493 | … | B1Ib | 7.59 | 0.23 | 289.01 | $-1.18$ | $2.15^{+0.14}_{-0.12}$ | $-0.044$
HD 97175 | … | B0.5III | 9.20 | 0.18 | 294.53 | $-9.17$ | $2.48^{+0.18}_{-0.14}$ | $-0.396$
HD 99857 | … | B0.5Ib | 7.49 | 0.35 | 294.78 | $-4.94$ | $1.80^{+0.08}_{-0.07}$ | $-0.155$
HD 99890 | … | B0IIIn | 9.26 | 0.24 | 291.75 | $+4.43$ | $2.53^{+0.18}_{-0.17}$ | $+0.195$
HD 100199 | … | B0IIIn | 8.17 | 0.30 | 293.94 | $-1.49$ | $1.73^{+0.14}_{-0.11}$ | $-0.045$
HD 101190 | … | O6IVf | 7.33 | 0.37 | 294.78 | $-1.49$ | $2.36^{+0.13}_{-0.12}$ | $-0.061$
HD 103779 | … | B0.5Iab | 7.22 | 0.21 | 296.85 | $-1.02$ | $2.02^{+0.10}_{-0.08}$ | $-0.036$
HD 104705 | DF Cru | B0Ib | 9.11 | 0.23 | 297.45 | $-0.34$ | $1.94^{+0.16}_{-0.15}$ | $-0.011$
HD 108639 | … | B0.2III | 8.57 | 0.37 | 300.22 | $+1.95$ | $1.98^{+0.10}_{-0.11}$ | $+0.067$
HD 109399 | … | B0.7II | 7.67 | 0.21 | 301.71 | $-9.88$ | $2.33^{+0.13}_{-0.14}$ | $-0.400$
HD 114886 | … | O9III | 6.89 | 0.40 | 305.52 | $-0.83$ | $1.83^{+0.95}_{-0.78}$ | $-0.026$
HD 115455 | … | O8IIIf | 7.97 | 0.47 | 306.06 | $+0.22$ | $1.80^{+0.09}_{-0.07}$ | $+0.007$
HD 116852 | … | O8.5II-IIIf | 8.47 | 0.21 | 304.88 | $-16.13$ | $3.42^{+0.39}_{-0.32}$ | $-0.949$
HD 121968 | … | B1V | 10.26 | 0.07 | 333.97 | $+55.84$ | $3.94^{+0.91}_{-0.63}$ | $+3.26$
HD 122879 | … | B0Ia | 6.50 | 0.36 | 312.26 | $+1.79$ | $2.22^{+0.16}_{-0.13}$ | $+0.069$
HD 124314 | … | O6IVnf | 6.64 | 0.53 | 312.67 | $-0.42$ | $1.61^{+0.11}_{-0.10}$ | $-0.012$
HD 137595 | … | B3Vn | 7.49 | 0.25 | 336.72 | $+18.86$ | $0.751^{+0.021}_{-0.022}$ | $+0.243$
HD 144965 | … | B2Vne | 7.11 | 0.35 | 339.04 | $+8.42$ | $0.258^{+0.002}_{-0.002}$ | $+0.038$
HD 147683 | V760 Sco | B4V+B4V | 7.05 | 0.39 | 344.86 | $+10.09$ | $0.290^{+0.002}_{-0.002}$ | $+0.051$
HD 147888 |  $\rho$ Oph D | B3V | 6.74 | 0.47 | 353.65 | $+17.71$ | $0.124^{+0.006}_{-0.006}$ | $+0.038$
HD 147933 |  $\rho$ Oph A | B2V | 5.05 | 0.45 | 353.69 | $+17.69$ | $0.137^{+0.003}_{-0.003}$ | $+0.042$
HD 148422 | … | B1Ia | 8.65 | 0.29 | 329.92 | $-5.60$ | $4.18^{+0.52}_{-0.31}$ | $-0.408$
HD 148937 | … | O6f?p | 6.71 | 0.65 | 336.37 | $-0.22$ | $1.15^{+0.03}_{-0.03}$ | $-0.004$
HD 151805 | … | B1Ib | 8.86 | 0.43 | 343.20 | $+1.59$ | $1.51^{+0.07}_{-0.06}$ | $+0.042$
HD 152590 | V1297 Sco | O7.5Vz | 9.29 | 0.48 | 344.84 | $+1.83$ | $1.68^{+0.08}_{-0.06}$ | $+0.053$
HD 156359 | … | B0Ia | 9.72 | 0.14 | 328.68 | $-14.52$ | $8.98^{+3.14}_{-2.11}$ | $-2.25$
HD 157857 | … | O6.5IIf | 7.78 | 0.43 | 12.97 | $+13.31$ | $2.22^{+0.20}_{-0.13}$ | $+0.511$
HD 163522 | … | B1Ia | 8.43 | 0.19 | 349.57 | $-9.09$ | $4.01^{+0.56}_{-0.47}$ | $-0.633$
HD 165246 | … | O8Vn | 7.60 | 0.38 | 6.40 | $-1.56$ | $1.19^{+0.04}_{-0.05}$ | $-0.032$
HD 165955 | … | B3Vn | 9.59 | 0.15 | 357.41 | $-7.43$ | $1.47^{+0.12}_{-0.09}$ | $-0.190$
HD 167402 | … | B0Ib | 9.03 | 0.21 | 2.26 | $-6.39$ | $4.94^{+0.83}_{-0.73}$ | $-0.550$
HD 168076 | … | O4IIIf | 8.25 | 0.78 | 16.94 | $+0.84$ | $1.65^{+1.51}_{-0.67}$ | $+0.024$
HD 168941 | … | O9.5IVp | 9.37 | 0.24 | 5.82 | $-6.31$ | $4.00^{+0.60}_{-0.53}$ | $-0.440$
HD 170740 | … | B2IV-V | 5.72 | 0.48 | 21.06 | $-0.53$ | $0.225^{+0.005}_{-0.005}$ | $-0.002$
HD 177989 | … | B0III | 9.34 | 0.23 | 17.81 | $-11.88$ | $2.41^{+0.20}_{-0.19}$ | $-0.496$
HD 178487 | … | B0Ib | 8.69 | 0.35 | 25.78 | $-8.56$ | $2.85^{+0.36}_{-0.25}$ | $-0.425$
HD 179407 | … | B0.5Ib | 9.44 | 0.28 | 24.02 | $-10.40$ | $4.44^{+0.69}_{-0.47}$ | $-0.802$
HD 185418 | … | B0.5V | 7.49 | 0.50 | 53.60 | $-2.17$ | $0.692^{+0.010}_{-0.009}$ | $-0.026$
HD 190918 | V1676 Cyg | O9.7Iab+WN4 | 6.75 | 0.45 | 72.65 | $+2.07$ | $1.78^{+0.07}_{-0.07}$ | $+0.064$
HD 191877 | … | B1Ib | 6.27 | 0.21 | 61.57 | $-6.45$ | $1.73^{+0.11}_{-0.13}$ | $-0.194$
HD 192035 | RX Cyg | B0III-IVn | 8.22 | 0.34 | 83.33 | $+7.76$ | $1.65^{+0.06}_{-0.06}$ | $+0.223$
HD 192639 | … | O7.5Iab | 7.11 | 0.66 | 74.90 | $+1.48$ | $1.81^{+0.07}_{-0.06}$ | $+0.047$
HD 195455 | … | B0.5III | 9.20 | 0.10 | 20.27 | $-32.14$ | $2.35^{+0.35}_{-0.24}$ | $-1.25$
HD 195965 | … | B0V | 6.97 | 0.25 | 85.71 | $+5.00$ | $0.790^{+0.023}_{-0.025}$ | $+0.069$
HD 198478 | 55 Cyg | B3Ia | 4.86 | 0.57 | 85.75 | $+1.49$ | $1.84^{+0.35}_{-0.22}$ | $+0.048$
HD 198781 | … | B0.5V | 6.45 | 0.35 | 99.94 | $+12.61$ | $0.915^{+0.022}_{-0.027}$ | $+0.200$
HD 201345 | … | ON9.2IV | 7.76 | 0.15 | 78.44 | $-9.54$ | $1.83^{+0.15}_{-0.11}$ | $-0.303$
HD 202347 | … | B1.5V | 7.50 | 0.17 | 88.22 | $-2.08$ | $0.764^{+0.023}_{-0.019}$ | $-0.028$
HD 203374 | … | B0IVpe | 6.67 | 0.53 | 100.51 | $+8.62$ | $2.04^{+1.42}_{-0.76}$ | $+0.306$
HD 203938 | … | B0.5IV | 7.08 | 0.74 | 90.56 | $-2.23$ | $2.96^{+2.26}_{-1.35}$ | $-0.115$
HD 206267 | … | O6Vf | 5.62 | 0.53 | 99.29 | $+3.74$ | $0.790^{+0.172}_{-0.112}$ | $+0.052$
HD 206773 | … | B0Vpe | 6.87 | 0.45 | 99.80 | $+3.62$ | $0.888^{+0.016}_{-0.014}$ | $+0.056$
HD 207198 | … | O8.5II | 5.94 | 0.62 | 103.14 | $+6.99$ | $0.978^{+0.034}_{-0.027}$ | $+0.119$
HD 207308 | … | B0.5V | 7.49 | 0.53 | 103.11 | $+6.82$ | $0.906^{+0.017}_{-0.013}$ | $+0.108$
HD 207538 | … | O9.7IV | 7.30 | 0.64 | 101.60 | $+4.67$ | $0.830^{+0.013}_{-0.013}$ | $+0.068$
HD 208440 | … | B1V | 7.91 | 0.28 | 104.03 | $+6.44$ | $0.877^{+0.019}_{-0.018}$ | $+0.098$
HD 209339 | … | O9.7IV | 8.51 | 0.36 | 104.58 | $+5.87$ | $0.936^{+0.028}_{-0.024}$ | $+0.096$
HD 210809 | … | O9Iab | 7.56 | 0.31 | 99.85 | $-3.13$ | $3.66^{+0.52}_{-0.34}$ | $-0.200$
HD 210839 |  $\lambda$ Cep | O6.5Infp | 5.05 | 0.57 | 103.83 | $+2.61$ | $0.833^{+0.066}_{-0.049}$ | $+0.038$
HD 212791 | V408 Lac | B3ne | 8.02 | 0.17 | 101.64 | $-4.30$ | $0.893^{+0.019}_{-0.016}$ | $-0.067$
HD 218915 | … | O9.2Iab | 7.20 | 0.30 | 108.06 | $-6.89$ | $2.97^{+0.33}_{-0.25}$ | $-0.357$
HD 219188 | … | B0.5IIIn | 7.06 | 0.13 | 83.03 | $-50.17$ | $2.10^{+0.25}_{-0.27}$ | $-1.61$
HD 220057 | … | B3IV | 6.94 | 0.23 | 112.13 | $+0.21$ | $0.385^{+0.004}_{-0.004}$ | $+0.001$
HD 224151 | V373 Cas | B0.5II-III | 6.00 | 0.44 | 115.44 | $-4.64$ | $1.89^{+0.13}_{-0.10}$ | $-0.153$
HDE 232522 | … | B1II | 8.70 | 0.27 | 130.70 | $-6.71$ | $3.46^{+0.41}_{-0.44}$ | $-0.404$
HDE 303308 | … | O4.5Vfc | 8.17 | 0.45 | 287.59 | $-0.61$ | $2.17^{+0.09}_{-0.10}$ | $-0.023$
HDE 308813 | … | O9.7IVn | 9.73 | 0.34 | 294.79 | $-1.61$ | $2.43^{+0.11}_{-0.09}$ | $-0.068$
BD$+$35 4258 | … | B0.5Vn | 9.46 | 0.25 | 77.19 | $-4.74$ | $2.21^{+0.14}_{-0.11}$ | $-0.182$
BD$+$53 2820 | … | B0IVn | 9.96 | 0.29 | 101.24 | $-1.69$ | $3.57^{+0.25}_{-0.17}$ | $-0.105$
CPD$-$59 2603 | V572 Car | O7Vnz | 8.81 | 0.46 | 287.59 | $-0.69$ | $2.63^{+0.16}_{-0.14}$ | $-0.032$
CPD$-$59 4552 | … | B1III | 8.24 | 0.38 | 303.22 | $+2.47$ | $2.04^{+0.07}_{-0.06}$ | $+0.088$
CPD$-$69 1743 | … | B0.5IIIn | 9.64 | 0.30 | 303.71 | $-7.35$ | $3.20^{+0.19}_{-0.16}$ | $-0.410$

###  2.1 Processing of the STIS Data

Archival STIS spectra acquired using either the medium-resolution echelle
grating (E140M) or the high-resolution grating (E140H) at central wavelength
settings that cover the relevant P ii and/or Cl i lines were obtained from the
Mikulski Archive for Space Telescopes (MAST). We focus only on sight lines
with reliable column densities of H i and H2 published in the literature
(e.g., Jenkins, 2019). An additional requirement for the Cl analysis is that
each sight line must have FUSE observations available.222An exception is made
in the case of HD 147933 ($\rho$ Oph A), which was not observed with FUSE but
does have high-resolution STIS spectra covering the Cl i $\lambda 1347$
transition. Moomey et al. (2012) obtained Cl i and Cl ii column densities for
this sight line using Copernicus observations. However, the total Cl abundance
derived from their analysis is unusually low. We therefore reexamine the
Copernicus data toward $\rho$ Oph A, in conjunction with the newer high-
resolution STIS spectra, in order to re-evaluate the Cl i and Cl ii column
densities in this direction. The final P and Cl samples differ from one
another slightly due to the constraints placed on measuring the various
absorption features (as discussed above). Basic information regarding the
background stars associated with the 107 sight lines that constitute the final
combined sample is presented in Table 1. The wavelengths and adopted
oscillator strengths for the interstellar lines of interest to our survey are
provided in Table 2.1.

After downloading the pipeline-processed archival files from MAST, the STIS
data were further reduced in a manner analogous to that described in Ritchey
et al. (2011, 2018). Multiple exposures of a given target acquired with the
same echelle grating were co-added to increase the signal-to-noise (S/N) ratio
in the final spectrum. When a feature of interest appeared in adjacent echelle
orders with sufficient continua on both sides of the line, the overlapping
portions of the two orders were averaged together. Portions of the co-added
spectrum surrounding interstellar lines of interest (typically 2 Å wide) were
cut from the data. These smaller spectral segments were then normalized by
fitting the continuum regions with a low-order polynomial function.

Table 2: Atomic Data

Species | $\lambda_{0}$ | $f$ | Ref.
---|---|---|---
| (Å) |  |
P ii | 1152.818 | 0.27200 | 1
| 1301.874 | 0.01960 | 2
| 1532.533 | 0.00737 | 3
Cl i | 1004.678 | 0.04730 | 4
| 1094.769 | 0.03850 | 4
| 1097.369 | 0.00880 | 5
| 1347.240 | 0.15300 | 5
| 1379.528 | 0.00269 | 6
Cl ii | 1071.036 | 0.01420 | 7

References. — (1) Federman et al. (2007), (2) Brown et al. (2018), (3) this
work, (4) Alkhayat et al. (2019), (5) Schectman et al. (1993), (6) Oliver &
Hibbert (2013), (7) Schectman et al. (2005).

In most cases, the process of normalizing the continuum was straightforward.
The P ii $\lambda 1532$ line and the Cl i $\lambda 1347$ and $\lambda 1379$
lines are relatively isolated and are easily distinguished from the underlying
stellar spectra. However, as already mentioned, the P ii $\lambda 1301$ line
is positioned very close to the strong O i $\lambda 1302$ resonance feature.
(The difference in wavelength between the two transitions amounts to 68 km
s-1.) Thus, great care had to be taken in choosing an appropriate
normalization for the P ii $\lambda 1301$ line. Any cases where the P ii line
was inextricably blended with O i were rejected. We considered the spectrum to
be “blended” if there was no obvious continuum region between the P ii feature
and the core of the O i line. Of the 107 sight lines in our combined sample,
only 48 were deemed to have unblended P ii $\lambda 1301$ absorption lines.
Furthermore, in many of the “unblended” cases, the P ii feature is
superimposed onto the gently sloping damping wing of the O i line. Thus, in
these cases, the continuum being fit is not the stellar continuum but the
Lorentzian damping wing of the O i resonance line.

###  2.2 Processing of the FUSE Data

All FUSE exposures for the sight lines in our survey were downloaded from the
MAST archive. For each detector segment, multiple exposures of a given target
were cross-correlated in wavelength space and then co-added by taking the
weighted mean of the measured intensities.333The registration and co-addition
of FUSE data was accomplished using IDL routines within the LTOOLS package,
available at the following URL:
https://archive.stsci.edu/fuse/analysis/idl_tools.html. The overlapping
portions of different detector segments that covered lines of interest to our
survey were then cross-correlated and co-added in the same manner. In this
way, we produced high S/N ratio spectra for the Cl ii $\lambda 1071$ line and
also, in some cases, for the Cl i $\lambda 1004$, $\lambda 1094$, and $\lambda
1097$ transitions.

A major complication when dealing with FUSE spectra are the numerous
absorption features arising from electronic transitions within the Lyman and
Werner bands of H2. Most concerning for this investigation is that the $R$(4)
line of the H2 (3$-$0) Lyman band at 1070.9 Å is very close to the Cl ii line
at 1071.0 Å. (The velocity separation between the two lines is only 38 km
s-1.) The difficulty in distinguishing the two features is compounded by the
relatively low resolution of the FUSE spectra ($\Delta v\approx 18$ km s-1).
The stellar continuum surrounding the H2 lines in the vicinity of the Cl ii
$\lambda 1071$ feature is easy to identify in most cases. The same continuum
fitting procedure used for the STIS data provided adequate solutions for the
normalization of the FUSE spectra. A more specialized procedure was required
to deblend the absorption features associated with the H2 and Cl ii lines near
1071 Å. This procedure is described in more detail in Section 3.2.2.

##  3 RESULTS ON COLUMN DENSITIES

In order to obtain column densities for P ii, Cl i, and Cl ii for the sight
lines in our survey, we employed the technique of multi-component Voigt
profile fitting using the code ISMOD (Sheffer et al., 2008). The profile
fitting routine treats the column density $N$, Doppler broadening parameter
$b$, and velocity $v$ of each component as a free parameter and seeks to
minimize the root mean square (rms) deviations between the observed spectrum
and the synthetic one. In the following subsections we describe in more detail
the specific procedures used to fit the P ii, Cl i, and Cl ii lines, each of
which presented us with a unique challenge.

Figure 1: Logarithmic difference in the column densities obtained from the P
ii $\lambda 1301$ and $\lambda 1532$ lines when the theoretical $f$-value from
Froese Fischer et al. (2006) is adopted for the $\lambda 1532$ transition and
the experimental $f$-value from Brown et al. (2018) is adopted for P ii
$\lambda 1301$. The weighted mean of the difference (0.022 dex, as indicated
by the solid horizontal line) yields an empirical $f$-value for the P ii
$\lambda 1532$ transition of 0.00737.

###  3.1 Phosphorus Column Densities

For our analysis of P ii column densities, we adopt the experimentally
determined oscillator strength for the P ii $\lambda 1301$ transition recently
published by Brown et al. (2018). Their $f$-value for P ii $\lambda 1301$
indicates that a downward revision of 0.188 dex should be applied to column
densities derived using the older Hibbert (1988) $f$-value (which is the one
recommended by Morton, 2003). The Brown et al. (2018) $f$-value for P ii
$\lambda 1301$ is consistent with more recent theoretical determinations
(Tayal, 2003; Froese Fischer et al., 2006), which indicate that the $f$-value
for P ii $\lambda 1532$ is also in need of substantial revision. The
theoretical calculations of Tayal (2003) and Froese Fischer et al. (2006)
suggest that a reduction in column density of $\sim$0.4 dex is needed for
results based on the Hibbert (1988) $f$-value for P ii $\lambda 1532$. Since
no recent experimental result is available for P ii $\lambda 1532$, we derived
an empirical $f$-value for this transition based on high-resolution STIS
spectra.

Table 3: Phosphorus Column Densities Star |  $W_{\lambda}$(1301) | log $N$(1301) |  $W_{\lambda}$(1532) | log $N$(1532) | log $N$(P ii)aaFinal P ii column density. In cases where both P ii lines are measured, the final column density is obtained from the weighted mean of the two results.
---|---|---|---|---|---
| (mÅ) |  | (mÅ) |  |
HD 108 | … | … | $32.6\pm 2.1$ | $14.49\pm 0.03$ | $14.49\pm 0.03$
HD 1383 | … | … | $40.0\pm 3.1$ | $14.56\pm 0.04$ | $14.56\pm 0.04$
HD 12323 | … | … | $28.4\pm 2.0$ | $14.45\pm 0.04$ | $14.45\pm 0.04$
HD 13268 | … | … | $39.9\pm 4.5$ | $14.52\pm 0.05$ | $14.52\pm 0.05$
HD 13745 | … | … | $44.1\pm 5.1$ | $14.53\pm 0.05$ | $14.53\pm 0.05$
HD 14434 | … | … | $52.2\pm 4.6$ | $14.75\pm 0.04$ | $14.75\pm 0.04$
HD 15137 | … | … | $31.7\pm 1.9$ | $14.44\pm 0.03$ | $14.44\pm 0.03$
HD 23478 | $18.8\pm 0.5$ | $13.95\pm 0.02$ | $11.9\pm 1.0$ | $13.96\pm 0.04$ | $13.95\pm 0.02$
HD 24190 | $28.1\pm 0.4$ | $14.15\pm 0.02$ | $17.2\pm 2.0$ | $14.13\pm 0.05$ | $14.15\pm 0.02$
HD 24534 | $19.6\pm 0.2$ | $14.05\pm 0.03$ | $13.3\pm 0.7$ | $14.05\pm 0.03$ | $14.05\pm 0.02$
HD 25443 | $64.5\pm 1.8$ | $14.70\pm 0.03$ | … | … | $14.70\pm 0.03$
HD 37903 | $16.9\pm 0.5$ | $13.88\pm 0.02$ | … | … | $13.88\pm 0.02$
HD 41161 | $49.9\pm 1.2$ | $14.35\pm 0.02$ | … | … | $14.35\pm 0.02$
HD 46223 | … | … | $44.1\pm 1.0$ | $14.61\pm 0.03$ | $14.61\pm 0.03$
HD 52266 | $50.3\pm 0.6$ | $14.46\pm 0.03$ | … | … | $14.46\pm 0.03$
HD 53975 | $41.8\pm 0.8$ | $14.27\pm 0.02$ | … | … | $14.27\pm 0.02$
HD 63005 | … | … | $27.3\pm 2.2$ | $14.47\pm 0.04$ | $14.47\pm 0.04$
HD 66788 | … | … | $31.6\pm 3.3$ | $14.39\pm 0.04$ | $14.39\pm 0.04$
HD 72754 | $34.8\pm 2.2$ | $14.33\pm 0.04$ | … | … | $14.33\pm 0.04$
HD 73882 | $33.9\pm 1.6$ | $14.22\pm 0.03$ | … | … | $14.22\pm 0.03$
HD 79186 | $51.6\pm 2.4$ | $14.46\pm 0.03$ | … | … | $14.46\pm 0.03$
HD 88115bbFirst line gives results from E140H data; second line gives results from E140M data. | $33.3\pm 1.8$ | $14.15\pm 0.03$ | … | … | $14.15\pm 0.03$
| $29.2\pm 2.8$ | $14.23\pm 0.05$ | $16.1\pm 4.2$ | $14.12\pm 0.10$ | $14.20\pm 0.05$
HD 89137 | $40.2\pm 1.8$ | $14.27\pm 0.03$ | … | … | $14.27\pm 0.03$
HD 90087 | $45.3\pm 0.8$ | $14.36\pm 0.02$ | $30.1\pm 1.7$ | $14.38\pm 0.03$ | $14.37\pm 0.02$
HD 91597 | … | … | $47.5\pm 4.1$ | $14.67\pm 0.04$ | $14.67\pm 0.04$
HD 91651 | … | … | $35.3\pm 3.3$ | $14.41\pm 0.04$ | $14.41\pm 0.04$
HD 91824 | $50.6\pm 0.9$ | $14.38\pm 0.02$ | … | … | $14.38\pm 0.02$
HD 91983 | $55.5\pm 2.1$ | $14.44\pm 0.02$ | … | … | $14.44\pm 0.02$
HD 92554 | … | … | $44.2\pm 6.7$ | $14.57\pm 0.06$ | $14.57\pm 0.06$
HD 93222 | … | … | $49.8\pm 1.3$ | $14.59\pm 0.01$ | $14.59\pm 0.01$
HD 93843 | … | … | $24.7\pm 2.3$ | $14.30\pm 0.04$ | $14.30\pm 0.04$
HD 94493 | $48.6\pm 1.8$ | $14.35\pm 0.02$ | … | … | $14.35\pm 0.02$
HD 97175 | $28.8\pm 2.2$ | $14.19\pm 0.04$ | … | … | $14.19\pm 0.04$
HD 99857 | … | … | $37.7\pm 2.4$ | $14.54\pm 0.04$ | $14.54\pm 0.04$
HD 99890 | … | … | $30.1\pm 3.9$ | $14.41\pm 0.05$ | $14.41\pm 0.05$
HD 100199 | … | … | $29.9\pm 2.8$ | $14.39\pm 0.04$ | $14.39\pm 0.04$
HD 101190bbFirst line gives results from E140H data; second line gives results from E140M data. | … | … | $25.0\pm 2.9$ | $14.31\pm 0.05$ | $14.31\pm 0.05$
| … | … | $24.7\pm 0.8$ | $14.31\pm 0.02$ | $14.31\pm 0.02$
HD 103779 | … | … | $29.3\pm 1.8$ | $14.36\pm 0.03$ | $14.36\pm 0.03$
HD 104705 | … | … | $28.6\pm 0.9$ | $14.34\pm 0.02$ | $14.34\pm 0.02$
HD 108639 | … | … | $46.5\pm 2.3$ | $14.57\pm 0.02$ | $14.57\pm 0.02$
HD 114886 | … | … | $38.0\pm 2.4$ | $14.48\pm 0.03$ | $14.48\pm 0.03$
HD 116852 | … | … | $12.2\pm 2.0$ | $14.06\pm 0.07$ | $14.06\pm 0.07$
HD 121968 | $12.4\pm 0.6$ | $13.71\pm 0.03$ | $7.1\pm 0.8$ | $13.71\pm 0.05$ | $13.71\pm 0.02$
HD 122879 | … | … | $45.6\pm 2.0$ | $14.63\pm 0.03$ | $14.63\pm 0.03$
HD 124314 | $73.2\pm 1.3$ | $14.60\pm 0.02$ | $51.9\pm 3.2$ | $14.64\pm 0.03$ | $14.61\pm 0.02$
HD 137595 | $29.0\pm 1.0$ | $14.13\pm 0.03$ | $15.2\pm 2.0$ | $14.05\pm 0.05$ | $14.11\pm 0.02$
HD 144965 | $18.1\pm 1.0$ | $13.95\pm 0.04$ | … | … | $13.95\pm 0.04$
HD 147683 | $31.6\pm 1.3$ | $14.21\pm 0.04$ | $19.7\pm 1.9$ | $14.19\pm 0.04$ | $14.21\pm 0.03$
HD 147888 | $24.7\pm 0.3$ | $14.26\pm 0.05$ | … | … | $14.26\pm 0.05$
HD 147933 | $25.6\pm 0.5$ | $14.32\pm 0.05$ | … | … | $14.32\pm 0.05$
HD 148422 | … | … | $48.9\pm 3.9$ | $14.68\pm 0.04$ | $14.68\pm 0.04$
HD 148937 | $81.2\pm 1.3$ | $14.65\pm 0.02$ | $54.5\pm 2.5$ | $14.65\pm 0.02$ | $14.65\pm 0.01$
HD 151805 | … | … | $58.8\pm 5.7$ | $14.70\pm 0.04$ | $14.70\pm 0.04$
HD 152590 | … | … | $28.4\pm 2.4$ | $14.45\pm 0.05$ | $14.45\pm 0.05$
HD 157857 | $43.6\pm 1.1$ | $14.42\pm 0.03$ | … | … | $14.42\pm 0.03$
HD 163522 | … | … | $31.9\pm 2.9$ | $14.40\pm 0.04$ | $14.40\pm 0.04$
HD 165955 | … | … | $20.4\pm 2.5$ | $14.21\pm 0.05$ | $14.21\pm 0.05$
HD 167402 | … | … | $30.8\pm 3.2$ | $14.38\pm 0.04$ | $14.38\pm 0.04$
HD 168941 | … | … | $19.5\pm 2.7$ | $14.18\pm 0.06$ | $14.18\pm 0.06$
HD 170740 | $28.4\pm 1.4$ | $14.36\pm 0.05$ | … | … | $14.36\pm 0.05$
HD 177989 | … | … | $17.0\pm 0.6$ | $14.10\pm 0.02$ | $14.10\pm 0.02$
HD 178487 | … | … | $42.2\pm 3.0$ | $14.55\pm 0.03$ | $14.55\pm 0.03$
HD 179407 | … | … | $37.6\pm 4.0$ | $14.54\pm 0.05$ | $14.54\pm 0.05$
HD 185418 | $43.3\pm 0.7$ | $14.43\pm 0.04$ | … | … | $14.43\pm 0.04$
HD 190918 | … | … | $50.5\pm 3.0$ | $14.60\pm 0.03$ | $14.60\pm 0.03$
HD 191877 | $41.6\pm 1.2$ | $14.30\pm 0.02$ | … | … | $14.30\pm 0.02$
HD 192035 | $50.9\pm 1.8$ | $14.57\pm 0.05$ | $33.5\pm 3.3$ | $14.50\pm 0.05$ | $14.53\pm 0.04$
HD 195455 | $14.4\pm 1.3$ | $13.73\pm 0.04$ | … | … | $13.73\pm 0.04$
HD 195965 | $39.9\pm 1.6$ | $14.28\pm 0.02$ | $24.6\pm 2.0$ | $14.27\pm 0.03$ | $14.28\pm 0.02$
HD 198478 | $57.1\pm 2.5$ | $14.78\pm 0.06$ | … | … | $14.78\pm 0.06$
HD 198781 | $38.4\pm 1.1$ | $14.24\pm 0.02$ | … | … | $14.24\pm 0.02$
HD 201345 | … | … | $18.6\pm 2.2$ | $14.14\pm 0.05$ | $14.14\pm 0.05$
HD 202347 | $21.2\pm 1.1$ | $13.96\pm 0.03$ | $13.6\pm 1.5$ | $14.00\pm 0.05$ | $13.97\pm 0.02$
HD 203374bbFirst line gives results from E140H data; second line gives results from E140M data. | $57.6\pm 1.1$ | $14.47\pm 0.02$ | … | … | $14.47\pm 0.02$
| $54.9\pm 1.5$ | $14.49\pm 0.03$ | $35.6\pm 2.7$ | $14.47\pm 0.04$ | $14.48\pm 0.02$
HD 203938 | $47.3\pm 2.2$ | $14.57\pm 0.05$ | … | … | $14.57\pm 0.05$
HD 206267 | $57.9\pm 1.4$ | $14.52\pm 0.03$ | … | … | $14.52\pm 0.03$
HD 206773 | $46.0\pm 1.0$ | $14.40\pm 0.03$ | … | … | $14.40\pm 0.03$
HD 207308 | $41.9\pm 1.3$ | $14.44\pm 0.05$ | $29.4\pm 1.3$ | $14.43\pm 0.03$ | $14.43\pm 0.03$
HD 207538 | $54.9\pm 0.9$ | $14.56\pm 0.04$ | $36.5\pm 1.4$ | $14.52\pm 0.03$ | $14.53\pm 0.02$
HD 209339 | $48.4\pm 1.1$ | $14.43\pm 0.03$ | $29.2\pm 1.6$ | $14.37\pm 0.03$ | $14.40\pm 0.02$
HD 210809 | … | … | $43.3\pm 4.7$ | $14.60\pm 0.05$ | $14.60\pm 0.05$
HD 210839 | $59.2\pm 1.0$ | $14.49\pm 0.03$ | … | … | $14.49\pm 0.03$
HD 218915 | … | … | $25.9\pm 1.1$ | $14.28\pm 0.02$ | $14.28\pm 0.02$
HD 219188 | $19.2\pm 0.4$ | $13.91\pm 0.02$ | … | … | $13.91\pm 0.02$
HD 220057 | $26.4\pm 1.2$ | $14.21\pm 0.05$ | … | … | $14.21\pm 0.05$
HD 224151 | $53.4\pm 1.8$ | $14.43\pm 0.02$ | $36.8\pm 3.2$ | $14.47\pm 0.04$ | $14.44\pm 0.02$
HDE 308813 | … | … | $27.8\pm 3.0$ | $14.35\pm 0.05$ | $14.35\pm 0.05$
BD$+$35 4258 | … | … | $29.7\pm 3.9$ | $14.35\pm 0.05$ | $14.35\pm 0.05$
BD$+$53 2820 | … | … | $37.5\pm 5.2$ | $14.58\pm 0.06$ | $14.58\pm 0.06$
CPD$-$59 2603bbFirst line gives results from E140H data; second line gives results from E140M data. | … | … | $53.4\pm 1.9$ | $14.64\pm 0.02$ | $14.64\pm 0.02$
| … | … | $51.7\pm 2.2$ | $14.78\pm 0.04$ | $14.78\pm 0.04$
CPD$-$59 4552 | $60.4\pm 2.5$ | $14.51\pm 0.03$ | … | … | $14.51\pm 0.03$
CPD$-$69 1743 | … | … | $35.5\pm 4.7$ | $14.50\pm 0.06$ | $14.50\pm 0.06$

There are thirteen sight lines in our sample that have high-resolution STIS
spectra covering both P ii transitions ($\lambda 1301$ and $\lambda 1532$) and
have a $\lambda 1301$ profile that is not significantly blended with O i
$\lambda 1302$. For these sight lines, we first fit the P ii $\lambda 1301$
profile using the $f$-value from Brown et al. (2018). We then adopted the
component structure found from the $\lambda 1301$ line to fit the $\lambda
1532$ profile using the $f$-value from Froese Fischer et al. (2006), which is
$f(1532)=0.00701$. The weighted mean of the (logarithmic) differences in the
column densities derived from the $\lambda 1301$ and $\lambda 1532$
transitions is $0.022\pm 0.012$ dex (see Figure 1). From this result, we
obtain an empirical $f$-value for the P ii $\lambda 1532$ transition of
0.00737, very similar to the theoretical results of Tayal (2003) and Froese
Fischer et al. (2006).

Our empirical $f$-value is adopted for the remaining sight lines in our sample
with observations of P ii $\lambda 1532$. (For the sight lines used in the
derivation of the empirical $f$-value, the column densities obtained from the
$\lambda 1532$ transition were adjusted downward by 0.022 dex.) The equivalent
widths and column densities derived from profile synthesis fits to the P ii
$\lambda 1301$ and $\lambda 1532$ transitions are presented in Table 3 for the
92 sight lines in our final P sample. Errors in equivalent width account for
uncertainties due to noise in the spectra and uncertainties in continuum
placement. The column density uncertainties include an additional term that
depends on the degree of saturation in the absorption profile. In cases where
both P ii transitions were analyzed, final P ii column densities were obtained
from a weighted mean of the individual results.

For four of the sight lines in our sample (HD 88115, HD 101190, HD 203374, and
CPD$-$59 2603), we were able to independently derive P ii column densities
from medium-resolution and high-resolution STIS spectra. Both results are
included in Table 3. In three out of the four cases, the two independent
results agree with one another within the uncertainties. However, for CPD$-$59
2603, the P ii column densities derived from fits to the E140H and E140M
spectra disagree at the $3.1\sigma$ level. Since the equivalent widths from
the two fits are consistent with one another, the difference in column density
is likely related to uncertainties regarding the degree of saturation in the
line profile. Indeed, the fit to the medium resolution spectrum results in a
smaller $b$-value for the dominant absorption component (1.3 versus 2.4 km
s-1) yielding a higher column density. The $b$-values are better constrained
in fits to the higher resolution spectra. Thus, for the depletion analysis, we
adopt the results from the E140H data for each of the four sight lines
discussed here.

Examples of our fits to the P ii $\lambda 1532$ and $\lambda 1301$ lines are
presented in Figures 2 and 3. In each case, we provide both the unnormalized
spectrum with the adopted continuum fit and the normalized spectrum with the
derived profile synthesis fit. For each of the examples shown, the component
structure obtained from the fit to the stronger P ii $\lambda 1301$ feature
was adopted in the fit to the $\lambda 1532$ line (as described above).

Figure 2: High-resolution STIS spectra in the vicinity of the P ii $\lambda
1532$ and $\lambda 1301$ lines toward HD 24534 (left panels) and HD 90087
(right panels). Upper panels show the unnormalized spectra with dashed lines
indicating the adopted continuum fits. Lower panels show the normalized
spectra with solid red lines representing fits to the P ii absorption
profiles.

Figure 3: Same as Figure 2 except toward HD 148937 (left panels) and HD 209339
(right panels).

In Table 3.1, we present a comparison between the P ii column densities that
we obtain and the column densities reported in the literature for sight lines
studied previously by Lebouteiller et al. (2005) and Cartledge et al. (2006).
All of the previous results on $N$(P ii) given in Table 3.1 have been adjusted
to be consistent with the set of $f$-values adopted in this investigation (see
Table 2.1). Also note that the P ii column densities from Lebouteiller et al.
(2005) shown in Table 3.1 are the ones that those authors obtain from STIS
observations of the P ii $\lambda 1532$ line. The Cartledge et al. (2006)
column densities were derived from STIS observations of the P ii $\lambda
1301$ line. Our values for $\log N({\rm P~{}\textsc{ii}})$ shown in Table 3.1
refer to the final P ii column densities from Table 3. In some cases, these
final values were obtained from the weighted mean of the results from the
$\lambda 1301$ and $\lambda 1532$ lines.

Table 4: Comparison with Previous Studies: P ii

Star | log $N$(P ii)
---|---
| This Work | Previous ResultaaAll previous results have been adjusted to be consistent with the set of $f$-values adopted in this investigation (see Table 2). | Ref.
HD 24534 | $14.05\pm 0.02$ | $14.03\pm 0.05$ | 1
HD 37903 | $13.88\pm 0.02$ | $13.91\pm 0.04$ | 2
HD 72754 | $14.33\pm 0.04$ | $14.27\pm 0.04$ | 2
HD 79186 | $14.46\pm 0.03$ | $14.55\pm 0.06$ | 2
HD 91824 | $14.38\pm 0.02$ | $14.38\pm 0.03$ | 2
HD 91983 | $14.44\pm 0.02$ | $14.44\pm 0.04$ | 2
HD 93222 | $14.59\pm 0.01$ | $14.43\pm 0.01$ | 1
HD 99857 | $14.54\pm 0.04$ | $14.25\pm 0.02$ | 1
HD 104705 | $14.34\pm 0.02$ | $14.15\pm 0.04$ | 1
HD 121968 | $13.71\pm 0.02$ | $13.63\pm 0.09$ | 1
HD 124314 | $14.61\pm 0.02$ | $14.57\pm 0.24$ | 1
HD 152590 | $14.45\pm 0.05$ | $14.46\pm 0.02$ | 2
HD 157857 | $14.42\pm 0.03$ | $14.43\pm 0.04$ | 2
HD 177989 | $14.10\pm 0.02$ | $14.08\pm 0.07$ | 1
HD 185418 | $14.43\pm 0.04$ | $14.45\pm 0.03$ | 2
HD 198478 | $14.78\pm 0.06$ | $14.50\pm 0.03$ | 2
HD 198781 | $14.24\pm 0.02$ | $14.24\pm 0.04$ | 2
HD 201345 | $14.14\pm 0.05$ | $14.23\pm 0.06$ | 2
HD 202347 | $13.97\pm 0.02$ | $13.94\pm 0.07$ | 1
HD 206773 | $14.40\pm 0.03$ | $14.42\pm 0.04$ | 2
HD 218915 | $14.28\pm 0.02$ | $14.14\pm 0.09$ | 1
HD 220057 | $14.21\pm 0.05$ | $14.15\pm 0.03$ | 2
HD 224151 | $14.44\pm 0.02$ | $14.48\pm 0.09$ | 1

References. — (1) Lebouteiller et al. (2005), (2) Cartledge et al. (2006).

Most of the previous results on $N$(P ii) for the sight lines shown in Table
3.1 are consistent with our values at approximately the $1\sigma$ level. For
HD 198478, however, the Cartledge et al. (2006) value is lower than ours by
$4.3\sigma$. Cartledge et al. (2006) report an equivalent width for the
$\lambda 1301$ line toward HD 198478 ($47.0\pm 1.4$) that is lower than our
equivalent width ($57.1\pm 2.5$). However, the difference in equivalent width
is considerably smaller than the difference in column density (0.28 dex). The
discrepancy seems to be caused, therefore, by a combination of continuum
placement uncertainties and differences in the optical depth correction (via
the $b$-value) for the dominant P ii absorption component.

Similarly, there are three cases where the P ii column density from
Lebouteiller et al. (2005) is significantly lower than our value. The P ii
column densities reported by Lebouteiller et al. (2005) for the sight lines to
HD 93222, HD 99857, and HD 104705 are lower than our values by $8.9\sigma$,
$6.9\sigma$, and $4.3\sigma$, respectively. While Lebouteiller et al. (2005)
do not provide equivalent widths with their column density determinations, the
continua surrounding the P ii $\lambda 1532$ lines in these directions are
fairly easy to discern. One exception is that, in our fit to the $\lambda
1532$ line toward HD 99857, we include very weak components at negative
velocities, which are seen in the $\lambda 1301$ line but are difficult to
discern at $\lambda 1532$. Lebouteiller et al. (2005) apparently assumed that
these features were part of the continuum. In general, we typically include
many more velocity components in our profile synthesis fits than Lebouteiller
et al. (2005) do in theirs. For HD 93222, HD 99857, and HD 104705, we include
6, 8, and 5 components, respectively, while Lebouteiller et al. (2005) include
3, 1, and 3 components. Generally speaking, more components will typically
result in smaller $b$-values and potentially larger column densities. However,
this does not seem to explain the significant discrepancies in the P ii column
densities noted for these three sight lines. If we simply integrate the
apparent optical depth profiles of the $\lambda 1532$ lines, a procedure which
should yield a lower limit to the true P ii column density (see, e.g., Savage
& Sembach, 1991), we find values of $\log N({\rm P~{}\textsc{ii}})=14.58\pm
0.02$ toward HD 93222, $14.47\pm 0.02$ toward HD 99857, and $14.34\pm 0.02$
toward HD 104705. These values are consistent with the results we obtain from
profile fitting but are considerably higher than the column densities reported
by Lebouteiller et al. (2005).

###  3.2 Chlorine Column Densities

####  3.2.1 Neutral Chlorine

Neutral chlorine column densities were obtained (in most cases) from
simultaneous profile synthesis fits to the Cl i $\lambda 1347$ line from STIS
spectra and one other weaker line. If STIS observations were available
covering the Cl i $\lambda 1379$ feature, then this line served as the weaker
line in the simultaneous fit. Along many sight lines, the strong Cl i $\lambda
1347$ feature shows weak absorption components displaced in velocity relative
to the main, optically-thick absorption component. In such cases, the $\lambda
1379$ line will typically only show absorption from the main component but the
line is optically thin. A simultaneous fit to both features therefore allows
us to probe the full velocity structure of the Cl i absorption and also
provides us with an accurate value for the total column density along the line
of sight.

Table 5: Chlorine Column Densities Star |  $W_{\lambda}$(1347) | LineaaIdentification and equivalent width of the weak Cl i line used to constrain a simultaneous profile synthesis fit with the $\lambda 1347$ line. |  $W_{\lambda}$aaIdentification and equivalent width of the weak Cl i line used to constrain a simultaneous profile synthesis fit with the $\lambda 1347$ line. | log $N$(Cl i) |  $W_{\lambda}$(1071) | log $N$(Cl ii) | log $N$(Cltot)
---|---|---|---|---|---|---|---
| (mÅ) |  | (mÅ) |  | (mÅ) |  |
HD 108 |  $116.5\pm 0.7$0 | 1379 |  $12.5\pm 1.6$0 | $14.47\pm 0.04$ | $20.1\pm 3.4$ | $14.28\pm 0.07$ | $14.69\pm 0.04$
HD 1383 |  $121.9\pm 0.9$0 | 1379 | $8.4\pm 2.4$ | $14.31\pm 0.05$ | … | … | …
HD 3827 |  0$8.8\pm 1.0$ |  | … | $12.57\pm 0.04$ |  0$7.0\pm 1.9$ | $13.72\pm 0.11$ | $13.75\pm 0.10$
HD 12323 | $79.1\pm 1.4$ | 1379 | $6.2\pm 1.4$ | $14.17\pm 0.05$ | $10.1\pm 2.9$ | $13.93\pm 0.11$ | $14.36\pm 0.05$
HD 13268 |  $132.9\pm 2.0$0 | 1379 |  $12.5\pm 2.5$0 | $14.48\pm 0.05$ | … | … | …
HD 13745 |  $122.8\pm 2.3$0 | 1379 | $8.1\pm 2.3$ | $14.28\pm 0.04$ | … | … | …
HD 14434 |  $124.3\pm 3.1$0 | 1379 |  $16.9\pm 3.5$0 | $14.64\pm 0.05$ | … | … | …
HD 15137 |  $107.5\pm 0.8$0 | 1097 |  $12.0\pm 6.4$0 | $14.14\pm 0.04$ | $13.9\pm 5.6$ | $14.06\pm 0.15$ | $14.40\pm 0.07$
HD 23478 | $35.0\pm 0.2$ | 1379 | $8.0\pm 0.4$ | $14.33\pm 0.02$ | $<9.5$ | $<13.82$ | …
HD 24190 | $54.8\pm 0.4$ | 1379 | $7.0\pm 0.9$ | $14.25\pm 0.04$ |  0$9.4\pm 1.7$ | $13.87\pm 0.07$ | $14.40\pm 0.04$
HD 24534 | $31.9\pm 0.1$ | 1379 |  $10.9\pm 0.1$0 | $14.51\pm 0.02$ |  0$4.6\pm 2.3$ | $13.55\pm 0.17$ | $14.55\pm 0.03$
HD 37903 | $25.5\pm 0.3$ | 1097 | $4.7\pm 3.1$ | $13.79\pm 0.08$ | $12.4\pm 3.9$ | $14.04\pm 0.12$ | $14.23\pm 0.09$
HD 41161 | $94.8\pm 0.7$ | 1004 |  $28.2\pm 2.9$0 | $13.92\pm 0.02$ | $15.1\pm 4.5$ | $14.06\pm 0.11$ | $14.29\pm 0.07$
HD 46223 | $89.5\pm 0.5$ | 1379 |  $13.3\pm 0.6$0 | $14.53\pm 0.02$ | $20.1\pm 4.8$ | $14.23\pm 0.09$ | $14.71\pm 0.04$
HD 52266 | $56.4\pm 0.6$ | 1004 |  $19.2\pm 9.1$0 | $13.79\pm 0.04$ |  0$18.8\pm 10.6$ | $14.20\pm 0.19$ | $14.35\pm 0.15$
HD 53975 | $40.7\pm 0.4$ |  | … | $13.46\pm 0.02$ | $12.4\pm 7.2$ | $13.97\pm 0.20$ | $14.09\pm 0.16$
HD 63005 | $65.7\pm 1.1$ | 1379 | $4.7\pm 1.8$ | $14.04\pm 0.06$ | $12.5\pm 3.6$ | $14.06\pm 0.11$ | $14.35\pm 0.07$
HD 66788 | $67.9\pm 1.4$ | 1379 | $2.2\pm 1.7$ | $13.70\pm 0.03$ | $17.8\pm 5.0$ | $14.15\pm 0.11$ | $14.28\pm 0.08$
HD 72754 | $39.8\pm 0.8$ |  | … | $13.74\pm 0.05$ | … | … | …
HD 73882 | $55.9\pm 0.7$ | 1379 |  $13.6\pm 1.2$0 | $14.62\pm 0.04$ | $20.1\pm 7.5$ | $14.25\pm 0.14$ | $14.77\pm 0.05$
HD 75309 | $53.3\pm 0.6$ | 1097 | $6.2\pm 3.8$ | $13.86\pm 0.05$ | $12.8\pm 3.1$ | $14.02\pm 0.10$ | $14.25\pm 0.06$
HD 88115 | $24.6\pm 0.7$ |  | … | $13.19\pm 0.04$ | $12.8\pm 1.2$ |  $14.09\pm 0.04$bbA small upward correction has been made to the Cl ii column density in this direction to account for blending between the Cl ii profile and the nearby H2 feature. | $14.14\pm 0.04$
HD 89137 | $49.9\pm 0.6$ | 1094 |  $19.9\pm 3.1$0 | $13.90\pm 0.05$ | $14.9\pm 3.3$ | $14.07\pm 0.09$ | $14.29\pm 0.06$
HD 90087 | $58.9\pm 0.6$ |  | … | $13.65\pm 0.04$ | $20.3\pm 1.6$ | $14.23\pm 0.03$ | $14.34\pm 0.03$
HD 91597 | $46.2\pm 2.2$ | 1379 |  $<6.4$ccA $3\sigma$ upper limit is provided because the equivalent width determined in the fit is less than the associated uncertainty. | $13.45\pm 0.03$ | $31.3\pm 3.0$ | $14.50\pm 0.04$ | $14.54\pm 0.04$
HD 91651 | $52.1\pm 1.4$ |  | … | $13.52\pm 0.02$ | … | … | …
HD 91824 | $67.7\pm 1.2$ |  | … | $13.80\pm 0.05$ | $15.4\pm 3.2$ | $14.07\pm 0.08$ | $14.26\pm 0.06$
HD 91983 | $64.4\pm 1.0$ |  | … | $13.79\pm 0.05$ | $17.1\pm 5.6$ | $14.13\pm 0.12$ | $14.29\pm 0.09$
HD 92554 | $51.8\pm 2.8$ | 1379 |  $<9.6$ccA $3\sigma$ upper limit is provided because the equivalent width determined in the fit is less than the associated uncertainty. | $13.51\pm 0.03$ | $46.4\pm 7.1$ | $14.70\pm 0.06$ | $14.72\pm 0.06$
HD 93129 | $96.9\pm 0.3$ | 1379 |  $10.0\pm 0.6$0 | $14.39\pm 0.02$ | … | … | …
HD 93205 | $66.0\pm 0.7$ | 1379 | $4.2\pm 1.3$ | $13.98\pm 0.06$ | … | … | …
HD 93222 | $48.6\pm 0.6$ | 1379 | $2.0\pm 0.8$ | $13.65\pm 0.04$ | … | … | …
HD 93843 | $56.5\pm 0.8$ |  | … | $13.55\pm 0.03$ | … | … | …
HD 94493 | $52.6\pm 0.8$ |  | … | $13.54\pm 0.03$ | $22.5\pm 2.8$ | $14.26\pm 0.05$ | $14.33\pm 0.04$
HD 97175 | $27.7\pm 0.9$ |  | … | $13.41\pm 0.04$ | $13.4\pm 4.6$ | $14.08\pm 0.13$ | $14.16\pm 0.11$
HD 99857 | $68.1\pm 0.5$ | 1097 | $8.2\pm 4.9$ | $13.99\pm 0.06$ | $16.7\pm 3.2$ |  $14.32\pm 0.08$bbA small upward correction has been made to the Cl ii column density in this direction to account for blending between the Cl ii profile and the nearby H2 feature. | $14.49\pm 0.06$
HD 99890 | $70.7\pm 0.2$ | 1379 | $2.3\pm 0.3$ | $13.70\pm 0.03$ | $20.6\pm 3.4$ | $14.21\pm 0.07$ | $14.33\pm 0.05$
HD 100199 | $70.7\pm 1.1$ | 1379 |  $13.0\pm 1.5$0 | $14.54\pm 0.04$ | … | … | …
HD 101190 | $66.2\pm 0.5$ | 1379 |  $10.6\pm 0.4$0 | $14.45\pm 0.02$ | … | … | …
HD 104705 | $53.4\pm 0.5$ | 1379 | $4.5\pm 0.7$ | $14.04\pm 0.05$ | … | … | …
HD 108639 | $75.5\pm 0.7$ | 1379 | $4.1\pm 1.0$ | $13.98\pm 0.05$ | … | … | …
HD 109399 | $53.7\pm 0.9$ |  | … | $13.56\pm 0.04$ | $15.4\pm 3.2$ | $14.11\pm 0.08$ | $14.22\pm 0.07$
HD 114886 |  $106.9\pm 0.9$0 | 1379 | $7.8\pm 1.3$ | $14.26\pm 0.05$ | … | … | …
HD 115455 | $91.5\pm 1.1$ |  | … | $13.89\pm 0.03$ | … | … | …
HD 116852 | $34.1\pm 0.3$ |  | … | $13.44\pm 0.04$ |  0$7.4\pm 1.1$ |  $13.97\pm 0.07$bbA small upward correction has been made to the Cl ii column density in this direction to account for blending between the Cl ii profile and the nearby H2 feature. | $14.08\pm 0.06$
HD 121968 |  0$6.9\pm 0.3$ | 1379 |  $<0.7$ccA $3\sigma$ upper limit is provided because the equivalent width determined in the fit is less than the associated uncertainty. | $12.57\pm 0.03$ | $10.7\pm 1.7$ | $13.96\pm 0.07$ | $13.98\pm 0.06$
HD 137595 | $59.3\pm 0.7$ | 1379 | $8.2\pm 0.9$ | $14.35\pm 0.04$ |  0$4.6\pm 2.1$ | $13.53\pm 0.16$ | $14.41\pm 0.04$
HD 144965 | $40.4\pm 0.7$ | 1094 |  $13.2\pm 3.1$0 | $13.63\pm 0.04$ | $<6.8$ | $<13.67$ | …
HD 147683 | $51.1\pm 0.5$ | 1379 |  $10.3\pm 0.7$0 | $14.46\pm 0.03$ | $10.4\pm 7.8$ | $13.92\pm 0.24$ | $14.57\pm 0.07$
HD 147888 | $25.8\pm 0.3$ | 1097 | $5.3\pm 3.4$ | $13.85\pm 0.08$ | $19.2\pm 4.0$ | $14.43\pm 0.09$ | $14.53\pm 0.08$
HD 147933ddData on Cl i $\lambda 1097$ and Cl ii $\lambda 1071$ toward HD 147933 ($\rho$ Oph A) are obtained from Copernicus observations. | $26.8\pm 0.2$ | 1097 | $4.1\pm 3.7$ | $13.71\pm 0.08$ | $15.6\pm 3.8$ | $14.31\pm 0.10$ | $14.41\pm 0.09$
HD 148422 | $84.9\pm 1.6$ | 1379 | $7.2\pm 1.9$ | $14.23\pm 0.05$ | … | … | …
HD 148937 |  $120.8\pm 0.6$0 | 1379 |  $13.2\pm 1.3$0 | $14.51\pm 0.03$ | … | … | …
HD 151805 | $81.0\pm 1.9$ | 1379 | $6.6\pm 2.4$ | $14.20\pm 0.05$ | … | … | …
HD 152590 | $71.5\pm 0.3$ | 1379 | $6.7\pm 1.7$ | $14.22\pm 0.05$ | $18.3\pm 4.3$ | $14.19\pm 0.09$ | $14.51\pm 0.05$
HD 156359 |  0$7.4\pm 1.1$ | 1379 |  $<4.1$ccA $3\sigma$ upper limit is provided because the equivalent width determined in the fit is less than the associated uncertainty. | $12.59\pm 0.06$ |  0$9.7\pm 3.0$ | $13.93\pm 0.12$ | $13.95\pm 0.11$
HD 157857 | $76.6\pm 0.6$ | 1097 | $9.4\pm 5.3$ | $14.04\pm 0.05$ |  0$7.0\pm 4.0$ | $13.72\pm 0.20$ | $14.21\pm 0.08$
HD 163522 | $52.2\pm 1.9$ | 1379 |  $<6.2$ccA $3\sigma$ upper limit is provided because the equivalent width determined in the fit is less than the associated uncertainty. | $13.51\pm 0.03$ | … | … | …
HD 165246 | $39.2\pm 0.3$ | 1094 |  $18.4\pm 4.1$0 | $13.99\pm 0.06$ | $20.7\pm 1.6$ |  $14.32\pm 0.04$bbA small upward correction has been made to the Cl ii column density in this direction to account for blending between the Cl ii profile and the nearby H2 feature. | $14.49\pm 0.03$
HD 165955 | $30.2\pm 1.5$ | 1379 |  $<4.8$ccA $3\sigma$ upper limit is provided because the equivalent width determined in the fit is less than the associated uncertainty. | $13.23\pm 0.03$ | $15.5\pm 2.7$ | $14.13\pm 0.07$ | $14.18\pm 0.06$
HD 167402 | $35.4\pm 1.0$ | 1379 | $5.3\pm 1.3$ | $14.13\pm 0.07$ | $12.7\pm 4.5$ | $13.98\pm 0.13$ | $14.37\pm 0.07$
HD 168076 |  $128.7\pm 1.8$0 | 1379 |  $18.0\pm 2.5$0 | $14.72\pm 0.04$ |  0$41.8\pm 17.1$ | $14.60\pm 0.15$ | $14.97\pm 0.07$
HD 168941 | $44.5\pm 1.0$ | 1379 | $7.1\pm 1.1$ | $14.27\pm 0.05$ | $11.3\pm 2.9$ | $13.95\pm 0.10$ | $14.44\pm 0.05$
HD 177989 | $56.9\pm 0.6$ | 1379 | $4.8\pm 0.5$ | $14.09\pm 0.04$ |  0$3.7\pm 2.5$ | $13.42\pm 0.22$ | $14.18\pm 0.06$
HD 178487 | $71.9\pm 1.5$ | 1379 | $8.7\pm 1.3$ | $14.34\pm 0.05$ | $<16.7$ | $<14.07$ | …
HD 179407 | $72.6\pm 1.8$ | 1379 | $8.0\pm 2.6$ | $14.30\pm 0.06$ | $<21.5$ | $<14.17$ | …
HD 185418 | $67.5\pm 0.4$ | 1097 |  $16.9\pm 2.4$0 | $14.38\pm 0.04$ | $<14.7$ | $<14.01$ | …
HD 190918 |  $115.2\pm 1.3$0 | 1379 | $7.6\pm 1.4$ | $14.25\pm 0.04$ | $29.6\pm 3.5$ | $14.38\pm 0.05$ | $14.62\pm 0.03$
HD 191877 | $69.1\pm 0.5$ | 1097 | $6.2\pm 0.8$ | $13.85\pm 0.03$ |  0$8.1\pm 1.1$ | $13.78\pm 0.06$ | $14.12\pm 0.03$
HD 192035 | $62.0\pm 2.3$ | 1379 | $8.4\pm 1.5$ | $14.40\pm 0.06$ | $10.6\pm 4.0$ | $13.93\pm 0.14$ | $14.52\pm 0.06$
HD 192639 | $67.6\pm 0.4$ | 1097 |  $14.3\pm 1.7$0 | $14.27\pm 0.04$ | $18.0\pm 3.6$ | $14.22\pm 0.08$ | $14.54\pm 0.04$
HD 195455 |  0$9.6\pm 0.8$ |  | … | $12.62\pm 0.03$ |  0$4.3\pm 3.7$ | $13.49\pm 0.27$ | $13.55\pm 0.24$
HD 195965 | $67.8\pm 0.9$ | 1379 | $5.1\pm 1.4$ | $14.08\pm 0.06$ |  0$5.4\pm 0.8$ | $13.59\pm 0.06$ | $14.20\pm 0.05$
HD 201345 | $53.3\pm 0.6$ |  | … | $13.54\pm 0.03$ | $16.2\pm 1.7$ | $14.12\pm 0.04$ | $14.22\pm 0.03$
HD 202347 | $62.5\pm 0.7$ | 1379 | $3.3\pm 2.2$ | $13.88\pm 0.04$ | … | … | …
HD 203374 | $89.1\pm 0.6$ | 1379 |  $10.4\pm 2.0$0 | $14.44\pm 0.05$ | $23.8\pm 2.2$ | $14.29\pm 0.04$ | $14.67\pm 0.04$
HD 203938 | $68.6\pm 0.6$ | 1379 |  $20.1\pm 1.3$0 | $14.80\pm 0.03$ | $<38.3$ | $<14.42$ | …
HD 206267 | $76.2\pm 0.6$ | 1097 | $9.4\pm 1.2$ | $14.04\pm 0.04$ | $12.9\pm 3.2$ | $14.00\pm 0.10$ | $14.32\pm 0.05$
HD 206773 | $73.1\pm 0.5$ | 1097 |  $12.4\pm 1.5$0 | $14.19\pm 0.04$ | $12.0\pm 2.3$ | $13.97\pm 0.08$ | $14.40\pm 0.04$
HD 207198 | $77.1\pm 0.4$ | 1097 |  $12.3\pm 3.7$0 | $14.17\pm 0.04$ | $13.8\pm 7.3$ | $14.04\pm 0.18$ | $14.41\pm 0.09$
HD 207308 | $67.5\pm 0.7$ | 1379 |  $15.6\pm 1.1$0 | $14.64\pm 0.03$ | $18.3\pm 7.1$ | $14.23\pm 0.14$ | $14.78\pm 0.05$
HD 207538 | $74.9\pm 0.9$ | 1379 |  $14.4\pm 0.6$0 | $14.67\pm 0.03$ | $13.9\pm 5.6$ | $14.05\pm 0.15$ | $14.77\pm 0.04$
HD 208440 | $90.5\pm 0.7$ | 1097 | $9.5\pm 1.8$ | $14.04\pm 0.03$ | $17.1\pm 2.8$ | $14.14\pm 0.07$ | $14.39\pm 0.04$
HD 209339 | $73.5\pm 0.5$ | 1097 | $7.7\pm 1.3$ | $13.94\pm 0.03$ | $13.5\pm 1.6$ | $14.03\pm 0.05$ | $14.29\pm 0.03$
HD 210809 |  $109.6\pm 0.7$0 | 1379 | $4.1\pm 2.4$ | $13.97\pm 0.02$ | … | … | …
HD 210839 |  $106.2\pm 0.7$0 | 1097 |  $12.4\pm 2.3$0 | $14.16\pm 0.04$ | $11.8\pm 3.4$ |  $14.01\pm 0.11$bbA small upward correction has been made to the Cl ii column density in this direction to account for blending between the Cl ii profile and the nearby H2 feature. | $14.39\pm 0.05$
HD 212791 | $62.2\pm 1.0$ |  | … | $13.65\pm 0.03$ | $11.4\pm 5.5$ | $13.93\pm 0.17$ | $14.11\pm 0.12$
HD 218915 | $77.3\pm 0.4$ | 1379 | $6.2\pm 1.0$ | $14.16\pm 0.04$ | … | … | …
HD 219188 | $40.0\pm 0.3$ |  | … | $13.45\pm 0.03$ | $11.1\pm 2.2$ | $13.95\pm 0.08$ | $14.07\pm 0.06$
HD 220057 | $40.1\pm 0.7$ | 1097 | $6.4\pm 1.6$ | $13.89\pm 0.05$ |  0$6.1\pm 2.6$ | $13.68\pm 0.16$ | $14.10\pm 0.07$
HDE 232522 | $82.2\pm 0.2$ | 1379 | $5.8\pm 0.5$ | $14.13\pm 0.03$ | … | … | …
HDE 303308 | $78.8\pm 0.7$ | 1379 | $8.5\pm 1.9$ | $14.32\pm 0.05$ | … | … | …
HDE 308813 | $76.0\pm 2.1$ | 1379 | $7.8\pm 2.0$ | $14.30\pm 0.05$ | … | … | …
BD$+$35 4258 | $66.5\pm 2.0$ | 1379 | $4.4\pm 2.0$ | $14.02\pm 0.06$ | $26.3\pm 4.1$ | $14.35\pm 0.06$ | $14.52\pm 0.05$
BD$+$53 2820 |  $104.1\pm 2.7$0 | 1379 | $3.6\pm 3.3$ | $13.91\pm 0.04$ | … | … | …
CPD$-$59 2603 | $58.9\pm 0.7$ | 1379 | $8.5\pm 1.3$ | $14.33\pm 0.05$ | … | … | …
CPD$-$59 4552 | $82.9\pm 1.4$ |  | … | $13.95\pm 0.04$ | … | … | …
CPD$-$69 1743 | $40.1\pm 2.1$ | 1379 |  $<7.1$ccA $3\sigma$ upper limit is provided because the equivalent width determined in the fit is less than the associated uncertainty. | $13.53\pm 0.04$ | … | … | …

Our Cl i fits employed the same profile fitting routine as was used to fit the
P ii lines, except that the two Cl i absorption profiles were fitted
simultaneously. This means that the column densities, $b$-values, and
velocities of the individual components included in the fits to the two
profiles are necessarily identical, but the best fitting values of the
parameters are still determined iteratively through an rms-minimizing
approach. The equivalent widths and total Cl i column densities derived
through our profile synthesis fits are provided in Table 5 for a total of 98
sight lines. Examples of our simultaneous fits to the Cl i $\lambda 1379$ and
$\lambda 1347$ lines are presented in Figures 4 and 5. In a few cases, the
$\lambda 1379$ line, although included in the simultaneous fit, is not
significantly detected (i.e., the derived equivalent width is smaller than the
associated uncertainty). An example of this is provided by the line of sight
to HD 121968 (Figure 5). In these situations, we list $3\sigma$ upper limits
to the equivalent width of the $\lambda 1379$ line in Table 5.

There are 19 sight lines included in Table 5 for which the Cl i $\lambda 1379$
line is not available in the archival STIS data, but the $\lambda 1347$
feature is deemed to be weak enough that optical depth effects are not a major
concern. In each of these cases, the relative intensity of the $\lambda 1347$
line does not drop below $\sim$0.05 and the strongest absorption component
retains a Gaussian shape. In other words, the shape of the profile does not
appear to be distorted by saturation in the line core. For these 19 sight
lines, a profile synthesis fit to the $\lambda 1347$ line alone yielded the
total Cl i column density.

Figure 4: High-resolution STIS spectra in the vicinity of the Cl i $\lambda
1379$ and $\lambda 1347$ lines toward HD 24534 (left panels) and HD 104705
(right panels). Solid red lines represent simultaneous profile synthesis fits
to the two Cl i lines in each direction.

Figure 5: Same as Figure 4 except toward HD 121968 (left panels) and HD 232522
(right panels).

Figure 6: Simultaneous profile synthesis fits to the Cl i $\lambda 1347$ and
$\lambda 1097$ lines toward HD 206267 (left panels) and HD 210839 (right
panels). The $\lambda 1347$ data are from high-resolution STIS spectra, while
the $\lambda 1097$ spectra are from FUSE observations.

There are also a significant number of sight lines where there are no
observations covering the Cl i $\lambda 1379$ line and where the $\lambda
1347$ line is much too strong to fit on its own. For many of these sight
lines, we use observations of the Cl i $\lambda 1004$, $\lambda 1094$, or
$\lambda 1097$ transition, available from FUSE spectra, to help constrain the
Cl i column density. Experimentally determined $f$-values for the Cl i
transitions at 1004.7 Å and 1094.8 Å were recently reported by Alkhayat et al.
(2019). These $f$-values, from beam-foil experiments, are in good agreement
with the empirical $f$-values derived for these transitions by Sonnentrucker
et al. (2006). A secure experimental $f$-value is also available for the Cl i
$\lambda 1097$ transition (Schectman et al., 1993).

The Cl i $\lambda 1097$ transition is the weakest of the three FUSE
transitions mentioned above. We therefore prefer to use this transition to
constrain the Cl i column density provided that the line can be reliably
measured. In cases where the $\lambda 1097$ line is too weak to detect or is
affected by noise in the spectrum, we chose one of the other two lines
instead, whichever appeared to provide the most reliable measurement. We then
performed a simultaneous profile synthesis fit using the $\lambda 1347$ line
from STIS data and whichever FUSE transition was chosen to constrain the total
column density. Since the FUSE wavelength scale is known to be rather poorly
calibrated, we first corrected the velocity scale of the FUSE spectrum so that
the centroid velocity of the Cl i absorption feature was consistent with the
(weighted mean) velocity of the $\lambda 1347$ line before proceeding with the
simultaneous fit. The velocity shift applied to the FUSE spectrum was
calculated either directly from the Cl i $\lambda 1347$ line or from the S i
$\lambda 1295$ transition if the $\lambda 1347$ line was too heavily
saturated. The S i $\lambda 1295$ transition (from high-resolution STIS
spectra) also provided initial values for the component parameters (i.e., the
relative velocities, $b$-values, and component fractions) in cases where the
Cl i $\lambda 1347$ line was extremely optically thick. Examples of our
simultaneous profile synthesis fits to the Cl i $\lambda 1347$ line from STIS
spectra and the $\lambda 1097$ feature from FUSE data are provided in Figure
6. The resulting total Cl i column densities are included in Table 5.

####  3.2.2 Singly Ionized Chlorine

Analysis of the Cl ii $\lambda 1071$ feature, obtained from FUSE observations,
is complicated by the presence of a nearby H2 absorption line (as described in
Section 2.2). The velocity separation between the Cl ii $\lambda 1071$ line
and the H2 (3$-$0) $R$(4) line is only 38 km s-1, with the H2 feature
positioned to the blue of the Cl ii line. Thus, if the Cl ii absorption
profile includes components that are blueshifted by $\sim$40 km s-1 (relative
to the main absorption component), then the Cl ii profile will be severely
blended with H2 absorption from the main component. Unfortunately, many of the
sight lines in our Cl i sample show components that are blueshifted with
respect to the main absorption complex (both in the Cl i $\lambda 1347$ line
and in lines from dominant ions, such as P ii $\lambda 1301$). Typically,
sight lines that exhibit multiple absorption complexes at different velocities
are those that probe material in a distant spiral arm (such as the
Sagittarius-Carina spiral arm or the Perseus spiral arm). In these cases, the
absorption complexes at relatively low velocity correspond to material in the
“local arm”, whereas the absorption complexes displaced to negative velocities
correspond to material in the more distant arm. (One example of this is
provided by the Cl i $\lambda 1347$ line toward HD 104705 shown in Figure 4.
The absorption near $v_{\sun}\approx-25$ km s-1 is associated with diffuse
molecular gas in the Sagittarius-Carina spiral arm.)

We do not attempt to derive Cl ii column densities along any sight lines where
it is likely that the Cl ii absorption profile is severely blended with H2
absorption. In order to determine whether a given sight line exhibits blended
absorption, we examined both the Cl i $\lambda 1347$ profile and the P ii
$\lambda 1301$ and/or $\lambda 1532$ profile. (The H2 absorption probably
closely follows the Cl i absorption, while the Cl ii profile should be more
similar to that of P ii.) We looked for sight lines where most of the
absorption was positioned at velocities less than $\sim$20 km s-1 relative to
the deepest part of the absorption profile. Sight lines with prominent
absorption complexes located more than $\sim$20 km s-1 from the main component
were rejected.

For sight lines that passed this initial inspection, we still needed to
“deblend” the Cl ii and H2 absorption lines. Due to the low resolution of the
FUSE spectrograph ($\sim$18 km s-1), the wings of the Cl ii and H2 lines will
be blended even if the intrinsic spread in absorption is rather narrow. To
deblend the two features, we started with a simple Voigt profile fit
consisting of two components, one component for the H2 line and one component
for the Cl ii feature. We then used the parameters derived for the H2
component to create a synthetic spectrum that was divided into the observed
spectrum, thereby removing the absorption associated with H2. Any remaining
absorption was then attributed to Cl ii $\lambda 1071$.

Figure 7: FUSE spectra in the vicinity of the Cl ii $\lambda 1071$ line toward
HD 24190 (left panel) and HD 88115 (right panel). Solid red lines represent
profile synthesis fits to the Cl ii feature. The smooth black curves show the
Voigt profile fits used to deblend and remove the H2 lines.

Figure 8: Same as Figure 7 except toward HD 137595 (left panel) and HD 156359
(right panel). Note that the H2 (3$-$0) $R$(4) line is not detected toward HD
156359.

Figure 9: Same as Figure 7 except toward HD 165246 (left panel) and HD 191877
(right panel).

We then proceeded to fit the Cl ii $\lambda 1071$ feature with our usual
profile fitting routine. However, because the low resolution of the FUSE
spectra makes it impossible to discern individual components, we adopted a
profile template for the component structure of the Cl ii $\lambda 1071$ line
based on either the P ii $\lambda 1532$ line or the Ge ii $\lambda 1237$
feature. Both of these lines (which are available from STIS observations) are
expected to have an equivalent width that is comparable to (but somewhat
larger than) that of Cl ii $\lambda 1071$. Given the difference in the cosmic
abundances of P and Cl, and taking into account the different wavelengths and
$f$-values of the two transitions, the expected equivalent width ratio between
the $\lambda 1071$ and $\lambda 1532$ lines is $W_{1071}/W_{1532}\approx 0.6$
(assuming optically thin absorption and no differences in depletion).
Similarly, the expected ratio between the Cl ii $\lambda 1071$ and Ge ii
$\lambda 1237$ lines is $W_{1071}/W_{1237}\approx 0.5$. By using P ii $\lambda
1532$ or Ge ii $\lambda 1237$ as a template, we can therefore be reasonably
well assured that our fits to the Cl ii $\lambda 1071$ feature include all of
the absorption components that are likely to be detectable.

For many of our sight lines, profile fitting results were already available
for P ii $\lambda 1532$. Thus, these served as the basis for fitting the Cl ii
$\lambda 1071$ feature. In cases where observations of the P ii $\lambda 1532$
line were not available, the Ge ii $\lambda 1237$ line was used instead.
Component structures for Ge ii were obtained either from new analyses of the
Ge ii $\lambda 1237$ line or from the component results published in Ritchey
et al. (2018). The relative velocities, $b$-values, and component fractions
(obtained from P ii or Ge ii) were held fixed in the fit to the Cl ii $\lambda
1071$ line. The only free parameters were the total Cl ii column density and
an overall velocity offset. The equivalent widths and total Cl ii column
densities derived from these fits are provided in Table 5. Examples of our
fits to the H2 and Cl ii features near 1071 Å are presented in Figures 7–9. In
these figures, the smooth black curve represents the Voigt profile fit used to
deblend and remove the H2 line, while the red curve indicates the final
profile synthesis fit to the Cl ii feature. Total chlorine column densities,
where $N({\rm Cl}_{\rm tot})=N({\rm Cl~{}\textsc{i}})+N({\rm
Cl~{}\textsc{ii}})$, are also provided in Table 5.

There are five sight lines where a relatively weak negative velocity
absorption component (or group of components) appears in the P ii or Ge ii
absorption profile (at a velocity more than $\sim$20 km s-1 from the main
component), but the column density associated with the blueshifted absorption
is no more than 20% of the total line-of-sight column density. These cases
were initially rejected for exhibiting blended H2 and Cl ii absorption.
Ultimately, however, we decided to fit the unblended portion of the Cl ii
absorption and apply a correction to the resulting column density to account
for the “missing” absorption. The magnitude of the correction factor was
determined from the fractional column density of the P ii or Ge ii components
not included in the profile fit. These correction factors amounted to upward
revisions in the total Cl ii column density of 0.08 dex for HD 88115, 0.10 dex
for HD 99857, 0.06 dex for HD 116852, 0.02 dex for HD 165246, and 0.05 dex for
HD 210839. In most of these cases, the applied correction is comparable to the
uncertainty in the Cl ii column density obtained from the profile fit.

For six sight lines (HD 23478, HD 144965, HD 178487, HD 179407, HD 185418, and
HD 203938), we are unable to derive a Cl ii column density, not because the Cl
ii $\lambda 1071$ feature is blended with H2 absorption, but because the Cl ii
line is not significantly detected. These may be cases where nearly all of the
chlorine is in neutral form. We therefore calculated 3$\sigma$ upper limits to
the equivalent width and column density of the Cl ii $\lambda 1071$ feature
based on an assumed component structure (from the observed P ii line). These
limits are provided in Table 5. Unfortunately, none of the derived Cl ii upper
limits are low enough that the Cl ii column density can be neglected in
deriving a total chlorine abundance. We therefore do not provide total Cl
column densities for these sight lines.

###  3.3 Atomic and Molecular Hydrogen Column Densities

Column densities of atomic and molecular hydrogen for the sight lines in our
combined phosphorus and chlorine sample were obtained either from Jenkins
(2019) or from the values tabulated in Jenkins (2009). These values are
provided in Table 3.3 along with the total hydrogen column densities, $N({\rm
H}_{\rm tot})=N({\rm H~{}\textsc{i}})+2N({\rm H}_{2})$, and molecular hydrogen
fractions, $f({\rm H}_{2})=2N({\rm H}_{2})/N({\rm H}_{\rm tot})$. For the line
of sight to HD 165955, Jenkins (2009) reports a molecular hydrogen column
density of $\log N({\rm H}_{2})=16.53\pm 0.04$. This value was originally
derived by Cartledge et al. (2004). However, based on the Cl i column density
we obtain for this sight line, $\log N({\rm Cl~{}\textsc{i}})=13.23\pm 0.03$,
and the correlations discussed in the next section, the reported value of
$N({\rm H}_{2})$ toward HD 165955 appears to be much too low. We therefore
undertook our own analysis of the FUSE data in this direction to determine a
more accurate value for the molecular hydrogen column density. From an
analysis of the $J=0$ and 1 lines of the H2 (1$-$0), (2$-$0), (3$-$0), and
(4$-$0) bands444For the analysis of H2 absorption toward HD 165955, we used
the IDL-based program H2GUI, which was originally created and used by
Tumlinson et al. (2002)., we find $\log N({\rm H}_{2})=19.09\pm 0.07$. A
representative fit to the (4$-$0) band is presented in Figure 10.

Table 6: Atomic and Molecular Hydrogen Column Densities

Star | log $N$(H i) | log $N$(H2) | log $N$(Htot) | log $f$(H2) |  $F_{*}$aaSight line depletion factor. | Ref.
---|---|---|---|---|---|---
HD 108 | $21.38\pm 0.05$ | $20.45\pm 0.08$ | $21.47^{+0.04}_{-0.05}$ | $-0.72\pm 0.09$ | $0.53\pm 0.04$ | 1
HD 1383 | $21.46\pm 0.05$ | $20.49\pm 0.07$ | $21.54^{+0.04}_{-0.05}$ | $-0.75\pm 0.08$ | $0.48\pm 0.05$ | 1
HD 3827 | $20.55\pm 0.07$ | $18.35\pm 0.25$ | $20.56^{+0.07}_{-0.08}$ | $-1.90\pm 0.25$ | $0.41\pm 0.07$ | 1
HD 12323 | $21.19\pm 0.04$ | $20.26\pm 0.08$ | $21.28^{+0.04}_{-0.04}$ | $-0.72\pm 0.09$ | $0.51\pm 0.05$ | 1
HD 13268 | $21.34\pm 0.07$ | $20.46\pm 0.08$ | $21.44^{+0.06}_{-0.07}$ | $-0.68\pm 0.10$ | $0.45\pm 0.07$ | 1
HD 13745 | $21.34\pm 0.05$ | $20.53\pm 0.07$ | $21.46^{+0.04}_{-0.05}$ | $-0.63\pm 0.08$ | $0.55\pm 0.05$ | 1
HD 14434 | $21.37\pm 0.09$ | $20.47\pm 0.07$ | $21.47^{+0.07}_{-0.09}$ | $-0.70\pm 0.10$ | $0.52\pm 0.04$ | 2
HD 15137 | $21.24\pm 0.08$ | $20.23\pm 0.08$ | $21.32^{+0.07}_{-0.08}$ | $-0.79\pm 0.10$ | $0.36\pm 0.06$ | 1
HD 23478 | $20.71\pm 0.17$ | $20.48\pm 0.07$ | $21.05^{+0.09}_{-0.12}$ | $-0.27\pm 0.11$ | $0.00\pm 0.50$ | 2
HD 24190 | $21.18\pm 0.06$ | $20.38\pm 0.07$ | $21.30^{+0.05}_{-0.06}$ | $-0.62\pm 0.08$ | $0.63\pm 0.24$ | 2
HD 24534 | $20.73\pm 0.06$ | $20.92\pm 0.04$ | $21.34^{+0.03}_{-0.04}$ | $-0.12\pm 0.05$ | $0.90\pm 0.06$ | 2
HD 25443 | $21.29\pm 0.06$ | $20.92\pm 0.08$ | $21.56^{+0.05}_{-0.06}$ | $-0.34\pm 0.09$ | $0.72\pm 0.05$ | 1
HD 37903 | $21.12\pm 0.10$ | $20.85\pm 0.07$ | $21.44^{+0.06}_{-0.07}$ | $-0.29\pm 0.09$ | $1.15\pm 0.03$ | 2
HD 41161 | $21.09\pm 0.05$ | $19.99\pm 0.09$ | $21.15^{+0.05}_{-0.05}$ | $-0.86\pm 0.10$ | $0.35\pm 0.04$ | 1
HD 46223 | $21.46\pm 0.04$ | $20.67\pm 0.06$ | $21.58^{+0.03}_{-0.04}$ | $-0.61\pm 0.07$ | $0.86\pm 0.03$ | 1
HD 52266 | $21.22\pm 0.04$ | $20.00\pm 0.08$ | $21.27^{+0.04}_{-0.04}$ | $-0.97\pm 0.09$ | $0.56\pm 0.05$ | 1
HD 53975 | $21.08\pm 0.04$ | $19.15\pm 0.09$ | $21.09^{+0.04}_{-0.04}$ | $-1.64\pm 0.10$ | $0.38\pm 0.03$ | 1
HD 63005 | $21.24\pm 0.03$ | $20.17\pm 0.09$ | $21.31^{+0.03}_{-0.03}$ | $-0.84\pm 0.09$ | $0.68\pm 0.04$ | 1
HD 66788 | $21.23\pm 0.04$ | $19.72\pm 0.14$ | $21.26^{+0.04}_{-0.04}$ | $-1.24\pm 0.14$ | $0.55\pm 0.04$ | 1
HD 72754 | $21.17\pm 0.12$ | $20.35\pm 0.10$ | $21.28^{+0.10}_{-0.13}$ | $-0.63\pm 0.13$ | $0.76\pm 0.10$ | 2
HD 73882 | $21.11\pm 0.11$ | $21.08\pm 0.10$ | $21.57^{+0.08}_{-0.09}$ | $-0.19\pm 0.12$ | $0.68\pm 0.07$ | 2
HD 75309 | $21.10\pm 0.03$ | $20.16\pm 0.06$ | $21.19^{+0.03}_{-0.03}$ | $-0.73\pm 0.07$ | $0.58\pm 0.04$ | 1
HD 79186 | $21.18\pm 0.09$ | $20.72\pm 0.09$ | $21.41^{+0.07}_{-0.08}$ | $-0.39\pm 0.11$ | $0.69\pm 0.03$ | 2
HD 88115 | $21.03\pm 0.06$ | $19.25\pm 0.14$ | $21.04^{+0.06}_{-0.07}$ | $-1.49\pm 0.15$ | $0.39\pm 0.08$ | 1
HD 89137 | $21.03\pm 0.07$ | $20.02\pm 0.09$ | $21.11^{+0.06}_{-0.07}$ | $-0.79\pm 0.11$ | $0.45\pm 0.04$ | 1
HD 90087 | $21.19\pm 0.05$ | $19.88\pm 0.07$ | $21.23^{+0.05}_{-0.05}$ | $-1.05\pm 0.08$ | $0.36\pm 0.07$ | 1
HD 91597 | $21.40\pm 0.06$ | $19.70\pm 0.05$ | $21.42^{+0.06}_{-0.07}$ | $-1.42\pm 0.07$ | $0.44\pm 0.05$ | 2
HD 91651 | $21.15\pm 0.06$ | $19.07\pm 0.03$ | $21.16^{+0.06}_{-0.07}$ | $-1.79\pm 0.07$ | $0.27\pm 0.04$ | 2
HD 91824 | $21.12\pm 0.04$ | $19.81\pm 0.09$ | $21.16^{+0.04}_{-0.04}$ | $-1.05\pm 0.10$ | $0.33\pm 0.04$ | 1
HD 91983 | $21.15\pm 0.06$ | $20.10\pm 0.07$ | $21.22^{+0.05}_{-0.06}$ | $-0.82\pm 0.09$ | $0.31\pm 0.05$ | 1
HD 92554 | $21.34\pm 0.09$ | $19.24\pm 0.07$ | $21.35^{+0.09}_{-0.11}$ | $-1.81\pm 0.11$ | $0.15\pm 0.10$ | 1
HD 93129 | $21.47\pm 0.07$ | $20.21\pm 0.07$ | $21.52^{+0.06}_{-0.07}$ | $-1.00\pm 0.09$ | $0.49\pm 0.05$ | 1
HD 93205 | $21.36\pm 0.05$ | $19.73\pm 0.12$ | $21.38^{+0.05}_{-0.05}$ | $-1.35\pm 0.13$ | $0.33\pm 0.05$ | 1
HD 93222 | $21.47\pm 0.03$ | $19.77\pm 0.09$ | $21.49^{+0.03}_{-0.03}$ | $-1.42\pm 0.09$ | $0.33\pm 0.04$ | 1
HD 93843 | $21.30\pm 0.05$ | $19.62\pm 0.10$ | $21.32^{+0.05}_{-0.05}$ | $-1.40\pm 0.11$ | $0.50\pm 0.04$ | 1
HD 94493 | $21.10\pm 0.06$ | $20.09\pm 0.06$ | $21.18^{+0.05}_{-0.06}$ | $-0.79\pm 0.08$ | $0.32\pm 0.05$ | 1
HD 97175 | $20.96\pm 0.06$ | $20.10\pm 0.09$ | $21.07^{+0.05}_{-0.06}$ | $-0.66\pm 0.10$ | $0.49\pm 0.05$ | 1
HD 99857 | $21.27\pm 0.07$ | $20.30\pm 0.05$ | $21.35^{+0.06}_{-0.07}$ | $-0.75\pm 0.08$ | $0.45\pm 0.05$ | 1
HD 99890 | $21.12\pm 0.05$ | $19.56\pm 0.09$ | $21.14^{+0.05}_{-0.05}$ | $-1.28\pm 0.10$ | $0.16\pm 0.05$ | 1
HD 100199 | $21.18\pm 0.06$ | $20.15\pm 0.08$ | $21.25^{+0.05}_{-0.06}$ | $-0.80\pm 0.09$ | $0.41\pm 0.09$ | 1
HD 101190 | $21.24\pm 0.03$ | $20.43\pm 0.04$ | $21.36^{+0.02}_{-0.03}$ | $-0.63\pm 0.05$ | $0.51\pm 0.05$ | 1
HD 103779 | $21.17\pm 0.05$ | $19.90\pm 0.08$ | $21.21^{+0.05}_{-0.05}$ | $-1.01\pm 0.09$ | $0.29\pm 0.05$ | 1
HD 104705 | $21.15\pm 0.06$ | $20.02\pm 0.07$ | $21.21^{+0.05}_{-0.06}$ | $-0.89\pm 0.09$ | $0.40\pm 0.05$ | 1
HD 108639 | $21.36\pm 0.04$ | $20.01\pm 0.10$ | $21.40^{+0.04}_{-0.04}$ | $-1.09\pm 0.11$ | $0.34\pm 0.04$ | 1
HD 109399 | $21.11\pm 0.05$ | $20.01\pm 0.09$ | $21.17^{+0.05}_{-0.05}$ | $-0.86\pm 0.10$ | $0.44\pm 0.05$ | 1
HD 114886 | $21.34\pm 0.06$ | $20.30\pm 0.07$ | $21.41^{+0.05}_{-0.06}$ | $-0.81\pm 0.09$ | $0.47\pm 0.05$ | 1
HD 115455 | $21.38\pm 0.05$ | $20.55\pm 0.07$ | $21.49^{+0.04}_{-0.05}$ | $-0.64\pm 0.08$ | $0.41\pm 0.06$ | 1
HD 116852 | $20.96\pm 0.04$ | $19.75\pm 0.09$ | $21.01^{+0.04}_{-0.04}$ | $-0.96\pm 0.10$ | $0.48\pm 0.04$ | 1
HD 121968 | $20.58\pm 0.12$ | $18.70\pm 0.10$ | $20.59^{+0.12}_{-0.16}$ | $-1.59\pm 0.15$ | $0.26\pm 0.06$ | 2
HD 122879 | $21.31\pm 0.06$ | $20.33\pm 0.07$ | $21.39^{+0.05}_{-0.06}$ | $-0.76\pm 0.08$ | $0.51\pm 0.04$ | 1
HD 124314 | $21.41\pm 0.06$ | $20.41\pm 0.09$ | $21.49^{+0.05}_{-0.06}$ | $-0.78\pm 0.10$ | $0.47\pm 0.05$ | 1
HD 137595 | $20.97\pm 0.06$ | $20.59\pm 0.06$ | $21.23^{+0.04}_{-0.05}$ | $-0.34\pm 0.07$ | $0.85\pm 0.05$ | 1
HD 144965 | $20.97\pm 0.09$ | $20.74\pm 0.07$ | $21.31^{+0.06}_{-0.07}$ | $-0.27\pm 0.09$ | $1.07\pm 0.09$ | 1
HD 147683 | $21.20\pm 0.15$ | $20.68\pm 0.12$ | $21.41^{+0.11}_{-0.14}$ | $-0.42\pm 0.15$ | $0.56\pm 0.47$ | 2
HD 147888 | $21.68\pm 0.08$ | $20.45\pm 0.05$ | $21.73^{+0.07}_{-0.09}$ | $-0.98\pm 0.09$ | $1.15\pm 0.12$ | 1
HD 147933 | $21.63\pm 0.09$ | $20.57\pm 0.15$ | $21.70^{+0.08}_{-0.10}$ | $-0.83\pm 0.16$ | $1.09\pm 0.08$ | 2
HD 148422 | $21.24\pm 0.09$ | $20.15\pm 0.11$ | $21.31^{+0.08}_{-0.10}$ | $-0.85\pm 0.13$ | $0.21\pm 0.09$ | 1
HD 148937 | $21.48\pm 0.06$ | $20.68\pm 0.06$ | $21.60^{+0.05}_{-0.05}$ | $-0.62\pm 0.08$ | $0.56\pm 0.07$ | 1
HD 151805 | $21.33\pm 0.05$ | $20.35\pm 0.07$ | $21.41^{+0.04}_{-0.05}$ | $-0.76\pm 0.08$ | $0.54\pm 0.04$ | 1
HD 152590 | $21.37\pm 0.07$ | $20.47\pm 0.10$ | $21.47^{+0.06}_{-0.07}$ | $-0.70\pm 0.11$ | $0.52\pm 0.05$ | 1
HD 156359 | $20.80\pm 0.10$ | $18.07\pm 0.21$ | $20.80^{+0.10}_{-0.13}$ | $-2.43\pm 0.22$ | $0.24\pm 0.09$ | 1
HD 157857 | $21.26\pm 0.09$ | $20.68\pm 0.10$ | $21.44^{+0.07}_{-0.08}$ | $-0.46\pm 0.12$ | $0.62\pm 0.04$ | 2
HD 163522 | $21.14\pm 0.08$ | $19.66\pm 0.19$ | $21.17^{+0.08}_{-0.09}$ | $-1.21\pm 0.20$ | $0.34\pm 0.08$ | 1
HD 165246 | $21.41\pm 0.03$ | $20.14\pm 0.07$ | $21.45^{+0.03}_{-0.03}$ | $-1.01\pm 0.07$ | $0.78\pm 0.04$ | 1
HD 165955 | $21.10\pm 0.06$ |  $19.09\pm 0.07$bbThe H2 column density derived by Cartledge et al. (2004), and adopted in Jenkins (2009), for the line of sight to HD 165955 is too low. The value listed here is from our own analysis of the H2 data in this direction. | $21.11^{+0.06}_{-0.07}$ | $-1.72\pm 0.09$ | $0.42\pm 0.04$ | 2
HD 167402 | $21.13\pm 0.05$ | $20.06\pm 0.10$ | $21.20^{+0.05}_{-0.05}$ | $-0.84\pm 0.11$ | $0.25\pm 0.07$ | 1
HD 168076 | $21.65\pm 0.23$ | $20.68\pm 0.08$ | $21.73^{+0.20}_{-0.37}$ | $-0.75\pm 0.21$ | $0.68\pm 0.17$ | 2
HD 168941 | $21.18\pm 0.05$ | $20.11\pm 0.08$ | $21.25^{+0.04}_{-0.05}$ | $-0.84\pm 0.09$ | $0.68\pm 0.06$ | 1
HD 170740 | $21.09\pm 0.06$ | $20.86\pm 0.08$ | $21.43^{+0.05}_{-0.06}$ | $-0.27\pm 0.09$ | $0.87\pm 0.09$ | 1
HD 177989 | $20.99\pm 0.05$ | $20.16\pm 0.12$ | $21.10^{+0.05}_{-0.05}$ | $-0.64\pm 0.13$ | $0.66\pm 0.05$ | 1
HD 178487 | $21.22\pm 0.04$ | $20.51\pm 0.09$ | $21.36^{+0.04}_{-0.04}$ | $-0.55\pm 0.10$ | $0.74\pm 0.07$ | 1
HD 179407 | $21.20\pm 0.06$ | $20.23\pm 0.09$ | $21.28^{+0.05}_{-0.06}$ | $-0.75\pm 0.10$ | $0.63\pm 0.09$ | 1
HD 185418 | $21.19\pm 0.05$ | $20.72\pm 0.07$ | $21.41^{+0.04}_{-0.05}$ | $-0.39\pm 0.08$ | $0.74\pm 0.04$ | 1
HD 190918 | $21.38\pm 0.06$ | $19.84\pm 0.08$ | $21.40^{+0.06}_{-0.07}$ | $-1.26\pm 0.10$ | $0.46\pm 0.03$ | 2
HD 191877 | $21.03\pm 0.05$ | $20.00\pm 0.17$ | $21.10^{+0.05}_{-0.06}$ | $-0.80\pm 0.17$ | $0.53\pm 0.05$ | 1
HD 192035 | $21.20\pm 0.04$ | $20.63\pm 0.09$ | $21.39^{+0.04}_{-0.05}$ | $-0.46\pm 0.10$ | $0.74\pm 0.07$ | 1
HD 192639 | $21.29\pm 0.09$ | $20.73\pm 0.10$ | $21.48^{+0.07}_{-0.08}$ | $-0.45\pm 0.12$ | $0.64\pm 0.04$ | 2
HD 195455 | $20.61\pm 0.04$ | $18.42\pm 0.19$ | $20.62^{+0.04}_{-0.04}$ | $-1.89\pm 0.19$ | $0.39\pm 0.07$ | 1
HD 195965 | $20.92\pm 0.05$ | $20.28\pm 0.08$ | $21.08^{+0.04}_{-0.05}$ | $-0.50\pm 0.09$ | $0.45\pm 0.05$ | 1
HD 198478 | $21.32\pm 0.12$ | $20.76\pm 0.33$ | $21.51^{+0.16}_{-0.26}$ | $-0.45\pm 0.35$ | $0.75\pm 0.14$ | 1
HD 198781 | $20.93\pm 0.07$ | $20.48\pm 0.08$ | $21.16^{+0.05}_{-0.06}$ | $-0.38\pm 0.09$ | $0.79\pm 0.05$ | 1
HD 201345 | $21.00\pm 0.05$ | $19.36\pm 0.12$ | $21.02^{+0.05}_{-0.05}$ | $-1.36\pm 0.13$ | $0.34\pm 0.06$ | 1
HD 202347 | $20.83\pm 0.08$ | $19.98\pm 0.10$ | $20.94^{+0.07}_{-0.08}$ | $-0.66\pm 0.12$ | $0.66\pm 0.08$ | 1
HD 203374 | $21.20\pm 0.05$ | $20.67\pm 0.07$ | $21.40^{+0.04}_{-0.05}$ | $-0.43\pm 0.08$ | $0.66\pm 0.04$ | 1
HD 203938 | $21.48\pm 0.15$ | $21.00\pm 0.06$ | $21.70^{+0.10}_{-0.13}$ | $-0.40\pm 0.11$ | $0.99\pm 0.64$ | 2
HD 206267 | $21.22\pm 0.06$ | $20.85\pm 0.08$ | $21.49^{+0.05}_{-0.06}$ | $-0.34\pm 0.09$ | $0.80\pm 0.05$ | 1
HD 206773 | $21.09\pm 0.07$ | $20.41\pm 0.06$ | $21.24^{+0.05}_{-0.06}$ | $-0.53\pm 0.08$ | $0.60\pm 0.04$ | 1
HD 207198 | $21.28\pm 0.07$ | $20.79\pm 0.05$ | $21.50^{+0.05}_{-0.05}$ | $-0.41\pm 0.07$ | $0.80\pm 0.05$ | 1
HD 207308 | $21.20\pm 0.06$ | $20.80\pm 0.06$ | $21.45^{+0.04}_{-0.05}$ | $-0.35\pm 0.07$ | $0.82\pm 0.05$ | 1
HD 207538 | $21.27\pm 0.06$ | $20.85\pm 0.06$ | $21.52^{+0.04}_{-0.05}$ | $-0.36\pm 0.07$ | $0.85\pm 0.05$ | 1
HD 208440 | $21.24\pm 0.06$ | $20.28\pm 0.11$ | $21.33^{+0.05}_{-0.06}$ | $-0.75\pm 0.12$ | $0.65\pm 0.05$ | 1
HD 209339 | $21.20\pm 0.04$ | $20.13\pm 0.06$ | $21.27^{+0.04}_{-0.04}$ | $-0.84\pm 0.07$ | $0.50\pm 0.05$ | 1
HD 210809 | $21.31\pm 0.06$ | $19.97\pm 0.12$ | $21.35^{+0.06}_{-0.06}$ | $-1.08\pm 0.13$ | $0.31\pm 0.06$ | 1
HD 210839 | $21.24\pm 0.05$ | $20.80\pm 0.05$ | $21.48^{+0.04}_{-0.04}$ | $-0.38\pm 0.06$ | $0.78\pm 0.04$ | 1
HD 212791 | $21.11\pm 0.12$ | $19.42\pm 0.11$ | $21.13^{+0.12}_{-0.16}$ | $-1.41\pm 0.15$ | $0.57\pm 0.08$ | 2
HD 218915 | $21.20\pm 0.07$ | $20.16\pm 0.06$ | $21.27^{+0.06}_{-0.07}$ | $-0.81\pm 0.08$ | $0.39\pm 0.06$ | 1
HD 219188 | $20.72\pm 0.07$ | $19.38\pm 0.11$ | $20.76^{+0.07}_{-0.08}$ | $-1.08\pm 0.12$ | $0.37\pm 0.06$ | 1
HD 220057 | $20.95\pm 0.14$ | $20.27\pm 0.09$ | $21.10^{+0.11}_{-0.14}$ | $-0.53\pm 0.13$ | $0.69\pm 0.13$ | 1
HD 224151 | $21.35\pm 0.05$ | $20.55\pm 0.06$ | $21.47^{+0.04}_{-0.05}$ | $-0.62\pm 0.07$ | $0.49\pm 0.05$ | 1
HDE 232522 | $21.12\pm 0.04$ | $20.18\pm 0.08$ | $21.21^{+0.04}_{-0.04}$ | $-0.73\pm 0.09$ | $0.41\pm 0.04$ | 1
HDE 303308 | $21.41\pm 0.03$ | $20.20\pm 0.05$ | $21.46^{+0.03}_{-0.03}$ | $-0.96\pm 0.06$ | $0.42\pm 0.05$ | 1
HDE 308813 | $21.20\pm 0.06$ | $20.23\pm 0.07$ | $21.28^{+0.05}_{-0.06}$ | $-0.75\pm 0.08$ | $0.52\pm 0.06$ | 1
BD$+$35 4258 | $21.24\pm 0.03$ | $19.61\pm 0.11$ | $21.26^{+0.03}_{-0.03}$ | $-1.35\pm 0.11$ | $0.38\pm 0.06$ | 1
BD$+$53 2820 | $21.35\pm 0.05$ | $20.10\pm 0.11$ | $21.40^{+0.05}_{-0.05}$ | $-1.00\pm 0.12$ | $0.44\pm 0.07$ | 1
CPD$-$59 2603 | $21.43\pm 0.04$ | $20.02\pm 0.12$ | $21.46^{+0.04}_{-0.04}$ | $-1.14\pm 0.12$ | $0.51\pm 0.05$ | 1
CPD$-$59 4552 | $21.28\pm 0.05$ | $20.45\pm 0.08$ | $21.39^{+0.04}_{-0.05}$ | $-0.64\pm 0.09$ | $0.40\pm 0.05$ | 1
CPD$-$69 1743 | $21.16\pm 0.04$ | $19.85\pm 0.15$ | $21.20^{+0.04}_{-0.04}$ | $-1.05\pm 0.15$ | $0.40\pm 0.07$ | 1

References. — (1) Jenkins (2019), (2) Jenkins (2009).

Figure 10: Determination of $N({\rm H}_{2})$ toward HD 165955. A synthetic
spectrum (red curve) is compared to observations of the $R$(0), $R$(1), and
$P$(1) lines of the $B$$-$$X$ (4$-$0) band. The solid red curve corresponds to
a total H2 column density of $\log N({\rm H}_{2})=19.09$, while the dashed red
curves indicate variations in the column density of $\pm 0.07$ dex. Figure 11:
Correlation between $\log N$(Cl i) and $\log N$(H2). Solid black circles
represent chlorine measurements obtained in this investigation from STIS and
FUSE observations; open circles represent measurements obtained in previous
studies from Copernicus observations (Moomey et al., 2012; Brown, 2015). Solid
red circles represent Cl i and H2 measurements obtained for high-redshift
absorption systems toward background quasars (Balashev et al., 2015). The
dashed black line represents a least-squares linear fit to the STIS+FUSE
sample, while the solid black line indicates a linear fit to the entire
Galactic sample (excluding the four discrepant sight lines with very low H2
column densities). The solid red line shows a least-squares linear fit to the
combined Galactic and extragalactic sample (excluding the four outliers).

##  4 ANALYSIS

###  4.1 Column Density Correlations

Owing to the unique chemistry that links neutral chlorine to molecular
hydrogen, the column densities of Cl i and H2 have long been expected to be
well correlated (e.g., Jura & York, 1978; Harris & Bromage, 1984; Moomey et
al., 2012). Most recently, Moomey et al. (2012), using Cl measurements from
the Copernicus satellite, found a correlation between $N({\rm
Cl~{}\textsc{i}})$ and $N({\rm H}_{2})$ that had a slope of $0.92\pm 0.19$ (in
$\log$-$\log$ space) and a correlation coefficient greater than 0.6. Balashev
et al. (2015) found that a similar relationship holds for high-redshift damped
Lyman-$\alpha$ (DLA) systems where both H2 and Cl i are detected. Other
correlations, such as that between $N({\rm Cl}_{\rm tot})$ and $N({\rm H}_{\rm
tot})$, have also been noted in the literature (Harris & Bromage, 1984; Moomey
et al., 2012). In light of our greatly expanded sample of Cl i and Cl ii
column densities (Table 5), it is worthwhile to revisit these correlations.

In Figure 11, we compare our new determinations of neutral chlorine column
densities (from STIS and FUSE observations) with the corresponding molecular
hydrogen column densities from Table 3.3. Also included in Figure 11 are the
Cl i measurements obtained in previous studies based on Copernicus
observations (Moomey et al., 2012; Brown, 2015)555Brown (2015) derived Cl i
and Cl ii column densities from Copernicus observations for a sample of sight
lines with very low H2 column densities as part of a Masters Thesis at the
University of Toledo. For convenience, these measurements are reported in
Appendix A. and the Cl i and H2 column densities derived for high-redshift
absorption systems toward background quasars (Balashev et al., 2015). A strong
correlation between $N({\rm Cl~{}\textsc{i}})$ and $N({\rm H}_{2})$ is found
for Galactic sight lines with $\log N({\rm H}_{2})\gtrsim 18$. Least-squares
linear fits666The least-squares linear fits described in this paper were
performed using the IDL procedure FITEXY, which accounts for uncertainties in
both the $x$ and $y$ coordinates (Press et al., 2007). to the “STIS+FUSE”
sample and the sample that also includes Copernicus measurements (“all
Galactic”) yield linear correlation coefficients of $\sim$0.8 and slopes of
$\sim$0.8–0.9 (see Figure 11 and Table 7), consistent with the results of
Moomey et al. (2012). Four sight lines with $\log N({\rm H}_{2})<17$ do not
follow the trend exhibited by the other Galactic sight lines and are not
included in the linear fits.

Table 7: Column Density Correlations

$Y$ | $X$ | $A$ | $B$ |  $r$aaLinear correlation coefficient. |  $N$bbNumber of sight lines in the sample. | Sample
---|---|---|---|---|---|---
Cl i | H2 | $0.824\pm 0.016$ | $-2.59\pm 0.31$ | 0.840 | 98 | STIS+FUSE
|  | $0.862\pm 0.015$ | $-3.37\pm 0.31$ | 0.841 | 118 | all GalacticccExcludes outliers: $\gamma$ Cas, $\psi$ Ori, $\epsilon$ Ori, and $\kappa$ Ori.
|  | $0.906\pm 0.012$ | $-4.26\pm 0.23$ | 0.897 | 130 | Galactic+DLAsccExcludes outliers: $\gamma$ Cas, $\psi$ Ori, $\epsilon$ Ori, and $\kappa$ Ori.
Cl i | Htot | $2.862\pm 0.106$ |  $-46.96\pm 2.25$0 | 0.731 | 98 | STIS+FUSE
|  | $2.478\pm 0.061$ |  $-38.74\pm 1.29$0 | 0.828 | 119 | all Galactic
Cltot | Htot | $1.065\pm 0.057$ | $-8.26\pm 1.21$ | 0.845 | 62 | STIS+FUSE
|  | $1.030\pm 0.022$ | $-7.52\pm 0.45$ | 0.931 | 88 | all Galactic

Note. — $\log N(Y)=B+A\times\log N(X)$

As already discussed by Balashev et al. (2015), the relationship between Cl i
and H2 exhibited by high-redshift DLAs is very similar to that found for
Galactic sight lines. A least-squares linear fit to the combined Galactic and
extragalactic sample (“Galactic+DLAs”) yields a slope of $0.906\pm 0.012$ and
a linear correlation coefficient of 0.897 (Table 7). The extragalactic
measurements help to extend the trend of Cl i versus H2 to lower neutral
chlorine column densities (i.e., below $\log N({\rm Cl~{}\textsc{i}})\sim
12.5$). The only Galactic sight lines with such low values of $N({\rm
Cl~{}\textsc{i}})$ (i.e., $\gamma$ Cas, $\psi$ Ori, $\epsilon$ Ori, and
$\kappa$ Ori) have H2 column densities that are much lower than the values
that would be predicted by the linear trends. These sight lines probe
molecular gas in the transition region where H2 is not yet fully self-
shielded. A similar pattern is seen when the column densities of other trace
neutral species (e.g., K i and Na i) are plotted versus $N({\rm H}_{2})$
(Welty & Hobbs, 2001). The K i and Na i column densities examined by Welty &
Hobbs (2001) exhibit a plateau for molecular hydrogen column densities in the
range $15\lesssim\log N({\rm H}_{2})\lesssim 19$ (see their Figures 26 and
27).

Figure 12: Correlation between $\log N$(Cl i) and $\log N$(Htot). The plotting
symbols have the same meaning as in Figure 11. Two sight lines with
anomalously low Cl i column densities are identified. (These sight lines were
not excluded from the linear fits.)

Welty & Hobbs (2001) find nearly linear relationships for $N({\rm
K~{}\textsc{i}})$ versus $N({\rm H}_{2})$ and $N({\rm Na~{}\textsc{i}})$
versus $N({\rm H}_{2})$ for $\log N({\rm H}_{2})>18.5$, very similar to our
findings for Cl i. The fact that the slope of the relationship between $N({\rm
Cl~{}\textsc{i}})$ and $N({\rm H}_{2})$ in the high column density regime is
slightly less than one is most likely related to the increase in the fraction
of hydrogen in molecular form in the portion of the cloud where the H2
resides. As discussed below (see Section 4.3), the gas-phase abundance (i.e.,
depletion) of chlorine shows very little dependence on the molecular hydrogen
fraction. Thus, while the ratio of neutral chlorine to total hydrogen remains
roughly constant in the molecular portion of the cloud, the fraction of
molecular hydrogen increases as a function of $N({\rm H}_{2})$.

Welty & Hobbs (2001) find nearly quadratic relationships between the column
densities of K i and Na i and the total hydrogen column density. A quadratic
dependence of $N({\rm K~{}\textsc{i}})$ and $N({\rm Na~{}\textsc{i}})$ on
$N({\rm H}_{\rm tot})$ is expected if the ionization equilibrium is dominated
by photoionization and radiative recombination, if the electron fraction is
roughly constant, and if the individual clouds have a roughly uniform
thickness (Hobbs, 1974; Welty & Hobbs, 2001). In Figure 12, we plot the
observed trend of $N({\rm Cl~{}\textsc{i}})$ versus $N({\rm H}_{\rm tot})$.
Least-squares linear fits to the “STIS+FUSE” and “all Galactic” samples yield
slopes in the range 2.5–2.9 (see Table 7), much steeper than the corresponding
trends involving K i and Na i. The steeper relationship for Cl i is due to the
conversion of Cl+ to Cl0 in molecule-rich sight lines, where Cl i becomes the
dominant ion. Indeed, the fraction of chlorine in neutral form, $f({\rm
Cl~{}\textsc{i}})=N({\rm Cl~{}\textsc{i}})/N({\rm Cl}_{\rm tot})$, increases
steeply as a function of $N({\rm H}_{\rm tot})$.

Two sight lines are identified in Figure 12 as having anomalously low Cl i
column densities relative to the total amount of hydrogen along the lines of
sight. The two sight lines, HD 147933 ($\rho$ Oph A) and HD 147888 ($\rho$ Oph
D), are part of the same stellar system and probe the same complex of diffuse
molecular clouds (e.g., Snow et al., 2008). Welty & Hobbs (2001) describe
these sight lines, and several others in the Sco-Oph region, as being
“discrepant” because they exhibit several anomalous column density ratios. For
example, both $\rho$ Oph A and $\rho$ Oph D have anomalously low K i and Na i
column densities relative to $N({\rm H}_{\rm tot})$ (Welty & Hobbs, 2001; Snow
et al., 2008). Furthermore, while both sight lines exhibit severe depletions
of many different elements from the gas-phase (e.g., Ritchey et al., 2023),
the molecular hydrogen fractions are unusually low. An enhanced UV radiation
field may be responsible for the anomalously low column densities of Cl i, K
i, Na i, and H2 in these directions. Indeed, from an analysis of carbon
ionization balance, Jenkins & Tripp (2011) find that the UV radiation field
toward HD 147888 is enhanced by a factor of $\sim$16 over the average
interstellar field.

Figure 13: Correlation between $\log N$(Cltot) and $\log N$(Htot). The
plotting symbols have the same meaning as in Figure 11. Several outliers are
indicated. (These sight lines were not excluded from the linear fits.)

In Figure 13, we plot the observed relationship between $N({\rm Cl}_{\rm
tot})$ and $N({\rm H}_{\rm tot})$. Least-squares linear fits to the
“STIS+FUSE” and “all Galactic” samples yield slopes of $\sim$1.0–1.1 and
linear correlation coefficients of $\sim$0.8–0.9. The tight linear correlation
between the total chlorine column density and the total hydrogen column
density is an indication that the abundance and/or depletion of chlorine does
not vary appreciably with $N({\rm H}_{\rm tot})$. The weighted mean value of
$\log[N({\rm Cl}_{\rm tot})/N({\rm H}_{\rm tot})]$ for the “all Galactic”
sample is $-6.893\pm 0.008$, which is consistent with (but somewhat higher
than) the value of $-6.99\pm 0.04$ reported by Moomey et al. (2012). The sight
lines to HD 147888 and HD 147933 do not appear to be significant outliers in
Figure 13 like they are in Figure 12. Still, the total Cl column densities
toward the $\rho$ Oph stars are $\sim$0.4–0.6 dex below that determined for
the line of sight to HD 168076, which has a similar total hydrogen column
density. The sight line to HD 64740 does appear to be an outlier in Figure 13,
although the hydrogen column density is not very well determined in this
direction (Bohlin et al., 1983; Jenkins, 2009).

Figure 14: Gas-phase P and Cl abundances as a function of the sight line
depletion factor ($F_{*}$). Solid squares represent abundances derived in this
work from STIS and FUSE observations; open diamonds represent measurements
obtained in previous studies from Copernicus and GHRS observations (Jenkins,
2009; Moomey et al., 2012; Brown, 2015). Grey symbols denote sight lines where
$\sigma(F_{*})\geq 0.30$. The solid orange lines represent linear fits to the
data derived according to the methodology of Jenkins (2009). The shaded gray
regions indicate the 1$\sigma$ errors on the fit parameters. The horizontal
dotted line in each panel gives the adopted solar system abundance from
Lodders (2003).

###  4.2 Phosphorus and Chlorine Depletion Parameters

The main objective of this investigation is the redetermination of depletion
parameters for the elements phosphorus and chlorine. Previous studies of P
abundances (e.g., Lebouteiller et al., 2005; Jenkins, 2009) used now outdated
oscillator strengths for the important P ii transitions, leading to an
overestimation of P ii column densities, and, consequently, an underestimation
of the depletion of P onto interstellar grains. In his analysis of Cl
depletion, Jenkins (2009) considered only Cl ii column densities, neglecting
any contribution from Cl i to the derived total Cl abundances. However, of the
68 sight lines in Table 5 with measurements (or upper limits) for both Cl i
and Cl ii, nearly half (32) have $N({\rm Cl~{}\textsc{i}})>N({\rm
Cl~{}\textsc{ii}})$. Moreover, the neutral chlorine fraction, $f({\rm
Cl~{}\textsc{i}})$, increases systematically with $f({\rm H}_{2})$ (see
Section 4.3). The neglect of Cl i will therefore significantly impact the
interpretation of any perceived trends involving the depletion of Cl onto
interstellar grains.

As in Ritchey et al. (2018), we derive element depletion parameters following
the methodology of Jenkins (2009). In this formalism, the logarithmic
depletion of element $X$, defined as $[X/{\rm H}]=\log(X/{\rm H})-\log(X/{\rm
H})_{\sun}$, depends on the sight-line depletion strength factor $F_{*}$
according to:

$[X/{\rm H}]=B_{X}+A_{X}(F_{*}-z_{X}),$ (1)

where the depletion parameters $A_{X}$, $B_{X}$, and $z_{X}$ are unique to
each element. The slope parameter $A_{X}$ indicates how quickly a given
element becomes depleted as the growth of dust grains progresses in
interstellar clouds. The intercept parameter $B_{X}$ indicates the expected
depletion of element $X$ at $F_{*}=z_{X}$, where $z_{X}$ represents a weighted
mean value of $F_{*}$ for the particular set of sight lines with depletion
measurements available for the element. Values of the coefficients $A_{X}$ and
$B_{X}$ are obtained for each element through the evaluation of a least-
squares linear fit, with $[X/{\rm H}]$ as the dependent variable and $F_{*}$
the independent variable. (The reason for the additional term involving
$z_{X}$ in Equation (1) is that for a particular choice of $z_{X}$ there is a
near zero covariance between the formal fitting errors for the solutions of
$A_{X}$ and $B_{X}$; see Jenkins, 2009).

Values of the depletion strength factor $F_{*}$ for the sight lines in our
combined P and Cl sample are obtained primarily from Jenkins (2019), although,
in some cases, we use the $F_{*}$ values provided by Jenkins (2009). The
adopted values are listed in Table 3.3. The $F_{*}$ values from Jenkins (2019)
are preferred because they were derived solely from depletion measurements for
Mg and Mn and thus are entirely independent from the elements considered in
this investigation. Regardless of our preferrence, however, the two sets of
$F_{*}$ values are comparable. For the 70 sight lines in our sample that have
depletion strength factors listed in both Jenkins (2019) and Jenkins (2009),
the differences in the values are typically at the 1$\sigma$ or 2$\sigma$
level. (The mean absolute difference is 1.1$\sigma$.)

In Figure 14, we plot the gas-phase P and Cl abundances777Note that $\log({\rm
P}/{\rm H})=\log N({\rm P~{}\textsc{ii}})-\log N({\rm H}_{\rm tot})$, while
$\log({\rm Cl}/{\rm H})=\log N({\rm Cl}_{\rm tot})-\log N({\rm H}_{\rm tot})$.
as a function of $F_{*}$. Measurements obtained in this investigation (from
the data presented in Tables 3, 5, and 3.3) are represented by solid squares,
while measurements for additional sight lines examined in previous
investigations are represented by open diamonds. In the case of phosphorus, we
include the P ii measurements tabulated by Jenkins (2009) for sight lines not
included in our STIS sample. Most of these additional P ii measurements were
derived from Copernicus or Goddard High Resolution Spectrograph (GHRS)
observations (see Jenkins, 2009, and references therein). All of the P ii
column densities obtained from Jenkins (2009) were corrected so as to be
consistent with the set of $f$-values adopted in this investigation (Table
2.1). However, many of these previously reported P abundances have large
associated uncertainties. We therefore retained only those abundance
measurements with logarithmic uncertainties less than or equal to 0.176 dex
(corresponding to a relative uncertainty of 50% or better). For our analysis
of chlorine depletions, we include, in addition to our own measurements, the
total Cl column densities reported by Moomey et al. (2012) and Brown (2015,
see also Appendix A), which were based on Copernicus observations.

Table 8: Element Depletion Parameters Elem. | log ($X$/H)☉aaRecommended solar system abundance from Lodders (2003). | $A_{X}$ | $B_{X}$ | $z_{X}$ | [$X$/H]0 | [$X$/H]1 | $\chi^{2}$ | $\nu$
---|---|---|---|---|---|---|---|---
P | $-6.46\pm 0.04$ | $-0.776\pm 0.035$ | $-0.418\pm 0.041$ | 0.520 | $-0.015\pm 0.045$ | $-0.791\pm 0.044$ | 297.4 | 126
Cl | $-6.67\pm 0.06$ | $-0.238\pm 0.046$ | $-0.223\pm 0.061$ | 0.593 | $-0.082\pm 0.067$ | $-0.320\pm 0.063$ | 182.7 | 81

Depletion parameters for P and Cl were determined through least-squares linear
fits to the trends of $[{\rm P}/{\rm H}]$ and $[{\rm Cl}/{\rm H}]$ versus
$F_{*}$, adopting the functional form expressed in Equation (1). The resulting
values of $A_{X}$, $B_{X}$, and $z_{X}$ for the two elements are given in
Table 8. The fits themselves are depicted by solid orange lines in Figure 14
with shaded gray regions indicating the 1$\sigma$ errors on the fit
parameters. The horizontal dashed lines in Figure 14 represent the adopted
solar system abundances of P and Cl (Lodders, 2003, see Table 8). Since the
values obtained for the $B_{X}$ parameters depend on the specific choice of
reference abundances, the uncertainties in the solar system abundances were
added in quadrature to the formal fitting errors for $B_{\rm P}$ and $B_{\rm
Cl}$ to arrive at the uncertainties listed for these parameters in Table 8.
(Note that the slope parameters $A_{X}$ are entirely independent of the
adopted solar reference abundances.)

The intercept parameters $B_{X}$ are sensitive to the specific set of
depletion measurements available for a given element (because they depend on
$z_{X}$, which will vary from one element to the next). Thus, in order to
facilitate a more straightforward comparison of the depletion results for
different elements, it is important to evaluate two additional depletion
parameters. The two parameters:

$[X/{\rm H}]_{0}=B_{X}-A_{X}z_{X}$ (2)

and

$[X/{\rm H}]_{1}=B_{X}+A_{X}(1-z_{X}),$ (3)

indicate the expected depletions at $F_{*}=0$ and $F_{*}=1$, respectively. We
interpret the $[X/{\rm H}]_{0}$ values as representing the initial amounts of
depletion present in the diffuse ISM before significant grain growth has
occurred (or after the outer portions of the grains have been destroyed by the
passage of an interstellar shock). The $[X/{\rm H}]_{1}$ values represent the
depletions associated with a prototypical diffuse molecular cloud. (Jenkins
(2009) used the $v_{\sun}=-15$ km s-1 component toward $\zeta$ Oph as the
standard for a cloud with $F_{*}=1$.) The values we obtain for $[{\rm P}/{\rm
H}]_{0}$, $[{\rm P}/{\rm H}]_{1}$, $[{\rm Cl}/{\rm H}]_{0}$, and $[{\rm
Cl}/{\rm H}]_{1}$ are provided in Table 8. (The errors in these quantities
were determined according to the relations given in Jenkins (2009)).

Figure 15: Abundances of Cl i (left panel) and Cl ii (right panel) plotted as
a function of the sight line depletion factor ($F_{*}$). Squares represent
abundances derived in this work from STIS and FUSE observations; diamonds
represent measurements obtained in previous studies from Copernicus and GHRS
observations (Jenkins, 2009; Moomey et al., 2012; Brown, 2015). Grey symbols
denote sight lines where $\sigma(F_{*})\geq 0.30$. Blue symbols are used to
identify the discrepant Sco-Oph sight lines (see Welty & Hobbs, 2001; Welty &
Crowther, 2010). The horizontal dotted line in each panel gives the adopted
solar system abundance from Lodders (2003).

The last two columns in Table 8 give the chi-squared value ($\chi^{2}$)
associated with each of the linear fits along with the number of degrees of
freedom $\nu$ (i.e., the number of observations minus two). In each case, the
reduced chi-squared value ($\chi^{2}/\nu$) is relatively high (i.e., 2.4 and
2.3 for the P and Cl fits, respectively). This could indicate that there is
real intrinsic scatter in the gas-phase P and Cl abundances at a given value
of $F_{*}$. Alternatively, the poor goodness-of-fit statistics could mean that
the errors associated with the P and Cl abundances have been underestimated.
The depletion trend for Cl, in particular, seems irregular. Some sight lines
with moderate-to-large depletion strengths (such as HD 157857 and HD 206267)
have relatively low gas-phase Cl abundances, while others with similar or even
larger depletion strengths (such as X Per, HD 207308, and HD 207538) have much
higher total Cl abundances. The sight lines with the largest values of $F_{*}$
(HD 37903, HD 147888, and HD 147933) all have Cl abundances that are lower
than expected based on the linear trend shown in Figure 14.

An important fundamental result of our depletion analysis is that the
“initial” depletions of P and Cl (i.e., the $[X/{\rm H}]_{0}$ values) no
longer indicate that the P and Cl abundances are “supersolar” at $F_{*}=0$, in
contrast to the results presented in Jenkins (2009). Both $[{\rm P}/{\rm
H}]_{0}$ and $[{\rm Cl}/{\rm H}]_{0}$ are slightly less than zero (but
consistent with zero at approximately the 1$\sigma$ level), indicating very
little (if any) P and Cl depletion in the low-density ISM. The depletion slope
for P is relatively steep, however ($A_{\rm P}\approx-0.8$), while the slope
for Cl is rather shallow ($A_{\rm Cl}\approx-0.2$). These results will be
discussed further, and compared with the results for many other elements, in
Section 5.1.

###  4.3 Additional Insights on Depletion

We can gain additional insight into the possible cause of the irregular
depletion trend for Cl by examining separately the abundance trends for
neutral and singly-ionized Cl. In Figure 15, we plot the abundances of Cl i
and Cl ii against the sight line depletion strength factors. In these plots,
we have explicitly identified (with blue symbols) the discrepant Sco-Oph sight
lines discussed in Welty & Hobbs (2001) (see also Welty & Crowther, 2010;
Welty et al., 2020). These sight lines include: 1 Sco, $\pi$ Sco, $\delta$
Sco, $\beta^{1}$ Sco, $\omega^{1}$ Sco, $\nu$ Sco, $\sigma$ Sco, $\rho$ Oph A,
and $\rho$ Oph D. For most of the sight lines in our sample, we find a very
steep increase in the Cl i abundance with increasing $F_{*}$ (presumably due
to the neutralization of Cl+ in increasingly depleted H2-rich gas). However,
the discrepant Sco-Oph sight lines (and several other sight lines, including
$\kappa$ Ori, HD 37903, HD 72754, HD 144965, HD 165246, HD 206267, HD 207198,
and HD 210839) do not follow this trend, exhibiting Cl i abundances that are
much lower than expected for sight lines having relatively large values of
$F_{*}$.888It should be noted that the sight-line depletion strength factor
$F_{*}$ increases systematically with increasing $f({\rm H}_{2})$ for $\log
f({\rm H}_{2})>-1.0$ if the discrepant Sco-Oph sight lines are excluded (see
Welty et al., 2020).

As shown in the righthand panel of Figure 15, the Cl ii abundance drops
precipitously with increasing $F_{*}$ (again, owing to the rapid
neutralization of Cl+ ions as the H2 concentration increases). The discrepant
Sco-Oph sight lines (plus several others, such as $\psi$ Ori, $\kappa$ Ori, HD
37903, and HD 165246) show enhanced Cl ii abundances relative to other sight
lines with similar values of $F_{*}$. The Sco-Oph sight lines, and the other
discrepant sight lines, likely probe regions with enhanced UV radiation
fields. This is certainly true of HD 37903, which probes the photodissociation
region (PDR) associated with the reflection nebula NGC 2023. Likewise, HD
206267 is the exciting star of the H ii region IC 1396 and HD 165246 probes
gas in the outskirts of the Lagoon Nebula (M8). An enhanced UV radiation field
would significantly reduce the neutral chlorine abundance (both through the
photoionization of Cl0 and the photodissociation of H2), and would enhance the
abundance of Cl+. These effects could then account for the anomalous Cl i and
Cl ii abundances seen along the discrepant sight lines noted above.

Figure 16: Gas-phase abundances of Cl i (upper left), Cl ii (upper right),
total Cl (lower left), and P ii (lower right) plotted as a function of the
line-of-sight molecular hydrogen fraction. Squares represent abundances
derived in this work from STIS and FUSE observations; diamonds represent
measurements obtained in previous studies from Copernicus and GHRS
observations (Jenkins, 2009; Moomey et al., 2012; Brown, 2015). Blue symbols
are used to identify the discrepant Sco-Oph sight lines (see Welty & Hobbs,
2001; Welty & Crowther, 2010). The black dotted line in each panel indicates
the adopted solar system abundance from Lodders (2003). The blue dashed line
and light blue shaded region represent the mean and standard deviation of the
abundances for sight lines with $\log f({\rm H}_{2})<-3.0$. Solid lines
indicate linear fits with approximate slopes of $+0.9$ for Cl i (for sight
lines with $\log f({\rm H}_{2})>-3.0$) and $-0.8$ for Cl ii and P ii (for
sight lines with $\log f({\rm H}_{2})>-1.0$).

If the Sco-Oph sight lines, and the other discrepant sight lines, are excluded
from consideration, and the Cl depletion trend shown in Figure 14 is re-
evaluated, the slope of the relation between the gas-phase Cl abundance and
the sight line depletion factor becomes completely flat (i.e., $A_{\rm
Cl}=+0.00\pm 0.05$). Conversely, if only the discrepant sight lines are
considered, the slope of the relation is relatively steep (i.e., $A_{\rm
Cl}=-0.59\pm 0.18$). If the same discrepant sight lines are excluded from the
depletion analysis for P, the slope of the relation between the gas-phase P
abundance and $F_{*}$ is relatively unchanged (i.e., $A_{\rm P}=-0.74\pm
0.04$). However, if only the discrepant sight lines are included, the
depletion slope for P becomes significantly steeper (i.e., $A_{\rm P}=-1.49\pm
0.18$). For both P and Cl, the gas-phase abundances are much better correlated
with $F_{*}$ when only the discrepant sight lines are considered. (For these
fits, the $\chi^{2}/\nu$ values are 0.7 for P and 1.1 for Cl.)

The implication is that the depletion rates are enhanced in regions with
enhanced ionization. This would then point to the importance of ion-grain
reactions in driving the depletion of elements from the gas phase. The
accretion of atoms onto grain surfaces may be enhanced by the presence of very
small negatively-charged dust grains (Weingartner & Draine, 1999). Larger
grains tend to be positively charged, yet most of the abundant refractory
elements in the ISM are singly ionized. The cross sections for collisions
between ions and small grains are therefore enhanced, while those for larger
grains are diminished (see Weingartner & Draine, 1999). Our results indicate
that there is no enhanced depletion of Cl in relatively dense (heavily
depleted) regions where Cl is predominantly neutral (e.g., toward X Per, HD
207308, and HD 207538). However, there is enhanced Cl depletion in heavily
depleted regions where Cl is mostly ionized (e.g., toward HD 37903, HD 147888,
and HD 147933).

We can also examine the relationships that exist between the gas-phase P and
Cl abundances and the (line-of-sight) molecular hydrogen fractions. In Figure
16, we plot the gas-phase abundances of Cl i, Cl ii, total Cl, and P ii as a
function of $f({\rm H}_{2})$.999Note that in all four panels of Figure 16,
there are sight lines that are not shown because they have $\log f({\rm
H}_{2})<-3.0$. The gas-phase abundances in these directions are generally
consistent with a single value for a given species such that the abundances
exhibit a “plateau” at low molecular fraction. As expected, the neutral
chlorine abundance (upper left panel) increases systematically with the
molecular hydrogen fraction. A linear fit indicates a slope of $\sim$0.9,
similar to results for other trace neutral species, such as Na i and C i
(e.g., Welty et al., 2020). There is considerable scatter in this relation,
however, presumably because the total hydrogen column densities include
varying amounts of atomic hydrogen not associated with the molecular portions
of the interstellar clouds (see Section 4.4). Conversely, the abundance of
singly-ionized chlorine (upper right panel of Figure 16) shows a steady
decrease with increasing $f({\rm H}_{2})$, especially for sight lines with
$\log f({\rm H}_{2})\gtrsim-1.0$. The observed trends of Cl i and Cl ii
(relative to total H) versus $f({\rm H}_{2})$ are a direct consequence of the
rapid conversion of Cl+ to Cl0 in molecule-rich gas (Jura, 1974; Neufeld &
Wolfire, 2009).

While the Cl i abundance increases with increasing $f({\rm H}_{2})$, and the
Cl ii abundance decreases, the total Cl abundance (lower left panel of Figure
16) shows virtually no dependence on the line-of-sight molecular fraction As
already mentioned, there are some sight lines with unusually low total Cl
abundances (such as HD 37903, HD 147888, HD 147933, HD 157857, and HD 206267),
but the overall trend is exceptionally flat. In contrast, the gas-phase P
abundances (lower right panel) show a relatively constant value for $\log
f({\rm H}_{2})\lesssim-1.0$ (with the exception of the Sco-Oph sight lines)
and a sharp decline for $\log f({\rm H}_{2})\gtrsim-1.0$. This behavior is
typical of elements that exhibit moderate-to-strong depletions in cold
interstellar clouds. For example, similar trends versus $f({\rm H}_{2})$ are
seen for the elements B (Ritchey et al., 2011), Mg (Cartledge et al., 2006;
Jensen & Snow, 2007b; Welty et al., 2020), Ti (Welty & Crowther, 2010; Welty
et al., 2020), and Fe (Jensen & Snow, 2007a; Welty et al., 2020). Such trends
are clear indications of the increase in gas-phase element depletions in
increasingly dense, molecule-rich gas.

The discrepant Sco-Oph sight lines are prominent outliers in the plot showing
P/H versus $f({\rm H}_{2})$ (lower right panel of Figure 16). These sight
lines are known to exhibit strong depletions of (for example) Mg, P, Ti, and
Fe (e.g., Welty & Crowther, 2010; Welty et al., 2020, this work), yet contain
relatively little H2. An enhanced UV radiation field may be the source of
these discrepancies. An enhanced flux of UV radiation (for example, if the
absorbing gas is situated in close proximity to the background star) would
contribute to H2 photodissociation and would raise the ionization level of
atomic species. An increase in the concentration of ionized species in the
denser parts of the clouds could then lead to enhanced depletion through ion-
grain reactions. These results tend to corroborate our findings on Cl
depletion, which indicate that enhanced depletion of Cl generally only occurs
along sight lines with elevated levels of Cl ionization.

###  4.4 Neutral Chlorine Fractions

Because the first ionization potential of chlorine is less than that of
hydrogen, chlorine is predominantly ionized in low-density, diffuse atomic
gas. However, in molecular gas (i.e., in regions where the H2 molecules are
optically thick), Cl+ reacts rapidly with H2, yielding HCl+ and H. The
molecular ion HCl+ then undergoes dissociative recombination or reacts further
with H2, yielding H2Cl+ and H. Finally, the H2Cl+ is destroyed through
dissociative recombination, producing HCl, Cl, and H (see Jura, 1974; Neufeld
& Wolfire, 2009). Since the rate coefficient for the initial reaction
($k=1.0\times 10^{-9}$ cm3 s-1), when combined with a typical density for
neutral diffuse gas ($n_{\rm H}\gtrsim 10$ cm-3), is considerably larger than
the Cl photoionization rate for the average interstellar radiation field
($\Gamma\lesssim 4.8\times 10^{-11}$ s-1), chlorine will be predominantly
neutral wherever H2 is abundant.

With these considerations, Jura & York (1978) proposed a simple model to
account for observations of Cl i and Cl ii in the ISM. Their model consists of
two zones: (1) a region where the H2 molecules are optically thick, chlorine
is entirely neutral, and there is some associated atomic hydrogen, and (2) a
region where the H2 molecules are optically thin, chlorine is ionized, and the
hydrogen is predominantly (neutral) atomic. Following Sonnentrucker et al.
(2002), Equation (A7) in Jura & York (1978) may be written

$N_{1}({\rm H~{}\textsc{i}})=N({\rm H}_{\rm tot})[f({\rm
Cl~{}\textsc{i}})-f({\rm H}_{2})],$ (4)

where $N_{1}({\rm H~{}\textsc{i}})$ is the column density of H i that is
associated with the optically-thick H2 region, $N({\rm H}_{\rm tot})$ is the
total column density of hydrogen along the line of sight, $f({\rm
Cl~{}\textsc{i}})$ is the neutral chlorine fraction for the line of sight, and
$f({\rm H}_{2})$ is the line-of-sight molecular hydrogen fraction. With this
simple model, it is possible to use integrated line-of-sight column densities
of Cl i, Cl ii, H i, and H2 to derive estimates for the total hydrogen column
density and molecular hydrogen fraction associated with just the optically-
thick H2 region.

We apply Equation (4) to all of the sight lines in our chlorine sample with
measurements of both Cl i and Cl ii (Table 5), using values of $N({\rm H}_{\rm
tot})$ and $f({\rm H}_{2})$ from Table 3.3. The results are presented in Table
9. Note, in particular, the last two columns of Table 9, which give
$N_{1}({\rm H}_{\rm tot})=N_{1}({\rm H~{}\textsc{i}})+2N({\rm H}_{2})$ and
$f_{1}({\rm H}_{2})=2N({\rm H}_{2})/N_{1}({\rm H}_{\rm tot})$, the total
hydrogen column density and molecular hydrogen fraction, respectively, in the
optically-thick H2 region. We also include in Table 9 those sight lines with
Cl i and Cl ii measurements from Copernicus observations (Moomey et al., 2012;
Brown, 2015, Appendix A).

Table 9: Neutral Chlorine Fractions and Derived Parameters Star | log $N$(Htot) | log $f$(Cl i) | log $f$(H2) | log $N_{1}$(H i)aaColumn density of atomic hydrogen associated with the optically-thick H2 region. | log $N$(H2) | log $N_{1}$(Htot)bbTotal hydrogen column density associated with the optically-thick H2 region. | log $f_{1}$(H2)ccFraction of hydrogen in molecular form in the optically-thick H2 region.
---|---|---|---|---|---|---|---
HD 108 | $21.47^{+0.04}_{-0.05}$ | $-0.22\pm 0.05$ | $-0.72\pm 0.09$ | $21.09\pm 0.09$ | $20.45\pm 0.08$ | $21.26\pm 0.07$ | $-0.50\pm 0.10$
HD 3827 | $20.56^{+0.07}_{-0.08}$ | $-1.17\pm 0.11$ | $-1.90\pm 0.25$ | $19.29\pm 0.15$ | $18.35\pm 0.25$ | $19.38\pm 0.14$ | $-0.73\pm 0.27$
HD 12323 | $21.28^{+0.04}_{-0.04}$ | $-0.20\pm 0.07$ | $-0.72\pm 0.09$ | $20.93\pm 0.11$ | $20.26\pm 0.08$ | $21.08\pm 0.09$ | $-0.52\pm 0.11$
HD 15137 | $21.32^{+0.07}_{-0.08}$ | $-0.26\pm 0.08$ | $-0.79\pm 0.10$ | $20.90\pm 0.13$ | $20.23\pm 0.08$ | $21.06\pm 0.10$ | $-0.53\pm 0.12$
HD 24190 | $21.30^{+0.05}_{-0.06}$ | $-0.15\pm 0.05$ | $-0.62\pm 0.08$ | $20.96\pm 0.10$ | $20.38\pm 0.07$ | $21.15\pm 0.07$ | $-0.46\pm 0.10$
HD 24534 | $21.34^{+0.03}_{-0.04}$ | $-0.05\pm 0.03$ | $-0.12\pm 0.05$ | $20.51\pm 0.26$ | $20.92\pm 0.04$ | $21.30\pm 0.06$ | $-0.08\pm 0.07$
HD 37903 | $21.44^{+0.06}_{-0.07}$ | $-0.45\pm 0.11$ | $-0.29\pm 0.09$ | …ddThe derived value of $N_{1}$(H i) for this sight line is negative. However, the data are consistent with $f_{1}$(H2) = 1. | $20.85\pm 0.07$ | $20.99\pm 0.18$ | $+0.16\pm 0.19$
HD 41161 | $21.15^{+0.05}_{-0.05}$ | $-0.38\pm 0.07$ | $-0.86\pm 0.10$ | $20.60\pm 0.12$ | $19.99\pm 0.09$ | $20.78\pm 0.09$ | $-0.49\pm 0.12$
HD 46223 | $21.58^{+0.03}_{-0.04}$ | $-0.18\pm 0.04$ | $-0.61\pm 0.07$ | $21.21\pm 0.08$ | $20.67\pm 0.06$ | $21.41\pm 0.06$ | $-0.43\pm 0.08$
HD 52266 | $21.27^{+0.04}_{-0.04}$ | $-0.55\pm 0.15$ | $-0.97\pm 0.09$ | $20.51\pm 0.23$ | $20.00\pm 0.08$ | $20.72\pm 0.16$ | $-0.42\pm 0.17$
HD 53975 | $21.09^{+0.04}_{-0.04}$ | $-0.63\pm 0.16$ | $-1.64\pm 0.10$ | $20.41\pm 0.18$ | $19.15\pm 0.09$ | $20.46\pm 0.16$ | $-1.01\pm 0.18$
HD 63005 | $21.31^{+0.03}_{-0.03}$ | $-0.31\pm 0.09$ | $-0.84\pm 0.09$ | $20.84\pm 0.13$ | $20.17\pm 0.09$ | $21.00\pm 0.10$ | $-0.52\pm 0.13$
HD 66788 | $21.26^{+0.04}_{-0.04}$ | $-0.58\pm 0.09$ | $-1.24\pm 0.14$ | $20.56\pm 0.12$ | $19.72\pm 0.14$ | $20.67\pm 0.10$ | $-0.65\pm 0.17$
HD 73882 | $21.57^{+0.08}_{-0.09}$ | $-0.15\pm 0.07$ | $-0.19\pm 0.12$ | $20.28\pm 0.76$ | $21.08\pm 0.10$ | $21.41\pm 0.15$ | $-0.03\pm 0.17$
HD 75309 | $21.19^{+0.03}_{-0.03}$ | $-0.39\pm 0.07$ | $-0.73\pm 0.07$ | $20.53\pm 0.14$ | $20.16\pm 0.06$ | $20.80\pm 0.09$ | $-0.34\pm 0.10$
HD 88115 | $21.04^{+0.06}_{-0.07}$ | $-0.96\pm 0.05$ | $-1.49\pm 0.15$ | $19.94\pm 0.11$ | $19.25\pm 0.14$ | $20.09\pm 0.09$ | $-0.54\pm 0.16$
HD 89137 | $21.11^{+0.06}_{-0.07}$ | $-0.39\pm 0.08$ | $-0.79\pm 0.11$ | $20.49\pm 0.15$ | $20.02\pm 0.09$ | $20.72\pm 0.10$ | $-0.40\pm 0.13$
HD 90087 | $21.23^{+0.05}_{-0.05}$ | $-0.68\pm 0.05$ | $-1.05\pm 0.08$ | $20.31\pm 0.10$ | $19.88\pm 0.07$ | $20.55\pm 0.07$ | $-0.37\pm 0.10$
HD 91597 | $21.42^{+0.06}_{-0.07}$ | $-1.09\pm 0.05$ | $-1.42\pm 0.07$ | $20.05\pm 0.12$ | $19.70\pm 0.05$ | $20.33\pm 0.07$ | $-0.33\pm 0.09$
HD 91824 | $21.16^{+0.04}_{-0.04}$ | $-0.46\pm 0.07$ | $-1.05\pm 0.10$ | $20.57\pm 0.10$ | $19.81\pm 0.09$ | $20.70\pm 0.08$ | $-0.59\pm 0.12$
HD 91983 | $21.22^{+0.05}_{-0.06}$ | $-0.50\pm 0.10$ | $-0.82\pm 0.09$ | $20.44\pm 0.19$ | $20.10\pm 0.07$ | $20.72\pm 0.12$ | $-0.32\pm 0.13$
HD 92554 | $21.35^{+0.09}_{-0.11}$ | $-1.21\pm 0.07$ | $-1.81\pm 0.11$ | $20.01\pm 0.12$ | $19.24\pm 0.07$ | $20.14\pm 0.10$ | $-0.59\pm 0.12$
HD 94493 | $21.18^{+0.05}_{-0.06}$ | $-0.79\pm 0.05$ | $-0.79\pm 0.08$ | …ddThe derived value of $N_{1}$(H i) for this sight line is negative. However, the data are consistent with $f_{1}$(H2) = 1. | $20.09\pm 0.06$ | $20.38\pm 0.11$ | $+0.01\pm 0.12$
HD 97175 | $21.07^{+0.05}_{-0.06}$ | $-0.75\pm 0.12$ | $-0.66\pm 0.10$ | …ddThe derived value of $N_{1}$(H i) for this sight line is negative. However, the data are consistent with $f_{1}$(H2) = 1. | $20.10\pm 0.09$ | $20.31\pm 0.18$ | $+0.09\pm 0.20$
HD 99857 | $21.35^{+0.06}_{-0.07}$ | $-0.50\pm 0.08$ | $-0.75\pm 0.08$ | $20.51\pm 0.19$ | $20.30\pm 0.05$ | $20.86\pm 0.10$ | $-0.26\pm 0.11$
HD 99890 | $21.14^{+0.05}_{-0.05}$ | $-0.63\pm 0.06$ | $-1.28\pm 0.10$ | $20.41\pm 0.09$ | $19.56\pm 0.09$ | $20.52\pm 0.07$ | $-0.66\pm 0.11$
HD 109399 | $21.17^{+0.05}_{-0.05}$ | $-0.66\pm 0.07$ | $-0.86\pm 0.10$ | $20.09\pm 0.22$ | $20.01\pm 0.09$ | $20.52\pm 0.11$ | $-0.21\pm 0.14$
HD 116852 | $21.01^{+0.04}_{-0.04}$ | $-0.64\pm 0.07$ | $-0.96\pm 0.10$ | $20.09\pm 0.15$ | $19.75\pm 0.09$ | $20.37\pm 0.09$ | $-0.32\pm 0.12$
HD 121968 | $20.59^{+0.12}_{-0.16}$ | $-1.41\pm 0.07$ | $-1.59\pm 0.15$ | $18.72\pm 0.29$ | $18.70\pm 0.10$ | $19.19\pm 0.14$ | $-0.18\pm 0.16$