# Low-ionization iron-rich Broad Absorption-Line Quasar SDSS J 1652$+$2650:
Physical conditions in the ejected gas from excited Fe ii and metastable He
i††thanks: Based on observations carried out in visitor mode at the European
Southern Observatory under ESO program ID 0103.A-0529(A) (PI: Ledoux).

S. A. Balashev,1, C. Ledoux,2 P. Noterdaeme,3,4 P. Boissé,4 J.-K. Krogager,5S.
López6 and K. N. Telikova1
1 Ioffe Institute, Polyteknicheskaya 26, 194021 Saint-Petersburg, Russia
2 European Southern Observatory, Alonso de Córdova 3107, Vitacura, Casilla
19001, Santiago, Chile
3 Franco-Chilean Laboratory for Astronomy, IRL 3386, CNRS and U. de Chile,
Casilla 36-D, Santiago, Chile
4 Institut d’Astrophysique de Paris, CNRS-SU, UMR 7095, 98bis bd Arago, 75014
Paris, France
5 Centre de Recherche Astrophysique de Lyon, UMR 5574, 9 avenue Charles André,
69230 Saint-Genis-Laval, France
6 Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago,
Chile E-mail<EMAIL_ADDRESS>
(Accepted 2023. Received 2023; in original form 2023)

###### Abstract

We present high-resolution VLT/UVES spectroscopy and a detailed analysis of
the unique Broad Absorption-Line system towards the quasar SDSS J
165252.67$+$265001.96. This system exhibits low-ionization metal absorption
lines from the ground states and excited energy levels of Fe ii and Mn ii, and
the meta-stable $2\,^{3}S$ excited state of He i. The extended kinematics of
the absorber encompasses three main clumps with velocity offsets of $-5680$,
$-4550$, and $-1770$ km s-1 from the quasar emission redshift, $z=0.3509\pm
0.0003$, derived from [O ii] emission. Each clump shows moderate partial
covering of the background continuum source, $C_{f}\approx[0.53;0.24;0.81]$.
We discuss the excitation mechanisms at play in the gas, which we use to
constrain the distance of the clouds from the Active Galactic Nucleus (AGN) as
well as the density, temperature, and typical sizes of the clouds. The number
density is found to be $n_{\rm H}\sim 10^{4}$ cm-3 and the temperature $T_{\rm
e}\sim 10^{4}$ K, with longitudinal cloudlet sizes of $\gtrsim 0.01$ pc.
Cloudy photo-ionization modelling of He i⋆, which is also produced at the
interface between the neutral and ionized phases, assuming the number
densities derived from Fe ii, constrains the ionization parameter to be $\log
U\sim-3$. This corresponds to distances of a few 100 pc from the AGN. We
discuss these results in the more general context of associated absorption-
line systems and propose a connection between FeLoBALs and the recently-
identified molecular-rich intrinsic absorbers. Studies of significant samples
of FeLoBALs, even though rare per se, will soon be possible thanks to large
dedicated surveys paired with high-resolution spectroscopic follow-ups.

###### keywords:

quasars: absorption lines; quasars: individual: SDSS J 165252.67$+$265001.96,
NVSS J 235953$-$124148; line: formation; galaxies: active.

††pubyear: 2023††pagerange: Low-ionization iron-rich Broad Absorption-Line
Quasar SDSS J 1652$+$2650: Physical conditions in the ejected gas from excited
Fe ii and metastable He i††thanks: Based on observations carried out in
visitor mode at the European Southern Observatory under ESO program ID
0103.A-0529(A) (PI: Ledoux).–A

##  1 Introduction

A key characteristic of around 20% of optically-selected quasars is the
occurrence of broad absorption-line (BAL) systems along the line-of-sight to
the quasar (Tolea et al. 2002; Hewett & Foltz 2003; Reichard et al. 2003;
Knigge et al. 2008; Gibson et al. 2009). BAL systems are typically associated
with highly-ionized metals, e.g., C iv and O vi, and their wide kinematic
spreads, velocity offsets, and partial covering factors all indicate that they
are produced by out-flowing gas. Observations of such outflows provide a
direct test of quasar feedback models.

One-tenth of BAL systems show associated wide Mg ii absorption (Trump et al.,
2006) and are called low-ionization BALs (hereafter LoBALs). An even smaller
fraction, totalling only $\sim 0.3\%$ of the global quasar population, in
addition, exhibits Fe ii absorption and is hence called FeLoBALs. A qualifying
feature of FeLoBALs is the detection of Fe ii in its various fine-structure
energy levels of the lowest electronic states. These levels may be excited by
collisions or UV pumping, and their relative abundance can provide robust
estimates of critical physical parameters. Interestingly, the modelling of
FeLoBALs indicates they contain some neutral gas and likely occur at the
interface between the ionized and neutral media (Korista et al., 2008).
Another feature of FeLoBALs, which was gradually recognised, is the presence
of absorption lines corresponding to transitions from the first excited level
of neutral helium, He i∗ (Arav et al., 2001; Aoki et al., 2011; Leighly et
al., 2011). This is observed in FeLoBALs but also more generally in LoBALs
(Liu et al., 2015). These lines have also been detected in the host galaxies
of a few GRBs (Fynbo et al., 2014). He i∗ is predominately populated by
recombination of He ii and the measured column densities of He i∗ provide a
measure of the total column density of the ionized medium. This can constrain
the physical conditions in the outflowing gas and determine the total mass
budget to draw a more complete physical picture of quasar activity. For
example, rapid cooling followed by the phase transition and subsequent
condensation in an outflowing medium can result in the escape of small chunks
of the medium from the outflowing gas. Such cloudlets can precipitate back
onto the central engine and sustain the formation of the broad-line region
around the central powering source (Elvis, 2017).

Because the incidence rate of FeLoBALs in quasars is low, only a small sample
of such systems was found in the Sloan Digital Sky Survey (SDSS) database
(e.g., Trump et al., 2006; Farrah et al., 2012; Choi et al., 2022). Most
importantly, only about a dozen such systems was studied so far by means of
high-resolution near-UV and visual spectroscopy, i.e.: Q 0059$-$2735 (Hazard
et al., 1987; Wampler et al., 1995; Xu et al., 2021), Q 2359$-$1241 (Arav et
al., 2001, 2008), FIRST 104459.6$+$365605 (Becker et al., 2000; de Kool et
al., 2001), FBQS 0840$+$3633 (Becker et al., 1997; de Kool et al., 2002b),
FIRST J 121442.3$+$280329 (Becker et al., 2000; de Kool et al., 2002a), SDSS J
030000.56$+$004828.0 (Hall et al., 2003), SDSS J 0318$-$0600 (Dunn et al.,
2010; Bautista et al., 2010), AKARI J 1757$+$5907 (Aoki et al., 2011), PG
1411$+$442 (Hamann et al., 2019), SDSS J 2357$-$0048 (Byun et al., 2022b),
SDSS J 1439$-$0106 (Byun et al., 2022a), SDSS J 0242$+$0049 (Byun et al.,
2022c), SDSS J 1130$+$0411 (Walker et al., 2022), Mrk 231 (Boroson et al.,
1991; Smith et al., 1995; Veilleux et al., 2016), and NGC 4151 (Crenshaw et
al., 2000; Kraemer et al., 2001). Among these, only a few systems exhibit mild
line saturation and overlapping, which allow one to resolve the fine-structure
lines and therefore derive robust constraints on the gas physical conditions
(Arav et al., 2008). Moreover, each previously-studied system appears to be
fairly specific, i.e., FeLoBALs show a broad range of properties, which means
any new observation and detailed analysis potentially bring new valuable clues
to understanding the physics and environmental properties of AGN outflows.

In this paper, we report the serendipitous discovery of a multi-clump FeLoBAL
towards SDSS J 165252.67$+$265001.96, which we refer to in the following as J
1652$+$2650\. We present high-quality VLT/UVES data of this quasar and the
spectroscopic analysis of the absorption system and discuss the excitation
mechanisms at play in the gas. Our goal is to infer the physical properties of
FeLoBAL clouds and estimate their distance from the central engine.

##  2 Observations and data reduction

We selected the bright quasar J 1652$+$2650 ($B=18.2$; $V=17.7$; $z_{\rm
em}=0.35$; Véron-Cetty & Véron, 2010) with the primary goal to search for CN,
CH, and CH+ molecules in absorption based on the detection of strong
associated Na i lines at $z\approx 0.33$ in the SDSS spectrum (Pâris et al.,
2018; Negrete et al., 2018), which is shown in Fig. 1. We observed the target
in visitor mode on the night of July 27, 2019, with UVES, the Ultraviolet and
Visual Echelle Spectrograph (Dekker et al., 2000) installed at the Nasmyth-B
focus of the ESO Very Large Telescope Unit-2, Kueyen. The total on-source
integration time was four hours, subdivided evenly into three exposures taken
in a row. The instrumental setup used Dichroic beam splitter #1 and
positioning of the cross-dispersers at central wavelengths of 390 nm and 590
nm in the Blue and Red spectroscopic arms, respectively. In each arm, the slit
widths were fixed to $1\arcsec$ and CCD pixels were binned $2\times 2$. While
observing, the weather conditions were excellent with clear sky transparency
and a measured Differential Image Motion Monitor (Sarazin & Roddier, 1990)
seeing of $0\aas@@fstack{\prime\prime}6$. Despite a relatively high airmass
(1.63–1.97), the source was recorded on the detectors with a spatial PSF trace
of only $1\aas@@fstack{\prime\prime}1$ FWHM in the Blue
($1\aas@@fstack{\prime\prime}0$ FWHM in the Red).

The raw data from the telescope was reduced offline applying the recipes of
the UVES pipeline v5.10.4 running on the ESO Reflex platform. During this
process, the spectral format of the data was compared to a physical model of
the instrument, to which a slight CCD rotation was applied ($-0.05\degr$ in
the Blue; $+0.05\degr$ in the Red). ThAr reference frames acquired in the
morning following the observations were used to derive wavelength-calibration
solutions, which showed residuals of 1.53 mÅ RMS in the Blue (4.25 mÅ RMS in
the Red). The object and sky spectra were extracted simultaneously and
optimally, and cosmic-ray hits were removed efficiently using a
$\kappa$-$\sigma$ clipping factor of five. The wavelength scale was converted
to the helio-vacuum rest frame. Individual 1D exposures were then scaled and
combined together by weighing each pixel by its S/N. The S/N of the final
science product is $\sim 15$ per pixel at $325<\lambda_{\rm obs}<455$ nm and
$\sim 32$ per pixel at $490<\lambda_{\rm obs}<690$ nm. With a delivered
resolving power of 50 000, the instrumental line-spread function is 6 km s-1
FWHM.

##  3 Data analysis

###  3.1 Quasar spectrum and systemic redshift

J 1652$+$2650 exhibits moderate reddening. Based on the SDSS spectrum, shown
in Fig. 1, and using the Type I quasar template from Selsing et al. (2016), we
followed a procedure similar to that employed by Balashev et al. (2017, 2019)
and derived that $A_{V}\approx 1.2$, assuming standard galactic extinction law
(Fitzpatrick & Massa, 2007). This is quite large compared to intervening
quasar absorbers, e.g., DLAs (Murphy & Bernet, 2016). We also note that this
quasar shows iron-emission line complexes in the spectral regions around 4600
Å and 5300 Å in the quasar rest frame, that are enhanced by a factor of $\sim
4$ relative to the fiducial quasar template (see Fig. 1).

Figure 1: SDSS spectrum of J 1652$+$2650\. The top panel shows the full range
of SDSS spectrum plotted by black line. The quasar emission lines are marked
by pink arrows, while pink line represents the composite QSO spectrum reddened
by dust (see Sect. 3.1). The bottom panels show a zoom view of Ca ii and Na i
absorption line regions (these lines were used to preselect J 1652$+$2650 from
SDSS database, see Sect. 2). The red, green and blue lines indicate a
absorption line complexes, that were resolved using UVES spectrum, see Sect.
3.2.

To determine the quasar emission redshift accurately, we followed the
recommendations of Shen et al. (2016) that Ca ii and [O ii] should be
considered the most reliable systemic-redshift indicators. In the case of J
1652$+$2650, the blue side of the Ca ii profile is affected by strong self-
absorption, so we are left with [O ii], which according to Shen et al. (2016)
is not significantly shifted relative to Ca ii. Based on a single-component
Gaussian fit, we measured $z_{\rm em}=0.3511(7)$ when considering the [O
ii],$\lambda$3727.092 transition line alone, and $z_{\rm em}=0.3506(7)$ when
using the mean wavelength of the [O ii],$\lambda\lambda$3727.092,3729.875
doublet. This translates into $z_{\rm em}=0.3509\pm 0.0003$, which we consider
as our most-accurate determination of the quasar systemic redshift. The
H$\beta$ emission line is observed at a redshift of $z\approx 0.3494$,
implying a velocity blue-shift of $\Delta V\sim-330$ km s-1 relative to [O
ii]. This is consistent with the findings of Shen et al. (2016) using their
own sample.

###  3.2 Absorption-line system overview

The FeLoBAL111Based on Mg ii lines this system do not satisfy standard BAL
definition Weymann et al. (1981), and should be attributed to the mini-BAL.
However, C iv lines (that are typically used in BAL definition and usually
indicate much wider profiles than Mg ii lines) out of the spectral range.
Therefore, we will keep denotation of this system as FeLoBAL, which is
supported by example of Q 2359$-$1241(Arav et al., 2001), for which Mg ii
lines indicate a similar width as in J 1652$+$2650, but HST observations
confirm large width of C iv lines. on the line-of-sight to J 1652$+$2650
consists of multiple prominent absorption lines from Mg ii, Ca ii, He i⋆
(i.e., the meta-stable excited state $2^{3}\rm S$), Mg i, Fe ii, and Mn ii,
all covered by the UVES spectrum. The system is composed of three main,
kinematically-detached absorption-line complexes,222When selecting this quasar
to search for intervening molecular absorption, we assumed the line-of-sight
could intersect three galaxies (possibly located in a cluster hosting the
quasar). It turns out that the gas is associated with the quasar active
nucleus itself in spite of its low ionization. i.e., at $z_{\rm abs}=0.32531$
($\Delta V\approx-5680$ km s-1), 0.33043 ($\Delta V\approx-4550$ km s-1), and
0.34292 ($\Delta V\approx-1770$ km s-1), where the reddest and bluest clumps
exhibit the strongest Mg ii and Ca ii absorption overall (see Fig. 2). In the
following, we refer to these three complexes as $A$, $B$, and $C$, in order of
increasing redshift. Each complex has at least a few velocity components
resolved by eye within its own profile. Weak absorption is also visible in Mg
ii at $z_{\rm abs}=0.3357$ ($\Delta V\approx-3390$ km s-1) and, tentatively,
also in Ca ii $\lambda$3934.

Figure 2: Portions of the normalized UVES spectrum showing the kinematics of
Mg ii (upper panel), Ca ii H and K, and He i⋆ $\lambda 3889$ (lower panel) in
the FeLoBAL towards J 1652$+$2650\. The Mg ii absorption-line complex at
$z\approx 0.3357$ is much weaker than complexes $A$, $B$, and $C$, and
therefore is not included in the following Voigt-profile fitting analysis. In
both panels, the top axis shows the velocity of the strongest transition,
i.e., Mg ii$\lambda 2796$ (upper panel) or Ca ii$\lambda 3934$ (lower panel),
relative to the quasar systemic redshift.

Fe ii,$\lambda\lambda$2586,2600 ground-state absorption lines as well as lines
from the fine-structure energy levels of the two LS states $\rm
3d^{6}4s\,a^{6}D$ (ground) and $\rm 3d^{6}4s\,a^{4}D$ (second excited state,
which is encompassing the Fe ii9∗–Fe ii12∗ levels) are detected in this system
(see Figs. 4, 12, 13, and 14). Transition lines from the first excited LS
state $\rm 3d^{7}\,a^{4}F$ (i.e., 5th to 8th excited levels above the ground
state) are not covered by our spectrum, being located bluewards of the
observed wavelength range. The Mn ii $\lambda\lambda\lambda$2576,2594,2606
triplet and Mn ii⋆ transition lines (i.e.,
$\lambda\lambda\lambda$2933,2940,2950) from the first excited level of Mn ii,
with an excitation energy of 9473 cm-1, are detected most clearly in the $C$
(reddest) and $A$ (bluest) clumps (see Fig. 15). Such transitions were
detected before only in a few FeLoBALs (e.g., FBQS J 1151$+$3822; Lucy et al.,
2014).

Ca i $\lambda$4227 and CH+ $\lambda\lambda$3958,4233 absorptions are not
detected. Using the oscillator strengths of CH+ lines from Weselak et al.
(2009), we derive a 2$\sigma$ upper limit on the column density of $N(\rm
CH^{+})=10^{13}\rm\,cm^{-2}$ for each of the three complexes. We detect
possible Na i emission lines in both the UVES and SDSS spectra. All absorption
lines in the UVES spectrum were identified, except a weak line at
$\lambda_{\rm obs}=5341$ Å, which has a similarly wide profile as the FeLoBAL
lines. Searching an identification in the NIST database at the different BAL
sub-redshifts did not provide any satisfactory solution, therefore it has
likely spurious nature.

###  3.3 Evidence of partial flux covering from Mg ii

Figure 3: Top to Bottom: Mg ii in absorption-line complexes $A$, $B$, and $C$,
at $z_{\rm abs}=0.32531$, 0.33043, and 0.34292, respectively. Upper insets in
each panel: Apparent optical-depth ratio of the Mg ii doublet. The regions
where the lines are not apparently blended are highlighted in colour (blue,
green, or red). The dashed and dotted horizontal lines correspond to the ratio
expected in the cases of complete line saturation ($\tau_{1}/\tau_{2}=1$) and
optically-thin lines ($\tau_{1}/\tau_{2}=2$), respectively.

Singly-ionized magnesium exhibits an overall similar velocity structure as Ca
ii, but much stronger lines, hence additional components are detected in the
Mg ii profiles. Since the Mg ii lines are saturated (see Fig. 2), it would not
be reasonable to fit such complex profiles as the velocity decomposition would
be unreliable and solutions degenerate. However, we can obtain qualitative
results by comparing the apparent optical-depth ratios of the Mg ii doublet
lines, $\tau_{1}/\tau_{2}$ (where $\tau_{1}$ and $\tau_{2}$ are the apparent
optical depths in Mg ii $\lambda 2796$ and Mg ii $\lambda 2803$,
respectively). This is shown in Fig. 3 for all three absorption-lines
complexes, $A$, $B$, and $C$. When the absorber fully covers the continuum
source, the apparent optical-depth ratio of the Mg ii doublet is expected to
be $\tau_{1}/\tau_{2}=f_{1}\lambda_{1}/f_{2}\lambda_{2}\approx 2$ (where
$f_{i}$ and $\lambda_{i}$ are the line oscillator strengths and wavelengths,
respectively), unless the lines are not fully saturated. One can see in Fig. 3
that in our case $\tau_{1}/\tau_{2}$ is close to unity along the entire
profiles, which, together with seemingly saturated profiles, is evidence for
partial flux covering. In the case of fully-saturated line profiles, partial
covering factors ($C_{f}$) can be roughly determined as $C_{f}\approx
1-e^{-\tau_{1}}$. Therefore, a value of $\tau_{1}/\tau_{2}$ close to unity
even in the line wings (where the optical depths $\tau_{1,2}<1$) indicates
that partial covering is likely changing through the profile, which may
additionally complicate line-profile fitting. Using the flux residuals
observed at the bottom of the profiles (where $\tau_{1}/\tau_{2}\approx 1$),
we derive upper limits on $C_{f}$ of $\sim 0.82$, 0.68, and 0.96, in complexes
$A$, $B$, and $C$, respectively. We also note that the Mg ii $\lambda\lambda
2796,2803$ lines are not blended with each other, with the exception of the
far wings of Mg ii $\lambda 2803$ and Mg ii $\lambda 2796$ in complexes $A$
and $B$, respectively. In addition, the velocity differences between complexes
$A$, $B$, and $C$, and the weaker complex at $z\approx 0.3357$, do not
correspond to any of the strong high-ionization line-doublet splitting, i.e.,
Si iv, C iv, N v, nor O vi. Therefore, line locking (see, e.g., Bowler et al.,
2014) is not clearly present in this system.

###  3.4 Voigt-profile fitting of Ca ii, Mg i, He i⋆, Fe ii, and Mn ii

We performed simultaneous fits to Ca ii, Mg i, He i⋆, Fe ii and Mn ii
absorption lines using multiple-component Voigt profiles. While even Ca ii
$\lambda\lambda 3934,3969$, Mg i $\lambda 2852$, and He i⋆
($\lambda\lambda\lambda 2945,3188,3889$) lines are located in spectral regions
of high S/N, the weakness of some of the velocity components prevents us from
fitting the lines individually. Additionally, Fe ii and Mn ii lines are
significantly blended with each other, not only between components of a given
complex ($A$, $B$, or $C$) but also between components pertaining to different
complexes. Therefore, to obtain internally consistent fits, we tied the
Doppler parameters in each component assuming them to be equal for each
species. This implicitly assumes that turbulent broadening dominates over
thermal broadening (micro-turbulence assumption), which is reasonable for the
wide (FWHM $>15$ km s-1) profiles of this system.

As we mentioned, absorption lines from Fe ii and Mn ii in the UVES spectrum
display a high degree of mutual blending and complexity, therefore, in order
to remove possible degeneracies, we assumed that the Fe ii levels are
populated by collision with electrons (as argued for the majority of
previously-studied FeLoBALs (Korista et al., 2008; Dunn et al., 2010; Bautista
et al., 2010; Byun et al., 2022a). This assumption also minimizes the number
of independent variables in the analysis. Thus, for each velocity component,
the column densities of Fe ii levels are set by the total Fe ii column
density, the electron density, and the temperature. The data for the strengths
of collisions with electrons were taken from the CHIANTI 9.0.1 database (Dere
et al., 2019) and the atomic data from the NIST database. We did not find any
data for the collisional excitation of Mn ii levels. Therefore, we could not
consider its excitation together with Fe ii and hence we derived the column
densities of the Mn ii and Mn ii⋆ levels, independently from Fe ii. The atomic
data for He i⋆ and Ca ii were taken from Drake & Morton (2007) and Safronova &
Safronova (2011), respectively. For lines from excited levels of Fe ii and Mn
ii, we used the data from Nave & Johansson (2013), Schnabel et al. (2004), and
Kling & Griesmann (2000), respectively. For other species, we used the atomic
data compiled by Morton (2003).

Similar to the analysis of Mg ii lines in Sect. 3.3, the apparent optical
depth of the Ca ii (as well as Fe ii) lines indicates partial covering.
Therefore, we need to include partial covering in the line-profile fitting
procedure and for this we used the simple model proposed by Barlow & Sargent
(1997; see also Balashev et al. 2011). Within this model, it is assumed that
the velocity components with non-unity covering factors spatially overlap. If
it were not the case, this would introduce additional covering factors to
describe mutual overlapping (see, e.g., the discussion in Ishita et al. 2021).
When many components are intertwined in wavelength space, this requires a
significant increase in the number of independent variables for the analysis
(up to the factorial of $n$, where $n$ is the total number of components).
This complicates the analysis and makes the derived results ambiguous.
Therefore, we made no attempt here to include mutual covering in the fitting
procedure but rather tied the covering factors of all the components within
the same complex (i.e., $A$, $B$, or $C$) to be the same. The model employed
here therefore can only provide coarse estimates of the covering factors.

The likelihood function was constructed using visually-identified regions of
the spectrum that are associated with the lines to be fitted assuming a normal
distribution of the pixel uncertainties. To obtain the posterior probability
functions on the fit parameters (i.e., column densities, redshifts, Doppler
parameters, and covering factors), we used a Bayesian approach with affine-
invariant sampler (Goodman & Weare, 2010). We used flat priors on redshifts,
Doppler parameters, covering factors, and logarithms of column densities. For
the electron temperature (which is relevant for lines from the excited Fe ii
levels), we used a Gaussian prior of $\log T_{e}=4.2\pm 0.5$, corresponding to
the typical electron temperatures of a fully-ionized medium, where excited Fe
ii levels are highly populated (e.g., Korista et al., 2008). The sampling was
performed on the cluster running $\approx 100$ processes in parallel (using
several hundred walkers) which typically took a few days until convergence.
While this approach allows us to constrain the full shape of the posterior
distribution function for each parameter, in the following we report the fit
results in a standard way. The point and interval estimates correspond to the
maximum posterior probability and the 0.683 credible intervals obtained from
1D marginalized posterior distribution functions.

Figure 4: Left to Right: Voigt-profile fits to selected Ca ii, He i⋆, Mg i, Mn
ii, and Fe ii absorption lines in complexes $A$ (Left), $B$ (Middle), and $C$
(Right), at $z_{\rm abs}=0.3253$, 0.3304, and 0.3429, respectively, towards J
1652$+$2650\. The coloured stripes show a 0.683 credible interval of the line
profiles sampled from posterior probability distributions of fitting
parameters. The yellow represents the total line profile, while the blue,
green, and red lines indicate individual components from complexes $A$, $B$,
and $C$, respectively. The vertical lines show the positions of each
component. Horizontal dashed lines and their surrounding grey areas indicate
the extent of partial covering determined by fitting each clump independently
with its own covering factor. The spectrum was rebinned to 0.1 Å scale for
presentation purposes. Note the different y-axis scaling in each column. The
original spectrum and all absorption-line profiles are displayed in Figs. 11,
12, 13, 14, and 15.

The results of Voigt-profile fitting are given in Table 1 and the modeled line
profiles are briefly shown in Fig. 4, and fully displayed in Figs. 11, 12, 13,
14, and 15. The derived Doppler parameters span a wide range, from several up
to two hundred km s-1. The largest Doppler parameters found here should
however be considered as upper limits only due to our inability to
unambiguously resolve the line profiles in individual components owing to
insufficient spectrum quality and few available transitions. While the column-
density ratios between components are not drastically varied some trends
likely appear. For example, the components in complex $A$ have a much larger
Ca ii-to-Mg i column-density ratio than the components in complexes $B$ or
$C$. This likely indicates that the physical conditions vary from one complex
to the other. Alternatively, this may indicate that the species under study
are not co-spatial, which would weaken the assumption of identical Doppler
parameters and location in velocity space for the considered species.
Therefore, the derived uncertainties on the column densities should be
considered with caution as they only describe uncertainties in a statistical
sense. In complexes $A$, $B$, and $C$, the covering factors, $C_{f}$, are
measured to be $0.53^{+0.01}_{-0.01}$, $0.24^{+0.01}_{-0.01}$, and
$0.81^{+0.01}_{-0.01}$, respectively, where again the quoted uncertainties
should be considered with caution given the assumptions discussed above. The
covering factors are mainly constrained by the Ca ii and Fe ii lines since the
He i⋆ and Mn ii lines are weak (and hence are less sensitive to partial
covering) and Mg i exhibits a single line. Therefore, the Mg i column
densities reported in Table 1 are only reliable if the Mg i-bearing gas is co-
spatial with Ca ii. One should also note that the covering factors derived
here are smaller than those found for Mg ii in Sect. 3.3. This indicates that
the spatial extent of the Mg ii-bearing gas is larger than that of Ca ii and
Fe ii.

Figure 5: Comparison of physical parameters in velocity complexes $A$, $B$,
and $C$ (in blue, green, and red, respectively) towards J 1652$+$2650\. The
values are taken from Table 1, except $\rm[Fe/H]$, which is the Fe gas-phase
abundance relative to the solar. For presentation purposes, we only show the
values that are reasonably well-constrained. In the lower panel, larger values
of $C_{f}$ correspond to larger covering factors.

In Fig. 5, we plot the physical parameters derived using the lines from
excited levels of Fe ii for velocity complexes $A$, $B$, and $C$. We found
that the electron densities in different components lie in the range
$n_{e}\approx[10^{2};10^{5}]$ cm-3, which is expected given the detection of
highly-excited levels (with energies $\lesssim 10^{4}$ cm-1). Interestingly,
the electron densities in complex $A$ are found to be systematically higher
than in $B$ or $C$. This suggests that complex $A$ is located closer to the
central engine in a harsher environment, which in turn is in agreement with
the assumption that these complexes are produced in a decelerated-wind medium.
We however note, that the line profiles are quite complex and therefore exact
velocity decomposition is quite complicated in this system and hence our
solution is not necessary to be unique.

The electron temperature is found to be in the wide range $T_{\rm e}\approx
10^{3.5}-10^{4.5}$ K, close to the observationally motivated chosen priors.
This is not surprising since collisional population is less sensitive to the
temperature than to the number density itself. Using the inferred electron
densities (representing the number density, since the excited Fe ii originates
from ionized gas (see, e.g., Korista et al., 2008), and the total column
density of Fe ii one can derive the longitudinal extent of the absorbing
clouds associated with each component, which we found to also span a wide
range. When inferring this value, one should take into account the Fe gas
abundances. If we assume a solar Fe ii abundance, we get characteristic values
of about $10^{15}$ and $10^{15.5}\,\rm cm$ for complexes $A$ and $B$,
respectively, and a wide range of values for complex $C$. These values should
be considered as lower limits only, since neither the metallicity nor the Fe
depletion, nor the Fe ii ionization correction, are known. Indeed, the
depletion of Fe is much less than $<1$ and a very sensitive function of
metallicity, and Fe ii can be a subdominant form of Fe, even where excited Fe
ii levels are populated (see, e.g., Korista et al., 2008). The ratio of Mn ii⋆
to Mn ii column densities was found to be around $0.1-0.3$, similar to Fe ii,
where complex $A$ exhibits slightly higher excitation than complex $C$ (Mn ii⋆
is very unconstrained in complex $B$). If fitted individually, the covering
factors derived here for each complex were found to be consistent between Ca
ii and Fe ii, which indicates that Fe ii and Ca ii-bearing clouds likely have
similar spatial extents.

Table 1: Results of simultaneous fits (shown in Figs. 4, 11, 12, 13, 14, and
15) to absorption lines from Fe ii, Ca ii, He i⋆, Mg i, and Mn ii in the
FeLoBAL towards J 1652$+$2650\.

Comp. | z |  $\Delta$v† | b | $\log n$ | $\log T$ |  $\log N_{\rm tot}$(Fe ii) |  $\log N$(CaII) |  $\log N$(HeI∗) |  $\log N$(MgI) |  $\log N$(MnII) |  $\log N$(MnII∗) |  $C_{f}$‡
---|---|---|---|---|---|---|---|---|---|---|---|---
|  | [km s-1] | [km s-1] | [cm-3] | [K] |  |  |  |  |  |  |
$A,1$ | $0.324862(^{+23}_{-15})$ | -5778 | $46^{+7}_{-5}$ | $<2.9$ | $>4.0$ | $13.44^{+0.23}_{-0.24}$ | $<10.6$ | $13.48^{+0.08}_{-0.04}$ | $<11.1$ | $<11.3$ | $<11.9$ | $0.531^{+0.012}_{-0.008}$
$A,2$ | $0.325195(^{+14}_{-20})$ | -5704 | $44.3^{+1.8}_{-2.4}$ | $4.28^{+0.11}_{-0.14}$ | $3.9^{+0.1}_{-0.1}$ | $14.60^{+0.05}_{-0.07}$ | $12.90^{+0.05}_{-0.08}$ | $13.90^{+0.05}_{-0.05}$ | $12.39^{+0.07}_{-0.05}$ | $13.26^{+0.07}_{-0.08}$ | $12.33^{+0.13}_{-0.14}$ | "
$A,3$ | $0.3253310(^{+30}_{-30})$ | -5674 | $29.0^{+0.8}_{-1.2}$ | $>4.8$ | $3.7^{+0.1}_{-0.3}$ | $14.27^{+0.10}_{-0.12}$ | $13.18^{+0.04}_{-0.04}$ | $13.59^{+0.10}_{-0.09}$ | $12.29^{+0.06}_{-0.10}$ | $13.04^{+0.10}_{-0.18}$ | $12.57^{+0.06}_{-0.06}$ | "
$A,4$ | $0.325432(^{+17}_{-17})$ | -5651 | $181^{+3}_{-7}$ | $4.55^{+0.08}_{-0.07}$ | $4.4^{+0.1}_{-0.2}$ | $14.69^{+0.05}_{-0.03}$ | $13.154^{+0.019}_{-0.011}$ | $13.59^{+0.04}_{-0.04}$ | $12.57^{+0.04}_{-0.04}$ | $13.23^{+0.12}_{-0.14}$ | $12.56^{+0.10}_{-0.21}$ | "
$B,5$ | $0.329670(^{+30}_{-40})$ | -4710 | $125^{+14}_{-9}$ | $3.7^{+0.3}_{-0.3}$ | $3.9^{+0.3}_{-0.2}$ | $14.91^{+0.10}_{-0.15}$ | $12.97^{+0.04}_{-0.06}$ | $13.55^{+0.07}_{-0.07}$ | $12.83^{+0.05}_{-0.06}$ | $13.29^{+0.10}_{-0.19}$ | $12.89^{+0.10}_{-0.12}$ | $0.235^{+0.011}_{-0.012}$
$B,6$ | $0.330355(^{+18}_{-21})$ | -4559 | $27^{+4}_{-4}$ | $<2.1$ | $>3.9$ | $15.5^{+0.4}_{-0.5}$ | $12.72^{+0.11}_{-0.15}$ | $13.41^{+0.11}_{-0.16}$ | $<12.1$ | $<12.9$ | $12.36^{+0.17}_{-0.23}$ | "
$B,7$ | $0.330433(^{+12}_{-13})$ | -4542 | $15^{+4}_{-4}$ | $<4.6$ | $4.1^{+0.6}_{-0.3}$ | $<13.2$ | $12.39^{+0.18}_{-0.26}$ | $12.9^{+0.3}_{-0.3}$ | $<12.4$ | $12.97^{+0.24}_{-0.32}$ | $<11.6$ | "
$B,8$ | $0.330513(^{+13}_{-17})$ | -4524 | $72^{+4}_{-3}$ | $4.0^{+0.1}_{-0.2}$ | $3.92^{+0.15}_{-0.03}$ | $15.09^{+0.10}_{-0.09}$ | $13.23^{+0.04}_{-0.05}$ | $13.98^{+0.03}_{-0.07}$ | $12.97^{+0.03}_{-0.08}$ | $13.27^{+0.15}_{-0.20}$ | $12.63^{+0.15}_{-0.25}$ | "
$C,9$ | $0.341623(^{+30}_{-30})$ | -2059 | $140^{+17}_{-29}$ | $4.8^{+0.6}_{-0.3}$ | $3.9^{+0.2}_{-0.1}$ | $14.09^{+0.05}_{-0.07}$ | $12.30^{+0.05}_{-0.08}$ | $<12.7$ | $12.07^{+0.08}_{-0.06}$ | $<12.5$ | $12.35^{+0.09}_{-0.13}$ | $0.807^{+0.010}_{-0.014}$
$C,10$ | $0.342551(^{+32}_{-28})$ | -1853 | $62^{+6}_{-7}$ | $3.14^{+0.21}_{-0.17}$ | $3.8^{+0.1}_{-0.1}$ | $14.42^{+0.08}_{-0.06}$ | $12.56^{+0.06}_{-0.06}$ | $13.35^{+0.05}_{-0.07}$ | $12.42^{+0.05}_{-0.08}$ | $12.94^{+0.07}_{-0.10}$ | $<12$ | "
$C,11$ | $0.342705(^{+6}_{-15})$ | -1819 | $20.2^{+1.5}_{-3.2}$ | $3.40^{+0.16}_{-0.20}$ | $3.9^{+0.2}_{-0.1}$ | $14.15^{+0.16}_{-0.11}$ | $12.29^{+0.05}_{-0.14}$ | $12.52^{+0.20}_{-0.18}$ | $12.15^{+0.06}_{-0.14}$ | $12.27^{+0.24}_{-0.38}$ | $<11$ | "
$C,12$ | $0.342821(^{+25}_{-15})$ | -1793 | $14^{+11}_{-3}$ | $3.8^{+1.1}_{-0.9}$ | $3.7^{+0.3}_{-0.3}$ | $13.6^{+0.5}_{-0.4}$ | $12.05^{+0.22}_{-0.21}$ | $<12.3$ | $11.93^{+0.25}_{-0.21}$ | $12.59^{+0.28}_{-0.16}$ | $11.84^{+0.14}_{-0.22}$ | "
$C,13$ | $0.3428840(^{+30}_{-30})$ | -1779 | $8.9^{+1.3}_{-1.2}$ | $3.7^{+0.5}_{-1.0}$ | $3.7^{+0.2}_{-0.3}$ | $13.67^{+0.18}_{-0.25}$ | $12.25^{+0.11}_{-0.12}$ | $12.42^{+0.22}_{-0.12}$ | $12.06^{+0.09}_{-0.13}$ | $<12.1$ | $<11.1$ | "
$C,14$ | $0.342930(^{+14}_{-14})$ | -1769 | $25^{+4}_{-4}$ | $2.8^{+0.3}_{-0.6}$ | $4.0^{+0.3}_{-0.1}$ | $14.91^{+0.26}_{-0.15}$ | $12.48^{+0.10}_{-0.19}$ | $12.99^{+0.09}_{-0.16}$ | $<12.1$ | $13.03^{+0.11}_{-0.27}$ | $<11.6$ | "
$C,15$ | $0.342996(^{+4}_{-5})$ | -1754 | $14.2^{+1.1}_{-1.0}$ | $3.6^{+0.6}_{-0.4}$ | $3.7^{+0.2}_{-0.3}$ | $13.74^{+0.23}_{-0.22}$ | $12.39^{+0.09}_{-0.08}$ | $12.75^{+0.12}_{-0.12}$ | $12.34^{+0.05}_{-0.08}$ | $12.60^{+0.18}_{-0.27}$ | $11.81^{+0.16}_{-0.19}$ | "
$C,16$ | $0.3430353(^{+59}_{-29})$ | -1745 | $68.6^{+3.2}_{-2.4}$ | $4.18^{+0.05}_{-0.07}$ | $4.0^{+0.1}_{-0.1}$ | $14.54^{+0.04}_{-0.03}$ | $12.67^{+0.02}_{-0.02}$ | $13.55^{+0.03}_{-0.02}$ | $12.27^{+0.05}_{-0.04}$ | $12.88^{+0.09}_{-0.11}$ | $12.15^{+0.14}_{-0.21}$ | "
$C,17$ | $0.343601(^{+10}_{-7})$ | -1620 | $19^{+4}_{-3}$ | $5.2^{+1.3}_{-1.0}$ | $4.2^{+0.5}_{-0.3}$ | $13.16^{+0.24}_{-0.22}$ | $11.48^{+0.08}_{-0.10}$ | $12.53^{+0.09}_{-0.10}$ | $<11.4$ | $<11.9$ | $<11.9$ | "
$C,18$ | $0.343902(^{+16}_{-20})$ | -1553 | $47^{+8}_{-6}$ | $4.5^{+0.5}_{-0.4}$ | $4.2^{+0.3}_{-0.5}$ | $13.71^{+0.07}_{-0.13}$ | $11.79^{+0.08}_{-0.04}$ | $12.75^{+0.06}_{-0.09}$ | $11.63^{+0.09}_{-0.09}$ | $<11.9$ | $<11.9$ | "

  * •

† relative to $z_{\rm sys}=0.3509$ (as determined in Sect. 3.1).

  * •

‡ covering factor

####  3.4.1 On the possibility of UV pumping

We tried to model the excitation of the observed Fe ii levels by UV pumping
instead of collisions with electrons since the UV flux can be very high for
gas in the vicinity of the central engine. To do this, we used the data of
transition probabilities from the NIST database to calculate the excitation
through UV pumping. We note that UV pumping can easily be incorporated into
the fit only in the optically-thin regime (corresponding to $\log
N(\text{Fe\,{\sc ii}})\lesssim 13$), which is not the case for most
components. Therefore, a complex multiple-zone excitation model should be
implemented that takes into account radiative transfer fully, since the UV
excitation at some position depends on the line profiles, and therefore on the
excitation balance at the regions closer to the radiation source. This
implementation however is impractical for such a complex line-profile fitting
as we have towards J 1652$+$2650\. However, we can draw qualitative
conclusions from the following two limiting cases: the optically-thin limit,
or assuming constant dilution of the excitation using typically-observed
column densities.

In the optically-thin case, we found that UV pumping does not provide
satisfactory fits, since it cannot reproduce the observed excitation of the Fe
ii levels as well as collisional excitation can do. To qualitatively
illustrate this, we plot in Fig. 6 the excitation of Fe ii levels as a
function of electron density and UV field. The UV field is expressed in terms
of distance to the central engine estimated from the observed $r$-band J
1652$+$2650 magnitude of $17.0$ assuming a typical quasar spectral shape. One
can see that the Fe ii excitation described by typically-estimated electron
densities (for example, at $n_{e}\approx 10^{4}\,\rm cm^{-3}$, which are shown
by a dashed line in each panel of Fig. 6), corresponds to roughly a factor of
two difference in distance (hence a factor of four difference in UV flux) that
is needed to describe the fine-structure levels of the ground term ($\rm
3d^{6}4s\,^{6}D$, representing excited levels from 1th to 4th) and the second
excited ($\rm 3d^{6}4s\,^{4}D$, representing excited levels from 9th to 12th)
Fe ii term. From this diagram, one can see that if the excitation is described
by UV pumping the absorbing gas must be located a few tens of parsec away from
the central engine. However, if UV pumping dominates the excitation of the Fe
ii levels, this results in very large values of the ionization parameter,
$\log U\gtrsim-1$, that is difficult to reconcile with the survival of Fe ii
and other associated low-ionization species (e.g. Ca ii and Na i).
Additionally, the observed column density of He i⋆ indicates $\log U\approx-3$
for dense gas, which is discussed later in Sect. 4. All this indicates that UV
pumping is unlikely to be the dominant excitation process at play in the gas.
Using calculations in the optically-thick regime and assuming $\log
N(\text{Fe\,{\sc ii}})=14$ we found less disagreement in the UV fluxes
required to populate the low and high Fe ii levels. However, the optically-
thick case implies smaller distances to the central engine and hence even
higher ionization parameters, in comparison to the optically-thin case.

Figure 6: Excitation of Fe ii levels as a function of electron density,
$n_{e}$, and distance to the central engine, $d$. The latter is calculated
from the measured photometric flux of J 1652$+$2650 assuming a typical quasar
spectral shape. Calculations were performed in the optically-thin limit, hence
any constraint on the distance that can be obtained if UV pumping dominates
the Fe ii excitation must be considered as an upper limit. Additionally, we
plot lines of constant ionizing parameter, $U$, calculated by scaling the UV
flux and assuming bluewards of the Lyman limit an AGN spectrum with power-law
shape with index $-1.2$. One can see that, for UV pumping to dominate the
excitation of Fe ii levels, the ionizing parameter must be larger than 0.01.

##  4 Photo-ionization model

We modelled the abundance of He i⋆ to estimate the physical conditions of the
gas associated with this small-detached narrow/low-ionization BAL towards J
1652$+$2650\. As discussed by, e.g., Arav et al. (2001) and Korista et al.
(2008), the meta-stable $2\,^{3}S$ level of He i⋆ is mostly populated through
He ii recombination and depopulated by radiative transition and collisional
de-excitation. Therefore, He i⋆ predominantly originates from a layer of
ionized gas where helium is in the form of He ii and $n_{e}\approx n_{\rm H}$,
and the He i⋆ column density is sensitive to the number density and ionizing
flux. In that sense, He i⋆ is an exceptional diagnostic of the physical
conditions, almost independent of metallicity and depletion, unlike other
metals. This is particularly relevant for the FeLoBAL under study since we can
measure neither the abundance of H i nor the total abundance of any metal
(i.e., only the singly-ionized state of each species is constrained), hence we
have neither a measurement of metallicity nor metal depletion in this system.
This limitation is also an issue for most of the previously studied FeLoBAL
systems, and the assumption of a particular metallicity value can
significantly affect the physical conditions derived from the photo-ionization
modelling (e.g., Byun et al., 2022b).

We used the latest public version of the Cloudy software package C17.02
(Ferland et al., 2017) to model a slab of gas in the vicinity of the AGN. Our
basic setup is a cloud of constant density that is illuminated on one side by
a strong UV field with a typical AGN spectrum. We assumed a metallicity of
$0.3$ solar333We took solar relative abundances of the metals, i.e. we did not
use any depletion factor. While in the most known FeLoBALs the metallicity is
found to be around solar value (e.g. Arav et al., 2001; Aoki et al., 2011;
Byun et al., 2022b), our chosen value 0.3 mimics possible Fe depletion, which
typically large (up to 2 dex) at solar metallicity., a characteristic value
for such clouds, but we also checked that the exact metallicity value has
little impact on the derived He i⋆ column densities. Temperature balance was
calculated self-consistently. As a stopping criterion, we used a total Fe ii
column density of $10^{15}$ cm-2 corresponding to the higher end of Fe ii
column densities observed within the FeLoBAL components.

We ran a grid of photo-ionization models by varying two main parameters: the
number density and the ionization parameter, within the ranges of $\log[1;6]$
and $\log[-4;0]$, respectively. Fig. 7 shows the constraints on each parameter
derived from the comparison of the modelled He i⋆ column density with the
fiducial value of $13.7\pm 0.1$, typical of high column-density components.
One can see that the modelling provides estimates on the number of ionizing
photons of $\log(Un_{\rm H})\approx 0.5$ for $\log n_{\rm H}\lesssim 4$, and
$\log U\approx-3$ for $\log n_{\rm H}\gtrsim 4$. The latter solution to
preferred value by the excitation of Fe ii, which provides an independent
constraint on $\log n_{\rm H}\approx 4$. Since the excited Fe ii levels
predominately arise from the ionized medium (as they are excited by collision
with electrons; see also Korista et al. 2008), this suggests that $n_{\rm
e}\approx n_{\rm H}$ and hence likely, the He i⋆ abundance provides an
estimate of $\log U\sim-3$. While the exact value of $\log U$ for each
component depends on the observed He i⋆ column density, we refrain from using
the latter because with such a modelling we cannot be confident regarding
constrained Fe ii column densities, as mutual covering may impact the derived
column densities. We also checked that the Cloudy modelling roughly reproduces
the Ca ii, Mg i, and Mn ii column densities. However, as we mentioned above,
using the abundance of these species is limited due to unconstrained total
metallicities and depletion patterns.

Figure 7: Constraints on the ionization parameter and number density from
Cloudy photo-ionization modeling of He i⋆ in the FeLoBAL towards J
1652$+$2650\. The region displayed in brown shows the estimated probability
density function constructed by comparing the calculated He i⋆ column
densities with the observed typical value of $\log N=13.7\pm 0.1$ (see Table
1). The region displayed in blue shows the range of $n_{\rm H}$ values
favoured by the measured Fe ii excitation.

##  5 Discussion

###  5.1 Case of FeLoBAL towards Q 2359$-$1241

One of the most comprehensive studies so far of a FeLoBAL by means of high-
resolution spectroscopy concerns Q 2359$-$1241 (Arav et al., 2001). A broad
and deep VLT/UVES spectrum of this quasar allowed Arav et al. (2008) to detect
Fe ii lines up to the $8^{\rm th}$ excited level (excitation energy of 7955
cm-1) above the ground state, and to constrain the physical conditions in the
associated medium (Korista et al., 2008). A sophisticated fitting model was
used, describing partial covering of the source by the absorbing clouds based
on a power-law distribution (see Arav et al., 2008). In our present study, we
used a model of uniform partial covering instead, which is dictated by the
observation of complex mutually-blended absorption-line profiles. In the case
of Q 2359$-$1241, the line velocity structure is simpler, with only a few
visually-distinct velocity components, and the lines are not significantly
saturated. This allows one to independently constrain the column density in
each Fe ii level and for each component, and then describe the population of
Fe ii levels to constrain the excitation mechanisms. Additionally, due to its
higher redshift, the FeLoBAL towards Q 2359$-$1241 allows one to constrain the
column density of intermediate Fe ii levels with energies between 1872$-$3117
cm-1 (corresponding to the $\rm a^{4}F$ term of the 3d7 configuration). We
found that these levels are important to disentangle between radiative pumping
and collisions with electrons. Certainly, such an independent determination of
Fe ii column densities is more robust than tightening them assuming a dominant
excitation mechanism, as we did for J 1652$+$2650\. Therefore, we endeavoured
to test our procedure by also fitting the FeLoBAL towards Q 2359$-$1241 using
the spectrum taken from the SQUAD database (Murphy et al., 2019).

To fit the Fe ii lines, we used an eight-component fit, out of which four
components exhibit higher excitation and are mutually blended, and four
components show a low level of excitation with only the first few excited
levels above the ground state detected Bautista et al. (2010). We used the
same Doppler parameters for all of the Fe ii levels in a given component. We
added an independent covering factor to each component, yet tying the covering
factor to be equal in two closely-associated weak components at $z=0.8611$.
The quasar continuum was re-constructed locally by interpolating the spectrum
free from absorption features. We note that the continuum placement may be
important for weak lines since the line profiles are fairly broad. The line-
fitting procedure we used is the same as described for J 1652$+$2650 (see
Sects. 3.4). The fitting results are listed in Table 2 and the modelled line
profiles are shown in the Appendix, in Figs. 16 to 20. In comparison to the
study of Arav et al. (2008) and Korista et al. (2008), we were able to
identify a larger number of Fe ii levels, up to the $12^{\rm th}$ excited
level of Fe ii (excitation energy of $\sim$8850 cm-1) above the ground state.

Table 2: Results of simultaneous fits to absorption lines from Fe ii ground
state and excited levels in the FeLoBAL towards Q 2359$-$1241.

Comp. | #1 | #2 | #3 | #4 | #5 | #6 | #7 | #8
---|---|---|---|---|---|---|---|---
$z_{\rm abs}$ | $0.859761(^{+10}_{-16})$ | $0.859922(^{+7}_{-7})$ | $0.8599318(^{+14}_{-9})$ | $0.860189(^{+6}_{-5})$ | $0.861056(^{+11}_{-6})$ | $0.861127(^{+26}_{-15})$ | $0.8618265(^{+17}_{-7})$ | $0.8626427(^{+13}_{-10})$
$\Delta v^{\rm\dagger}$ [km s-1] | -1325 | -1299 | -1298 | -1256 | -1117 | -1105 | -992 | -861
$b$ [km s-1] | $85.3^{+3.4}_{-1.9}$ | $47.5^{+2.0}_{-1.8}$ | $13.5^{+0.3}_{-0.3}$ | $15.9^{+1.0}_{-1.2}$ | $7.5^{+3.2}_{-1.5}$ | $20.1^{+2.5}_{-3.1}$ | $5.7^{+0.4}_{-0.4}$ | $7.5^{+0.3}_{-0.3}$
$\log N$(Fe ii,g.s.) | $15.39^{+0.12}_{-0.06}$ | $14.39^{+0.07}_{-0.04}$ | $13.63^{+0.03}_{-0.03}$ | $13.40^{+0.09}_{-0.17}$ | $13.45^{+0.18}_{-0.18}$ | $13.58^{+0.13}_{-0.17}$ | $13.01^{+0.07}_{-0.10}$ | $13.24^{+0.05}_{-0.04}$
$\log N$(Fe ii,j1) | $14.76^{+0.06}_{-0.07}$ | $13.91^{+0.07}_{-0.05}$ | $13.21^{+0.04}_{-0.04}$ | $13.01^{+0.13}_{-0.15}$ | $12.53^{+0.17}_{-0.25}$ | $12.56^{+0.26}_{-0.55}$ | $12.01^{+0.12}_{-0.12}$ | $12.48^{+0.03}_{-0.06}$
$\log N$(Fe ii,j2) | $14.50^{+0.06}_{-0.07}$ | $13.78^{+0.06}_{-0.09}$ | $13.10^{+0.04}_{-0.04}$ | $12.81^{+0.14}_{-0.13}$ | $12.63^{+0.12}_{-0.29}$ | $12.60^{+0.25}_{-0.58}$ | $11.68^{+0.27}_{-0.39}$ | $12.12^{+0.09}_{-0.09}$
$\log N$(Fe ii,j3) | $14.64^{+0.05}_{-0.08}$ | $13.43^{+0.10}_{-0.05}$ | $12.97^{+0.04}_{-0.03}$ | $12.73^{+0.10}_{-0.13}$ | $12.40^{+0.21}_{-0.42}$ | $12.48^{+0.21}_{-0.96}$ | $11.55^{+0.26}_{-0.33}$ | $12.20^{+0.06}_{-0.09}$
$\log N$(Fe ii,j4) | $14.25^{+0.06}_{-0.08}$ | $13.13^{+0.13}_{-0.14}$ | $12.67^{+0.07}_{-0.03}$ | $<12.5$ | $<12.2$ | $<12.1$ | $<11.7$ | $<11.5$
$\log N$(Fe ii,j5) | $15.33^{+0.10}_{-0.08}$ | $14.26^{+0.10}_{-0.08}$ | $13.70^{+0.05}_{-0.04}$ | $13.22^{+0.32}_{-0.83}$ | $<13.5$ | $<13.7$ | $12.64^{+0.19}_{-0.30}$ | $<12.6$
$\log N$(Fe ii,j6) | $14.92^{+0.09}_{-0.09}$ | $<13.5$ | $13.29^{+0.07}_{-0.09}$ | $<13.0$ | … | … | … | …
$\log N$(Fe ii,j7) | $14.63^{+0.20}_{-0.16}$ | $13.93^{+0.19}_{-0.31}$ | $<13.1$ | $<12.8$ | … | … | … | …
$\log N$(Fe ii,j8) | $14.88^{+0.08}_{-0.11}$ | $<12.80$ | $<13.0$ | $<12.2$ | … | … | … | …
$\log N$(Fe ii,j9) | $14.02^{+0.12}_{-0.08}$ | $<12.8$ | $12.55^{+0.05}_{-0.05}$ | $12.51^{+0.20}_{-0.30}$ | … | … | … | …
$\log N$(Fe ii,j10) | $14.20^{+0.08}_{-0.06}$ | $12.99^{+0.16}_{-0.12}$ | $12.16^{+0.09}_{-0.12}$ | $<11.8$ | … | … | … | …
$\log N$(Fe ii,j11) | $13.75^{+0.12}_{-0.15}$ | $<12.2$ | $12.00^{+0.09}_{-0.15}$ | $<11.3$ | … | … | … | …
$\log N$(Fe ii,j12) | $13.46^{+0.09}_{-0.11}$ | $<12.3$ | $11.50^{+0.24}_{-0.54}$ | $12.87^{+0.30}_{-0.44}$ | … | … | … | …
$\log N_{\rm tot}$ | $15.95^{+0.04}_{-0.04}$ | $14.86^{+0.05}_{-0.05}$ | $14.23^{+0.03}_{-0.04}$ | $13.80^{+0.12}_{-0.14}$ | $13.58^{+0.22}_{-0.15}$ | $13.83^{+0.12}_{-0.21}$ | $13.20^{+0.08}_{-0.10}$ | $13.39^{+0.05}_{-0.05}$
$C_{f}$‡ | $0.050^{+0.002}_{-0.005}$ | $0.13^{+0.01}_{-0.01}$ | $0.36^{+0.01}_{-0.02}$ | $0.14^{+0.03}_{-0.03}$ | $0.11^{+0.01}_{-0.01}$ | $0.11^{+0.01}_{-0.01}$ | $0.32^{+0.04}_{-0.04}$ | $0.31^{+0.02}_{-0.02}$

  * •

† relative to $z_{\rm sys}=0.868$ (as reported by Brotherton et al., 2001).

  * •

‡ covering factor

We used the measured population of the Fe ii levels to constrain the physical
conditions in the absorbing medium. We used the same model as for J
1652$+$2650, where we considered the competition between collisions (with
electrons) and radiative excitation (by UV pumping). In Figs. 8 and 9, we show
the excitation diagrams of the different Fe ii levels together with the
constrained region of the parameter space of physical conditions, i.e.,
electron density, $n_{\rm e}$, and UV field strength. As in Sect. 3.4.1, the
UV field is expressed in terms of distance to the central engine as estimated
from the $r$-band Q 2359$-$1241 magnitude of $\sim 17.0$ assuming a typical
quasar spectral shape (Selsing et al., 2016). The 2D posterior parameter
distributions were obtained using the likelihood function assumed to be a
product of individual likelihoods of the comparison between the modelled and
measured Fe iii∗/Fe ii ratios444The latter ratios were obtained directly from
the fitting procedure. During the fit of the Fe ii lines, the column densities
of the different energy levels of Fe ii were fitting parameters. However,
since we used the MCMC method to sample the posterior distributions, we could
simply obtain the posterior distributions for the aforementioned ratios.. We
also assumed flat priors on $\log n_{e}$ and $\log d$ emulating a wide prior
distribution for these two parameters. In Figs. 8 and 9, one can see that in
the components that have a large enough number of measured Fe ii levels, the
excitation is better reproduced by the model of collisions with electrons
only, and therefore these components provide robust constraints on the
electron density. For the components at $\Delta v=-1325$, $-1299$, and $-1298$
km s-1, we found $n_{e}$ to be in the range between $5\times 10^{3}$ and
$3\times 10^{4}$ cm-3. For the other components, the constrained posterior
does not indicate a preferable source of excitation leaving the physical
conditions poorly constrained over a wide range. However, for the weaker and
most-redshifted components at $\Delta v>-1200$ km s-1, the excitation of the
Fe ii levels is less and if it were dominated by collisions this would result
in a significantly lower electron density, $\log n_{e}\sim 3.5$, than for the
main components. In the three bluest components, where $n_{e}$ is robustly
measured, we can get an upper limit on the ionization parameter which was
found to be $\log U\lesssim-3$, $-1.5$, and $-2.5$ for the components at
$\Delta v=-1325$, $-1299$, and $-1298$ km s-1, respectively. These values are
reasonably consistent with the constraint of $\log U\sim-2.4$ obtained from
photo-ionization modelling of this system by Korista et al. (2008). It is also
in line with the characteristic values we obtained in Sects. 3 and 4 in the
case of the FeLoBAL towards J 1652$+$2650.

---
Figure 8: Left panels: Excitation diagrams of Fe ii levels in the FeLoBAL
towards Q 2359$-$1241\. The y-axes indicate the ratio of the measured column
density divided by statistical weight of i-th level to the ground level, while
the x-axes provide the energy of the levels. The text in the upper part of
each panel indicates the level terms. Each panel corresponds to a given
velocity component as indicated in a green box on top of the panel. Right
panels: Constrained physical conditions using the excitation of the Fe ii
levels shown in the left panels. The solid and dashed lines correspond to the
$1\sigma$ and $2\sigma$ confidence intervals of the 2D posterior probability
function, respectively. The violet and red curves on the top x-axis and right
y-axis, respectively, show the 1D marginalized probability functions. The red
and violet hatched regions below them indicate the approximate solutions where
the population of the levels is dominated by collisions or radiative
excitation, respectively. For illustrative purposes, the corresponding regions
of the excitation diagrams are shown in the left panels using the same colours
and hatch code.
---
Figure 9: Continuation of Fig. 8.

In comparison with Korista et al. (2008) and Bautista et al. (2010), we
derived significantly higher column densities for Fe ii levels in all
components (except at $-992$ km s-1) and total column density as well, which
is dominated by the velocity components at $-1325$ and $-1299$ km s-1. This
discrepancy is explained by the small covering factors of these two
components, which allow the lines to be significantly saturated. However, the
relative excitation of the Fe ii levels even in these saturated components
remains similar. Furthermore, the two central (in terms of apparent optical
depth) components at $-1298$ and $-1256$ km s-1 indicate an excitation
structure consistent with the results of Korista et al. (2008). We note that
determining the exact profile decomposition is not trivial in such systems. We
attempted to increase the number of fitted velocity components in Fe ii lines
but obtained more or less similar results, as additional components turned out
to be weak if any. Moreover, the systematic uncertainty is most likely
dominated by our choice of partial covering model, which is uniform with no
mutual intersection (see Sect. 3.4). In terms of number density, we obtained
similar results as Korista et al. (2008) who reported that $\log n_{\rm H}\sim
4.4\pm 0.1$ considering total Fe ii column densities only and a smaller number
of energy levels than we do.

Overall, our approach (multi-component model with a uniform covering factor)
provides well-consistent results with those from earlier works, especially for
derived physical quantity. While FeLoBALs in most cases are quite complicated
for spectral analysis, such systems as one towards Q 2359$-$1241 provide an
important example and testbed for the assumptions (e.g. collisionally-
dominated excitation) that can ease the analysis of more complex objects.

###  5.2 Similarities and differences between FeLoBALs

Previous studies indicate a wide range of physical conditions in FeLoBALs (and
other intrinsic Fe ii absorbers) with some kind of bimodal distribution, where
part of the population is located at mild distances of $\sim 0.1-10$ kpc and
has number densities of $<10^{5}$ cm-3, while the second part shows more
extreme properties with number densities $>10^{8}$ cm-3 and distances to the
nuclei down to $\sim 1$ pc. This may be in line with state-of-the-art models
of AGN outflows formation (Faucher-Giguère & Quataert, 2012; Costa et al.,
2020), where the wind is a complex phenomenon that is driven by different
mechanisms at different scales. It can be launched in the close vicinity of
the accretion disk by radiative pressure and at a much larger distance by
shocks, produced either by a wind from the accretion disk or the jet (e.g.
Proga et al., 2000; Costa et al., 2020).

On the other hand, we note that ionization parameters around $-1..-3$ have
been found for almost all detected FeLoBALs, while one would expect a wider
range of values. This can be a selection effect, where such values of the
ionization parameter are favourable for FeLoBAL observation. However, since
$U\propto n^{-1}d^{-2}$ the observed bimodality can be an artefact of improper
constraints on the physical conditions from the modelling. Indeed, the
modelling of FeLoBALs is quite complex and ambiguous, and there are many
factors that cannot be resolved using line-of-sight observations. FeLoBALs
always exhibit a multi-component structure, which makes the derivation of the
physical conditions a difficult task, since several solutions are possible. A
typical question arising from the photo-ionization modelling setup is what is
the relative position in physical space of the "gaseous clouds" associated
with each component? This impacts both the column density estimation, due to
unknown mutual partial covering, and ionization properties, since the closest
to AGN clouds will shield the farthest located ones. A relevant example of
this situation is the FeLoBAL towards J 104459.6$+$365605, where Fe ii
excitation and simplistic photo-ionization modelling suggest relatively low
number densities, $\log n_{\rm e}<4$, and a distance of $\sim 700$ pc, while
more sophisticated wind models (Everett et al., 2002), which take shielding
effects into account, yield number densities $10^{4}$ times higher, and a
distance of $\sim 4$ pc. We note however that in the latter model, the
excitation of Fe ii levels is expected to be much higher than what is
observed. In that sense, the usage of the excited fine-structure levels and He
i* may provide less degenerated constraints than the relative abundances of
the ions, the latter is also suffering from complications due to the unknown
metallicity and depletion pattern (see discussion in Sect. 4).

###  5.3 Relation between FeLoBALs and other intrinsic absorbers

It is worth mentioning that a recently-identified class of associated quasar
absorbers bearing H2 molecules exhibits distances to the AGN of $\sim 1-10$
kpc (Noterdaeme et al., 2019, 2021, 2023), slightly higher, but comparable to
what is derived for FeLoBALs. Interestingly, while the medium in such systems
is neutral (and even at the H i-H2 transition), in contrast with FeLoBALs
which arise in the ionized phase or at the boundary of the ionization front
(e.g., Korista et al., 2008), they exhibit number densities of $\gtrsim
10^{4}$ cm-3, similar to FeLoBALs. In the case of H2-bearing systems, such
number densities are required for H2 to survive in the vicinity of the AGN
(Noterdaeme et al., 2019) as the radiation fields are greatly enhanced. While
in the case of J 1652$+$2650 we were not able to get constraints on the H2
column density since the lines are out of the range of the spectrum, there is
no H2 detection in other FeLoBAL so far. Additionally, presented Cloudy
modelling (given in Sect. 4) suggests that the column densities and ionization
parameter are not enough in FeLoBAL towards J 1652$+$2650 to expect the
presence of the H2 in such kind of medium.

Therefore, the difference between FeLoBALs and H2-bearing systems may be
related to the lower ionization parameters of the latter (akin to higher lower
incident UV flux, or to their larger distances), which makes it possible for
H2 to survive or imply reasonable timescales to form H2. To elaborate on this
we ran a grid of Cloudy models to see how the conditions for the presence of
H2 and excited Fe ii levels compared in the physical parameter space. We
considered an isobaric model of the medium with 0.3 metallicity relative to
solar (we additionally scaled the Fe abundance by 0.3 to emulate typical
depletion at such metallicity) exposed by the AGN-shaped radiation field and
regular cosmic ray ionization rate, $2\times 10^{-16}\,\rm s^{-1}$ (for atomic
hydrogen). We varied the ionization parameter and thermal pressure in ranges
$\log U=-6..0$ (with 0.2 dex step) and $P_{\rm
th}=10^{5}..10^{10}\,\rm[K\,cm^{-3}]$ (with 0.5 dex step), respectively. We
stopped the calculations either when the total H2 or Fe ii column densities
reached characteristic values of $\log N(\rm H_{2})=20$ or $\log
N(\mbox{Fe\,{\sc ii}})=15$, respectively. The obtained contours of Fe ii
excitation and the total H2 column density are shown in Fig. 10.

One can see that indeed, large H2 column densities are found mostly outside
the region of the parameter space where Fe ii is highly excited, i.e.
$\log\mbox{Fe\,{\sc ii}}^{*}/\mbox{Fe\,{\sc ii}}>-1$ (which is typical for the
observed FeLoBAL systems), except only for the very high thermal pressures,
$P_{\rm th}\gtrsim 10^{9}\,\rm K\,cm^{-3}$ and low ionization parameters $\log
U\lesssim-4$. The presence of H2 is mostly limited by the distances to the
AGN, which in the case of J 1652$+$2650 corresponds to the values of several
hundreds of pc555Only a stationary cloud is considered here, i.e., we do not
take into account time-dependent effects, which may be important for H2 in
harsh environments.. In turn, the values measured Fe ii∗/Fe ii coupled with
the constrained ionization parameter $\log U\sim-3$ points to a low molecular
fraction. We note that this modelling is indicative only, and one needs to be
careful when comparing the measured excitation of the Fe ii levels with the
modelled one in this simulation. Indeed, in this particular modelling case, we
stopped at a relatively large total Fe ii column density, at which we may
include a significant part of the neutral medium, where Fe ii excitation is
not so high as in the ionized shell.

Figure 10: Constraints on the ionization parameter and thermal pressure from a
grid of Cloudy photo-ionization models (for setup, see text). The violet
colour gradient and dashed contours show the excitation of Fe ii, defined as
$\log N(\mbox{Fe\,{\sc ii}}^{*})/N(\mbox{Fe\,{\sc ii}})$. The characteristic
range of Fe ii excitation observed in J 1652$+$2650 is shown by the hatched
region. The red solid contours indicate the total H2 column density, marked by
values of $\log N(\rm H_{2})$. The blue dotted contours show the distance to
the central engine calculated using the parameters of J 1652$+$2650\.

The region in $(\rm U,P_{th})$ parameter space where $\rm H_{2}$ is present in
copious amount is adjacent to the region where Fe ii is reasonably excited
(e.g. $\log N(\mbox{Fe\,{\sc ii}}^{*})/N(\mbox{Fe\,{\sc ii}})>-1$). That
indicate that we may witness the appearance of a global natural sequence among
associated absorbers. The behaviour of this sequence is probably coupled with
hydrodynamics processes in the outflowing gas, which set the thermal pressures
and the sizes of the clumps and their dependence on the distance. Likely, a
part of this sequence was observationally noticed by Fathivavsari (2020),
regarding Coronographic (Finley et al., 2013) and Ghostly (Fathivavsari et
al., 2017) DLAs, which also exhibit high excitation of fine-structure
levels666There are examples of similar associated systems with fine-structure
excitation observed in the past, e.g., (Hamann et al., 2001)., but that do not
show Fe ii∗. The physical connection between these different classes of
absorbers seems to be evident, with FeLoBALs representing predominantly gas at
the ionization front, Coronographic and Ghostly DLAs being predominantly
neutral, and associated H2-bearing absorbers tracing the H i-to-H2 transition.
Importantly, both the ionization and photo-dissociation fronts are controlled
by the ratio of UV flux to number density, i.e., the ionization parameter,
while the appearance of a certain class of absorber also depends on the
ambient thermal pressure and the total column density of the medium. In that
sense, it will be important to observe and study systems of intermediate
classes, displaying mixed properties. The search for such rare systems will be
supported by the upcoming next-generation wide-field spectroscopic surveys
such as 4MOST (Krogager et al., 2023), DESI (Chaussidon et al., 2023), and
WEAVE (Jin et al., 2023).

##  6 Summary

In this paper, we presented an analysis of the serendipitously-identified
FeLoBAL system at z=0.3509 towards J 1652$+$2650 performed using a high-
resolution UVES spectrum. The main aim was to derive constraints on the
physical conditions in the absorbing medium located near the AGN central
engine.

The absorption system consists of three kinematically-detached absorption
complexes spanning -1700 to -5700 $\rm km\,s^{-1}$ relative to the QSO
redshift. We detected lines profiles of Mg ii, Mg i, Ca ii, He i⋆, Mn ii and
Fe ii. For, the latter species we detected lines from the various fine-
structure levels of the ground and second excited electronic states, with
energies up to $\approx 8850\,\rm cm^{-1}$. The lines indicate a partial
coverage of the continuum emission source, with a covering factor in the range
from 0.98 to 0.2. A relatively simple kinematic structure (in comparison to
the majority of known FeLoBALs) and an intermediate saturation allow us to
perform joint multi-component Voigt profile fitting of the aforementioned
species (except Mg ii) to derive column densities in the simplistic
homogeneous partial coverage assumption. Using an additional assumption during
the fit, that excitation of Fe ii levels dominated by the collisions with
electrons, we obtained the constraints on the electron density in the medium
to be $\sim 10^{4}\,\rm cm^{-3}$ with $\sim 1$ dex dispersion. We also
detected the lines from the first excited level of Mn ii and constrained Mn
ii⋆/Mn ii column density ratios to be in the range from 0.1 to 0.5 across
velocity components. However, the lack of collisional coefficients data for Mn
ii did not allow us to use Mn ii excitation to infer the physical condition in
the medium.

Among the other elements detected in this FeLoBAL, He i⋆ is most important
since it allows obtaining constraints on the combination of ionization
parameter and number density, even without measurement of the hydrogen column
density, metallicity and depletion pattern, which is the case of J
1652$+$2650\. We used Cloudy code to model the characteristic column densities
of He i⋆ and obtained a value of the ionization parameter $\log U\sim-3$
assuming the number density derived from Fe ii. Such values are typically
measured in FeLoBAL systems, which likely represents the similarity among
them, while line profiles in FeLoBALs can be drastically different. With the
estimate of the UV flux from J 1652$+$2650, this translates to a constraint on
the distance between the absorbing medium and the continuum source of $\sim
100$ pc.

We also discuss the connection of FeLoBAL systems with other types of
intrinsic absorbers, including Coronographic and recently identified proximate
H2-bearing DLAs. The latter indicates a similar value of the number densities
as measured in FeLoBAL, $\gtrsim 10^{4}$ cm-3. Using Cloudy modelling we
showed that the FeLoBAL and H2-bearing proximate systems located in the
adjacent regions in the parameter space, representing the main global
characteristics of the medium: thermal pressure and ionization parameter (or
the number density and the distance to the AGN, which are mutually
interconnected to them). This indicates a global natural sequence among
associated absorbers, where FeLoBALs represent predominantly gas at the
ionization front, Coronographic DLAs are predominantly neutral, and associated
H2-bearing absorbers trace the H i-to-H2 transition. This likely will be
comprehensively explored with upcoming next-generation wide-field
spectroscopic surveys. This shall greatly enhance our understanding of AGN
feedback and cool-gas flows from the AGN central engine.

## Data Availability

The data published in this paper is available through Open Access via the ESO
scientific archive, and in the SQUAD database (Murphy et al., 2019). The
reduced and co-added spectra can be shared upon request to the corresponding
author.

## Acknowledgements

We thank the anonymous referee for constructive report, that significantly
improved the paper. SAB acknowledges the hospitality and support from the
Office for Science of the European Southern Observatory in Chile during a
visit when this project was initiated. SAB and KNT were supported by RSF grant
23-12-00166. S.L. acknowledges support by FONDECYT grant 1231187.

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##  Appendix A Ca ii, Mg i, and He i⋆ fits

Here we present the figures of all the fitted Ca ii, Mg i, He i⋆, Fe ii and Mn
ii lines.

Figure 11: Left to Right: Voigt-profile fits to Ca ii, Mg i, and He i⋆
absorption lines in complexes $A$ (Left), $B$ (Middle), and $C$ (Right), at
$z_{\rm abs}=0.3253$, 0.3304, and 0.3429, respectively, towards J
1652$+$2650\. The graphical information is the same as in Fig. 4. The red
lines at the top of each panel show the residuals between the spectrum and
fit, with green dashed horizontal lines corresponding to 2$\sigma$ deviation.

Figure 12: Left to Right: Voigt-profile fits to Fe ii absorption lines in
complexes $A$ (Left), $B$ (Middle), and $C$ (Right), at $z_{\rm abs}=0.3253$,
0.3304, and 0.3429, respectively, towards J 1652$+$2650\. Horizontal dashed
lines and surrounding grey areas indicate the extent of partial covering
determined by fitting each clump independently with its own covering factor.

Figure 13: Same as Fig. 12 for higher-excitation Fe ii lines.

Figure 14: Same as Fig. 12 for higher-excitation Fe ii lines (continued).

Figure 15: Same as Fig. 12 for Mn ii lines.  Figure 16: Voigt-profile fits to
Fe ii absorption lines at $z_{\rm abs}=0.85993$ towards Q 2359$-$1241\. The
graphical information is similar to that in Fig. 11.  Figure 17: Same as Fig.
16 for higher-excitation Fe ii lines.  Figure 18: Same as Fig. 16 for higher-
excitation Fe ii lines (continued).  Figure 19: Same as Fig. 16 for higher-
excitation Fe ii lines (continued).  Figure 20: Same as Fig. 16 for higher-
excitation Fe ii lines (continued).  Figure 21: Same as Fig. 16 for higher-
excitation Fe ii lines (continued).