# Observations of the Bright Star in the Globular Cluster 47 Tucanae (NGC 104)

William V. Dixon Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  Pierre Chayer Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  Marcelo Miguel Miller Bertolami Instituto de Astrofísica de La Plata, CONICET-UNLP, Paseo del Bosque s/n, FWA, B1900 La Plata, Provincia de Buenos Aires, Argentina  Facultad de Ciencias Astronómicas y Geofísicas, Universidad Nacional de La Plata, Paseo del Bosque s/n, FWA, B1900 La Plata, Provincia de Buenos Aires, Argentina  Valentina Sosa Fiscella Facultad de Ciencias Astronómicas y Geofísicas, Universidad Nacional de La Plata, Paseo del Bosque s/n, FWA, B1900 La Plata, Provincia de Buenos Aires, Argentina  CCT La Plata, CONICET, Calle 8 No 1467, B1904 La Plata, Provincia de Buenos Aires, Argentina  Robert A. Benjamin Department of Physics, University of Wisconsin-Whitewater, 800 West Main Street, Whitewater, WI 53190, USA  Andrea Dupree Center for Astrophysics | Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138, USA

(Accepted 2021 June 23)

###### Abstract

The Bright Star in the globular cluster 47 Tucanae (NGC 104) is a post-AGB
star of spectral type B8 III. The ultraviolet spectra of late-B stars exhibit
a myriad of absorption features, many due to species unobservable from the
ground. The Bright Star thus represents a unique window into the chemistry of
47 Tuc. We have analyzed observations obtained with the Far Ultraviolet
Spectroscopic Explorer (FUSE), the Cosmic Origins Spectrograph (COS) aboard
the Hubble Space Telescope, and the MIKE Spectrograph on the Magellan
Telescope. By fitting these data with synthetic spectra, we determine various
stellar parameters ($T_{\rm eff}=10,850\pm 250$ K, $\log g=2.20\pm 0.13$) and
the photospheric abundances of 26 elements, including Ne, P, Cl, Ga, Pd, In,
Sn, Hg, and Pb, which have not previously been published for this cluster.
Abundances of intermediate-mass elements (Mg through Ga) generally scale with
Fe, while the heaviest elements (Pd through Pb) have roughly solar abundances.
Its low C/O ratio indicates that the star did not undergo third dredge-up and
suggests that its heavy elements were made by a previous generation of stars.
If so, this pattern should be present throughout the cluster, not just in this
star. Stellar-evolution models suggest that the Bright Star is powered by a
He-burning shell, having left the AGB during or immediately after a thermal
pulse. Its mass ($0.54\pm 0.16M_{\odot}$) implies that single stars in 47 Tuc
lose 0.1–0.2 $M_{\odot}$ on the AGB, only slightly less than they lose on the
RGB.

stars: abundances — stars: atmospheres — stars: individual (NGC 104 Bright
Star (catalog Cl* NGC 104 R 1)) — ultraviolet: stars

††journal: AJ††facilities: FUSE, HST (COS), Magellan:Clay (MIKE spectrograph)

##  1 Introduction

The Bright Star in the globular cluster 47 Tucanae (NGC 104) is the cluster’s
brightest member at both ultraviolet and optical wavelengths. It is Star No. i
in the catalog of Feast & Thackeray (1960), who refer to it as “the bright
blue star.” It is star R1 in the catalog of Lloyd Evans (1974). In his study
of UV-bright stars in globular clusters, de Boer (1985) calls the star “BS.”

This blue giant (B8 III) is a post-asymptotic giant branch (post-AGB) star.
Having ascended the AGB, it is moving across the color-magnitude diagram
toward the tip of the white-dwarf cooling sequence. Dixon et al. (1995)
observed the star with the Hopkins Ultraviolet Telescope (HUT; Davidsen et
al., 1992). By comparing the shape of its far-ultraviolet spectrum with those
predicted by Kurucz (1992) models, they derived an effective temperature
$T_{\rm eff}$ = 11,000 K and a surface gravity $\log g$ = 2.5.

While Kurucz models with [M/H] = $-1.0$ adequately reproduce the stellar
spectrum across most of the HUT bandpass (900–1850 Å), they over-predict the
flux at wavelengths shorter than Lyman $\beta$ (1026 Å). To understand this
discrepancy, Dixon & Chayer (2013) observed the star with the Far Ultraviolet
Spectroscopic Explorer (FUSE). They determined that the peculiar absorption
features in the stellar continuum are sculpted by resonances in the
photoionization cross section of excited-state neutral nitrogen in the stellar
atmosphere.

Table 1: Summary of FUSE, COS, and Magellan Observations

Instrument | Grating | Wavelength | $R\equiv\lambda/\Delta\lambda$ | Exp. Time | Obs. Date | Data ID | P.I.
---|---|---|---|---|---|---|---
|  | (Å) |  | (s) |  |  |
FUSE | $\cdots$ |  0905–1187 | 20,000 | 14,058 | 2006 Jun 7 | G0730101 | Dixon
COS | G130M | 1130–1430 | 18,000 |  00.524 | 2015 Mar 26 | LCKV01010 | Benjamin
COS | G160M | 1410–1780 | 18,000 |  00.764 | 2015 Mar 26 | LCKV01020 | Benjamin
MIKE-Blue | $\cdots$ | 3350–5000 | 36,000 |  00.900 | 2010 Jun 30 | Obj_mb0063 | Dupree
MIKE-Red | $\cdots$ | 4900–9300 | 26,000 |  00.900 | 2010 Jun 30 | Obj_mr0057 | Dupree

Note. — The MAST data used in this work are available at
https://doi.org/10.17909/t9-6s10-5z38 (catalog 10.17909/t9-6s10-5z38).

None of these studies attempted to derive the Bright Star’s chemical
abundances. The spectra of late-B stars exhibit a myriad of metal-line
absorption features. The Bright Star thus represents a unique window into the
chemistry of 47 Tuc. To take advantage of this opportunity, we have combined
the star’s FUSE spectrum with data obtained with the Cosmic Origins
Spectrograph (COS) aboard the Hubble Space Telescope and the MIKE Spectrograph
on the Magellan Telescope. A summary of these observations is presented in
Table 1.

In Section 2, we discuss each of our data sets. In Section 3, we present our
atmospheric models and use them to derive stellar parameters and abundances.
In Section 4, we consider various problems with the models, some that we are
able to solve, and some that we are not. In Section 5, we discuss our results:
we address the question of cluster membership in Section 5.1, we derive the
stellar mass and luminosity in Section 5.2, we consider the evolutionary state
of the star in Section 5.3, and we consider the implications of our chemical
abundances in Section 5.4. We summarize our conclusions in Section 6.

##  2 Observations and Data Reduction

###  2.1 FUSE Spectroscopy

FUSE provides medium-resolution spectroscopy from 1187 Å to the Lyman limit
(Moos et al., 2000; Sahnow et al., 2000). The Bright Star was observed through
the FUSE $30\arcsec\times 30\arcsec$ aperture. The data were reduced using
v3.2.2 of CalFUSE, the standard data-reduction pipeline software (Dixon et
al., 2007), and retrieved from the Mikulski Archive for Space Telescopes
(MAST). The signal-to-noise ratio (S/N) per resolution element is $<5$ at
wavelengths shorter than 1026 Å and 10–15 at longer wavelengths.

###  2.2 COS Spectroscopy

The Cosmic Origins Spectrograph (COS) enables high-sensitivity, medium- and
low-resolution spectroscopy in the 1150–3200 Å wavelength range (Green et al.,
2012). The Bright Star was observed with COS using both the G130M and G160M
gratings. The fully-reduced spectra, processed with CALCOS version 3.3.7
(Rafelski et al., 2018), were retrieved from MAST. The S/N per resolution
element varies between 20 and 40 in the G130M spectrum. In the G160M spectrum,
the S/N falls monotonically from $\sim 55$ at 1430 Å to $\sim 20$ at 1770 Å.

In the course of our analysis, we will have some questions about the
interstellar medium (ISM) along the line of sight to the Bright Star.
Fortunately, the cluster star UIT 14 (O’Connell et al., 1997) has been
observed with both the G130M and G160M gratings of COS (Program 11592; P.I.
Lehner). The star is only 1′ from our target, so samples the same line of
sight. Its high effective temperature ($T_{\rm eff}$ $\sim 50,000$ K) ensures
that all low-ionization features in its spectrum are interstellar.

###  2.3 MIKE Spectroscopy

The Bright Star was observed with the Magellan/Clay telescope at the Las
Campanas Observatory using the Magellan Inamori Kyocera Echelle (MIKE)
spectrograph (Bernstein et al., 2003) with a $0.70\times 5$ arcsecond slit.
The data were obtained as part of the observing run documented in Dupree et
al. (2011) and were reduced using the procedures described therein. The
authors quote a spectral resolution of $\sim$ 26,000 on the blue side and
$\sim$ 36,000 on the red, but these values are reversed; the resolution is
higher on the blue side. The S/N per resolution element varies within each
spectral order and with wavelength, reaching a maximum of about 450 at 4700 Å
on the blue side and 7500 Å on the red. The spectrum extends from the Balmer
edge to the Paschen series (lines P10 and higher) of hydrogen.

The orders of an echelle spectrum have a distinctive shape, which we must
remove before attempting to fit individual features. Dupree et al. (2011)
normalized their spectra with a cubic spline fit to the spectrum of the
standard white dwarf LTT 9491, but this curve does not perfectly reproduce the
shape of the Bright Star spectrum. In scaling the curve to match our data, we
have optimized the fit in regions around the hydrogen Balmer lines. Elsewhere,
particularly near the edges of an order, a small, residual tilt is often
apparent in the normalized spectrum.

###  2.4 Contamination

Globular clusters are crowded fields, and the Bright Star lies within an arc
minute of the cluster center. While the COS aperture is only 2$\farcs$5 in
diameter, the FUSE aperture is much larger. Fortunately, most bright cluster
stars are red giants and invisible in the ultraviolet. GALEX FUV images of the
cluster reveal no bright objects within 20″ of our target, and the cluster’s
FUV vs. (FUV–NUV) color-magnitude diagram shows that the Bright Star is nearly
three magnitudes brighter than any other star in the field (Schiavon et al.,
2012). In the optical, contamination is a larger threat. The MIKE aperture was
deliberately positioned to exclude two bright, nearby stars. (See the
identification chart in Forte et al. 2002.) As discussed in Section 3.5, the
fact that we derive the same high nitrogen abundance from both FUV and near-IR
features suggests that contamination of the MIKE spectrum is insignificant.

##  3 Spectral Analysis

###  3.1 Model Atmospheres

Because 47 Tuc has a relatively high metallicity ([Fe/H] = $-0.74$; Kovalev et
al. 2019), and because the spectra of B-type stars exhibit many absorption
features, we expect line opacity to play an important role in the atmospheric
structure of the Bright Star. We therefore employ the model atmospheres of
Castelli & Kurucz (2003). These models provide a thorough treatment of line
blanketing by metals, though the assumption of local thermodynamic equilibrium
(LTE) neglects non-LTE effects, which can be important for hot, low-gravity
stars. We use the $\alpha$-enhanced models ([$\alpha$/Fe] = +0.4) that are
generally appropriate for stars in a globular cluster. We use models with
metallicity [Fe/H] = $-1.0$, microturbulent velocity $\xi$ = 2.0 km s-1, and
mixing-length parameter $l/H$ = 1.25. From these model atmospheres, we compute
synthetic spectra using version 51 of the program SYNSPEC (Hubeny, 1988). Our
atomic models are similar to those used by Lanz & Hubeny (2007) to compute
their grid of B-type stars.

Because we use pre-computed model atmospheres, we cannot change their chemical
abundances. Instead, we adjust the abundances used by SYNSPEC to calculate the
model spectra. This technique makes the implicit assumption that small changes
in the photospheric abundances do not significantly alter the star’s
atmospheric structure. While this is certainly true for most metals, it is
probably not the case for helium, as discussed below.

Before comparison with the data, synthetic spectra must be convolved with the
appropriate line-spread function. For the FUSE spectrum, we use a Gaussian of
FWHM = 0.06 Å. For the COS spectra, we employ the tabulated line-spread
functions appropriate for data obtained at Lifetime Position #3, which are
available from the COS website. For the MIKE spectra, we assume a resolving
power of R = 36,000 on the blue side and 26,000 on the red, convolving the
spectrum with a Gaussian of FWHM = $\lambda/$R, where $\lambda$ is the
wavelength of the feature that we wish to fit. The resulting synthetic spectra
successfully reproduce the observed absorption features; in particular, we
find no evidence for rotational line broadening.

###  3.2 Spectral Fitting

Given a grid of synthetic spectra, our fitting routine linearly interpolates
among them—in one, two, or three dimensions, as appropriate—determining the
best fit to the data via chi-squared minimization. The uncertainties quoted
for parameters derived from individual line fits are $1\sigma$ errors computed
from the covariance matrix returned by the fitting routine; we refer to these
as statistical errors. In most cases, we ignore statistical errors when
combining the results from multiple line fits. Instead, we simply quote the
mean and standard deviation of the various measurements. We believe that this
approach is more appropriate when systematic uncertainties—such as the shape
and level of the stellar continuum—are the dominant source of error.

###  3.3 Stellar Parameters

We begin by constructing a three-dimensional grid of synthetic spectra with
effective temperatures $T_{\rm eff}$ = 10,000 to 11,750 K in steps of 250 K,
surface gravities $\log g$ = 2.0 and 2.5, and helium abundances $\log N({\rm
He})/N({\rm H})$ = $-0.6$ to $-1.4$ in steps of 0.2. Models with higher
temperatures or lower gravities are not available. The Castelli-Kurucz models
assume a helium abundance of 10% by number. We compute synthetic spectra with
other He/H ratios by adjusting the atmospheric abundances within SYNSPEC.

Figure 1: High-order Balmer lines in the Bright Star’s optical spectrum. The
black curve represents the data, the red curve our best-fit model.

Figure 2: Low-order Balmer lines in the Bright Star’s optical spectrum. Black
curves represent the data, red curves our best-fit model.

Numerical experiments show that the shape of the hydrogen series near the
Balmer jump is exquisitely sensitive to the surface gravity, while the low-
order Balmer features provide a better probe of effective temperature. We thus
determine these parameters using an iterative technique: (1) fix the
temperature and fit the high-order lines (3665–3700 Å) to determine the
gravity; (2) fix the gravity and fit the low-order lines to determine the
temperature (in this step, we consider only the H6–H10 lines, as discussed
below); (3) repeat until the values converge. Our best-fit parameters are
$T_{\rm eff}$ = 10,850 K and $\log g$ = 2.20. For these fits, we fix the
helium abundance at $\log N({\rm He})/N({\rm H})$ = $-1.0$.

Our results are presented in Figs. 1 and 2. We see that the high-order and
H6–H10 lines are well fit by our model spectrum. We are less successful with
the H$\alpha$–H$\epsilon$ lines: the model features are too wide and too
shallow. The H$\alpha$ line is particularly troublesome, because the fitting
routine drives the best-fit model beyond the limits of our grid. We will
return to this topic in Section 4; for now, we consider only the H6–H10 lines.

To estimate the uncertainty in these parameters, we determine them again in a
slightly different way. We re-fit each of the H6–H10 lines, this time allowing
both the effective temperature and surface gravity to vary. The resulting
parameters are $T_{\rm eff}$ = 11,100 K and $\log g$ = 2.33, where these
values are averages from fits to each of our six lines. (Six, because the H09
line appears in two orders, and we fit each copy separately.) We take the
difference between these and our previous results to be the uncertainty in
each parameter. Our final parameters are thus $T_{\rm eff}$ = $10,850\pm 250$
K and $\log g$ = $2.20\pm 0.13$.

Before fitting the high-order lines, we normalize the appropriate echelle
order (#26 in the blue spectrum) by hand, identifying the peaks between each
absorption line, fitting a spline to them, and dividing by the spline. We use
the same technique to normalize the model spectra.

The echelle orders containing the low-order lines are normalized as discussed
in Section 2.3. A residual tilt is apparent in the some regions of the
normalized spectra. To account for this tilt, we multiply our model spectra by
a linear function whose slope and y-intercept are free parameters in the fit.
We mask all nearby metal lines to prevent them from affecting the continuum
fit.

###  3.4 Helium Abundance

To constrain the star’s helium abundance, we fit selected He I features in its
optical spectrum. When features appear in multiple orders, we fit each copy
separately. Features blended with metal lines are omitted. We use the model
grid described above, but fix $T_{\rm eff}$ and $\log g$ at their best-fit
values and allow only the He/H ratio to vary. Results are presented in Table 2
and Fig. 3. We derive a helium abundance $\log N({\rm He})/N({\rm H})$ =
$-0.82\pm 0.16$, where the values represent the mean and standard deviation of
our twelve measurements.

This value of the helium abundance corresponds to a mass fraction $Y=0.37\pm
0.08$, which is considerably greater than the range $Y=0.26$–0.29 derived for
the various subpopulations of 47 Tuc (Heyl et al., 2015; Salaris et al., 2016;
Denissenkov et al., 2017). This apparent helium enhancement may be a non-LTE
effect. Takeda & Takada-Hidai (2000) determined the helium abundances of a
sample of late-B and early-A type supergiants. For the four stars whose
effective temperature and surface gravity are closest to those of the Bright
Star, they found that LTE models over-predict the helium abundance by
0.57–0.79 dex relative to non-LTE models.

Another possible explanation reflects a limitation in our models: Because our
synthetic spectra are derived from atmospheric models with a fixed helium
abundance (10% by number), they do not reflect variations in the atmospheric
structure that would accompany changes in the He/H ratio. Indeed, synthetic
spectra computed assuming $\log N({\rm He})/N({\rm H})$ = $-0.82$ yield the
same values of the mean N and Fe abundance as spectra computed assuming $\log
N({\rm He})/N({\rm H})$ = $-1.0$.

Table 2: Helium Features  $\lambda_{\rm lab}$ | Spectrum | Order | $\log N({\rm He})/N({\rm H})$
---|---|---|---
3819.6 | Blue | 22 | $-0.86\pm 0.04$
4026.2 | Blue | 18 | $-0.62\pm 0.03$
4026.2 | Blue | 17 | $-0.61\pm 0.03$
4120.8 | Blue | 16 | $-1.12\pm 0.06$
4143.8 | Blue | 15 | $-0.78\pm 0.02$
4387.9 | Blue | 10 | $-0.72\pm 0.02$
4471.5 | Blue |  09 | $-0.71\pm 0.02$
4471.5 | Blue |  08 | $-0.68\pm 0.03$
4713.1 | Blue |  05 | $-1.04\pm 0.03$
4713.1 | Blue |  04 | $-0.96\pm 0.05$
5875.6 | Red | 12 | $-0.91\pm 0.03$
5875.6 | Red | 13 | $-0.85\pm 0.02$

###  3.5 Metal Abundances

We derive the star’s metal abundances by fitting synthetic spectra to selected
absorption features in its observed spectrum. For each element, we use the
line list produced by SYNSPEC to identify the strongest features in the
bandpass of each instrument. We omit features that are likely to be blended
with geocoronal emission or photospheric, interstellar, or telluric absorption
lines. A few exceptions are described below. Our final list of absorption
features is presented in Tables 5 and 6, found in the Appendix.

Because we lack a model atmosphere with our best-fit stellar parameters, we
take the four Castelli-Kurucz models that are closest in parameter space
($T_{\rm eff}$= 10,750 K and 11,000 K; $\log g$= 2.0 and 2.5), generate
synthetic spectra for each, and interpolate among them to our best-fit values.
We adjust the photospheric abundances within SYNSPEC. We repeat the process to
generate a grid of spectra with a range of abundance values for the element of
interest.

At FUV wavelengths, our synthetic spectra suffer a serious flaw. As discussed
by Dixon & Chayer (2013), the Bright Star’s spectrum exhibits broad absorption
features due to resonances in the photoionization cross section of excited-
state neutral nitrogen. The Lanz & Hubeny (2007) atomic models include these
resonances, but smoothed over many tens of Ångstroms. The resulting synthetic
spectra exhibit strange, sloping continua. Fortunately, they are easily
repaired: For each synthetic spectrum, SYNSPEC writes a continuum model to a
separate file. We divide by this continuum, then multiply by the continuum of
a spectrum generated with no metals—and thus no resonant absorption.

Figure 3: Helium lines in the Bright Star’s optical spectrum. Black curves
represent the data, red curves our best-fit models. Lines appearing in
multiple orders are fit separately.

Continuum placement is the dominant uncertainty in our fits to the FUSE and
COS spectra. The FUV spectra of B-type stars are riddled with absorption
lines. Particularly in the FUSE bandpass, the apparent continuum lies far
below the true continuum level. After considerable experimentation, we
developed the following scheme to estimate the FUV continuum. We scale the
synthetic spectrum by a Fitzpatrick (1999) extinction curve assuming $E(B-V)$
= 0.04 and an ISM absorption model assuming an H I column density log $N$(H I)
= 21.90 [cm-2]. (Our choice of H I column density is discussed in Section
4.2.) Finally, we scale the model to reproduce the continuum in the region
around 1500 Å, where the S/N is highest.

As discussed below, we determine the iron abundance from the star’s FUV
spectrum by dividing it into 10 Å pieces, masking the lines not due to iron,
and fitting model spectra to each piece. Considering only the well-behaved
long-wavelength segment of the G160M spectrum, we derive an iron abundance
that is consistent with that derived from the optical. By inverting this
process, we can test the accuracy of our continuum model. Continuing to
shorter wavelengths, we determine the factor by which our model spectrum must
be scaled to yield the optically-determined iron abundance. We find that this
scale factor is within a few percent of unity down to about 1180 Å, at which
point our model continuum becomes less reliable. When possible, we avoid using
lines with wavelengths shorter than 1180 Å. When we must use such lines, we
multiply our models by a scale factor determined from nearby iron features.

We estimate an uncertainty of 5% in our FUV continuum placement. To see how
this uncertainty propagates into our abundance estimates, we perform each fit
twice, once with the model spectrum scaled as described and again with the
model scaled by an additional factor of 0.95. The difference in the two
abundances is an estimate of the systematic error due to uncertainties in the
continuum level. We add this term and the statistical error in quadrature to
compute the uncertainty reported for each feature. At optical wavelengths, we
use only the statistical error.

We fit each line or group of lines separately. We measure the equivalent width
for the optical lines, but not for the FUV features, which are often blended
and for which the continuum level is uncertain. Because the scale factors of
both the data and the synthetic spectra are fixed, the only free parameter in
our fits is the abundance of the element in question. For each element, the
final abundance and its uncertainly represent the mean and the standard
deviation of our individual measurements. Results are presented in Table 3 and
plotted in Fig. 4. Notes on individual elements follow.

Carbon: The Bright Star’s FUV spectrum exhibits absorption from both C I and C
II. The C I lines yield an abundance $\log N({\rm C})/N({\rm H})$ = $-6.16\pm
0.15$, a value more than 1.5 dex below the cluster mean. The star’s lone C II
line yields $\log N({\rm C})/N({\rm H})$ = $-5.26\pm 0.13$. Alexeeva et al.
(2016) compute non-LTE models of stars with effective temperatures ranging
from 9550 to 17,500 K, though with surface gravities and metallicities higher
than those of the Bright Star. In each of their models, C II dominates the
total carbon abundance throughout the line-formation region of the atmosphere.
The fraction of neutral carbon does not exceed a few tenths of a percent, and
in the line-formation layers, it is smaller in non-LTE than in LTE models due
to ultraviolet (UV) overionization. We therefore adopt the carbon abundance
derived from the C II line as the stellar value.

Nitrogen: The Bright Star’s spectrum exhibits numerous N I features. The
optical lines yield $\log N({\rm N})/N({\rm H})$ = $-3.48\pm 0.11$, and the
FUV lines yield $\log N({\rm N})/N({\rm H})$ = $-3.78\pm 0.40$. The spectrum
contains three strong N II lines, one in the FUV and two in the optical; they
yield $\log N({\rm N})/N({\rm H})$ = $-3.44\pm 0.01$. As these values are
consistent within the uncertainties, we use the mean of all 20 lines as the
stellar value.

Przybilla & Butler (2001) compare LTE and non-LTE models for several late-B
and early-A type supergiants. For the two stars whose parameters are most
similar to the Bright Star, they find that LTE models over-predict the
nitrogen abundance. For $\eta$ Leo ($T_{\rm eff}$ = 9600 K, $\log g$ = 2.00),
the N I lines modeled in LTE overestimate the abundance by 0.36 dex on average
and the N II lines by 0.15 dex. For $\beta$ Ori ($T_{\rm eff}$ = 12,000 K,
$\log g$ = 1.75), the corresponding values are 0.58 dex for N I and 0.27 dex
for N II. Thus, our nitrogen abundance may be slightly overestimated.

As discussed in Section 2.4, the Bright Star lies near other, redder stars
with which its spectrum may be blended at long wavelengths. To determine
whether our MIKE spectrum is contaminated, we consider the nitrogen abundance
derived from lines at the longest wavelengths. Nitrogen is useful for this
test, because its abundance is greatly enhanced relative to the cluster mean.
All of the optical N I features have wavelengths longer than 7400 Å. As we
have seen, they yield an abundance consistent with that of the FUV features,
indicating that contamination of the MIKE spectrum is insignificant.

To fit the nitrogen lines with wavelengths between 1165 and 1178 Å, we rescale
our model spectra by a factor of 0.93, a value determined from fits to nearby
iron lines.

Oxygen: We adopt the mean abundance derived from all of our O I lines as the
stellar value. Franchini et al. (2021) point out that, among the optical O I
features, the high-excitation 6158 Å line is less susceptible to non-LTE
effects than the O I triplet at 7772, 7774, and 7775 Å. The 6158 Å line yields
an abundance $\log N({\rm O})/N({\rm H})$ = $-3.97\pm 0.03$, slightly lower
than our adopted value.

To fit the O I $\lambda 1152$ features, we rescale our model spectra by a
factor of 0.86, again determined from fits to nearby iron lines.

Aluminum: The strongest aluminum feature in the Bright Star’s FUV spectrum is
the Al II line at 1670 Å. Being a resonance transition, it is also a prominent
ISM feature. We fit a Gaussian profile to the ISM line in the G160M spectrum
of UIT 14. Including this ISM line in our fit to the Al II $\lambda 1670$ line
in the Bright Star spectrum yields $\log N({\rm Al})/N({\rm H})$ $=-6.20\pm
0.09$. Ignoring the ISM contribution raises the derived abundance to $-6.17$,
a change smaller than the uncertainty in our fit. We adopt the mean value from
all Al II features as the stellar abundance.

Phosphorus: To fit the P II features near 1150 Å, we rescale our model spectra
by a factor of 0.87.

Sulfur: All of the available sulfur lines in the star’s FUV spectrum are due
to S I; they yield a mean abundance $\log N({\rm S})/N({\rm H})$ $=-6.11\pm
0.14$, a full dex lower than the cluster mean. The optical lines are due to S
II and yield a mean abundance $\log N({\rm S})/N({\rm H})$ $=-5.25\pm 0.22$.
We attribute this discrepancy to non-LTE effects. At the effective temperature
of the Bright Star, S II is the dominant species; we thus adopt the optically-
derived abundance for this element.

Chlorine: Our model spectra predict numerous chlorine features in the star’s
FUV spectrum, but all are hopelessly blended with other photospheric lines
except for a pair of Cl II lines at 1075 and 1079 Å. Of these, the 1075 Å
feature falls in a region where the continuum is depressed by a broad bound-
free transition of N I (Fig. 3 of Dixon & Chayer 2013) that is not included in
our model. Our chlorine abundance is thus derived from a single Cl II line at
1079 Å. For this feature, we rescale our model spectra by a factor of 1.13, a
value determined from fits to nearby iron lines.

Calcium: Only the Ca II $\lambda 1288$ feature yields an abundance consistent
with the cluster mean. The line falls at the edge of a COS detector segment,
where the S/N ratio is low and the wavelength scale uncertain. The other two
FUV lines and all three of the optical features, also due to Ca II, yield
consistent and much lower abundances. Again, non-LTE effects may be
responsible. At the effective temperature of the Bright Star, Ca III is the
dominant species (Sitnova et al., 2018). Unfortunately, there are no Ca III
features in our spectra, so we quote the mean of our Ca II results.

Titanium: The two ionization states of titanium yield different abundances.
The optical lines are due to Ti II and yield $\log N({\rm Ti})/N({\rm H})$
$=-8.43\pm 0.19$, while the FUV lines are due to Ti III and yield $\log N({\rm
Ti})/N({\rm H})$ $=-7.86\pm 0.18$. The cluster abundance is $\log N({\rm
Ti})/N({\rm H})$ = $-7.50\pm 0.07$ (Kovalev et al., 2019). Sitnova et al.
(2016) compute non-LTE models of stars with $T_{\rm eff}$ between 5780 and
12800 K. They find that non-LTE effects for Ti II are small for the coolest
models, but grow toward higher $T_{\rm eff}$ and lower $\log g$ and [Fe/H]. In
their hottest model, Ti III is the dominant species. We adopt the Ti III
abundance for this element.

Iron: We fit 17 individual Fe II lines in the MIKE spectrum; they yield a mean
abundance $\log N({\rm Fe})/N({\rm H})$ = $-5.46\pm 0.18$. There are hundreds
of iron lines in the FUV spectrum. Instead of fitting individual features, we
divide the long-wavelength segment of the G160M spectrum into 10 Å pieces. We
mask ISM features and all stellar features that are not due to iron and fit
our model spectra to each piece. We omit a few pieces that contain broad
absorption features or lack unblended iron lines. The 14 remaining pieces,
spanning the wavelength range 1610–1760 Å, yield a mean iron abundance $\log
N({\rm Fe})/N({\rm H})$ = $-5.43\pm 0.19$, consistent with our optically-
determined value.

Nickel: The large error bar on our nickel abundance, 0.51 dex, reflects the
wide range of values derived from individual features. Several features, Ni II
$\lambda 1501,\lambda 1502,\lambda 1536.9,$ and $\lambda 3849$, yield values
consistent with the cluster mean, but the rest are much lower. We adopt the
mean of all available measurements as the stellar value.

Gallium: Nielsen et al. (2005) fit LTE models to the Ga II and Ga III features
in a high-resolution STIS spectrum of the star $\chi$ Lup ($T_{\rm eff}$ =
10,650 K, $\log g$ = 3.8). They found that the ground-state resonance lines Ga
II $\lambda 1414$, Ga III $\lambda 1495$, and Ga III $\lambda 1534$ yield
abundances much greater than the lines from excited states of Ga II. In our
data, Ga II $\lambda 1414$ yields an abundance almost 2.0 dex less than the
excited-state transitions, while Ga III $\lambda 1495$ and $\lambda 1534$
yield abundances consistent with them. Castelli et al. (2017) report that Ga
II is the dominant species in the photosphere of HR 6000 ($T_{\rm eff}$ =
13,450 K, $\log g$ = 4.3), so we assume that the same is true for the Bright
Star. We thus use only the excited-state transitions of Ga II to derive the
abundance of this species. Note: Our original line list included only the Ga
II $\lambda 1414$ feature. We employ the line parameters published by Nielsen
et al. to model the gallium features in our spectrum.

Table 3: Photospheric Abundances

Species | Bright Star | 47 Tuc | Sun
---|---|---|---
He | $-0.82\pm 0.16$ |  $\cdots$ | $-1.07\pm 0.01$
C | $-5.26\pm 0.13$ | $-4.51\pm 0.24$ | $-3.57\pm 0.05$
N | $-3.60\pm 0.31$ | $-3.92\pm 0.41$ | $-4.17\pm 0.05$
O | $-3.76\pm 0.15$ | $-3.49\pm 0.18$ | $-3.31\pm 0.05$
Ne | $-3.92\pm 0.08$ |  $\cdots$ | $-4.07\pm 0.10$
Mg | $-4.99\pm 0.21$ | $-4.78\pm 0.05$ | $-4.40\pm 0.04$
Al | $-6.16\pm 0.28$ | $-6.12\pm 0.11$ | $-5.55\pm 0.03$
Si | $-5.05\pm 0.27$ | $-4.81\pm 0.06$ | $-4.49\pm 0.03$
P | $-7.21\pm 0.10$ |  $\cdots$ | $-6.59\pm 0.03$
S | $-5.25\pm 0.22$ | $-5.19\pm 0.14$ | $-4.88\pm 0.03$
Cl | $-7.13\pm 0.35$ |  $\cdots$ | $-6.50\pm 0.30$
Ca | $-6.96\pm 0.34$ | $-6.10\pm 0.13$ | $-5.66\pm 0.04$
Sc | $-9.76\pm 0.39$ | $-9.42\pm 0.06$ | $-8.85\pm 0.04$
Ti | $-7.86\pm 0.18$ | $-7.50\pm 0.07$ | $-7.05\pm 0.05$
Cr | $-7.59\pm 0.29$ | $-7.17\pm 0.08$ | $-6.36\pm 0.04$
Mn | $-7.65\pm 0.30$ | $-7.55\pm 0.08$ | $-6.57\pm 0.04$
Fe | $-5.44\pm 0.19$ | $-5.24\pm 0.03$ | $-4.50\pm 0.04$
Co | $-8.01\pm 0.09$ | $-7.79\pm 0.10$ | $-7.01\pm 0.07$
Ni | $-7.10\pm 0.51$ | $-6.68\pm 0.03$ | $-5.78\pm 0.04$
Cu | $-8.45\pm 0.38$ | $-8.73\pm 0.18$ | $-7.81\pm 0.04$
Ga | $-8.59\pm 0.36$ |  $\cdots$ | $-8.96\pm 0.09$
Pd | $-10.00\pm 0.10$ |  $\cdots$ | $-10.43\pm 0.10$
In | $-11.31\pm 0.43$ |  $\cdots$ | $-11.20\pm 0.20$
Sn | $-9.88\pm 0.43$ |  $\cdots$ | $-9.96\pm 0.10$
Hg | $-10.73\pm 0.39$ |  $\cdots$ | $-10.83\pm 0.08$
Pb | $-10.48\pm 0.53$ |  $\cdots$ | $-10.25\pm 0.10$

Note. — Abundances relative to hydrogen: $\log N({\rm X})/N({\rm H})$. LTE
cluster values for C, N, O, Si, and Ca from Mészáros et al. (2020) and for Sc,
Cr, Mn, Co, Ni, and Cu from Thygesen et al. (2014). Non-LTE cluster values for
Al from Thygesen et al., for S from Duffau et al. (2017), and for Mg, Ti, and
Fe from Kovalev et al. (2019).

Figure 4: Photospheric abundances of the Bright Star (stars), 47 Tuc (red
lines), and the solar photosphere (black lines). See Table 3 for references.

##  4 Details of the Models

###  4.1 Balmer Lines

Our models cannot reproduce the low-order hydrogen Balmer line profiles. The
model features are broad and shallow, while the observed features have narrow
wings and deep cores. This discrepancy reflects the fact that the low-order
lines are formed farther out in the stellar atmosphere, where the low particle
density facilitates photoionization, while the high-order lines are formed
deeper in the atmosphere, where collision processes dominate and the
assumption of LTE is more appropriate.

In Fig. 5, we plot the H$\alpha$ profile of the Bright Star, together with
three synthetic spectra. The green curve, reproduced from Fig. 2, is our best-
fit LTE model. The orange curve represents a simple, non-LTE model computed
using the program TLUSTY (Hubeny & Lanz, 1995). The model treats H and He in
non-LTE, treats CNO in LTE, and neglects heavier elements. Metal lines are
included in the synthetic spectrum assuming our best-fit abundances. We see
that this model spectrum reproduces the deep, narrow core of the H$\alpha$
profile, but over-predicts the wings of the line. The blue curve represents a
model from the BSTAR grid of Lanz & Hubeny (2007), a set of metal line-
blanketed, non-LTE, plane-parallel, hydrostatic model atmospheres with
parameters appropriate for B-type stars. Unfortunately, the coolest BSTAR
models have effective temperatures $T_{\rm eff}$ = 15,000 K, considerably
hotter than the Bright Star. We have selected the model with $T_{\rm eff}$ =
15,000 K and $\log g$ = 2.25, parameters closest to those of our target. This
model does an excellent job of reproducing the general shape of the line, but
the line core is too shallow, because the model is too hot. These results
suggest that a line-blanketed, non-LTE model with the appropriate stellar
parameters could reproduce the Bright Star’s H$\alpha$ profile.

Figure 5: The H$\alpha$ line profile. The black spectrum represents the data,
and colored spectra are various models, as described in the text.

###  4.2 Lyman Lines

The Bright Star’s Ly$\alpha$ profile is remarkably broad. To reproduce it, we
employ an ISM absorption model with an H I column density log $N$(H I) = 21.90
[cm-2]. The actual column density toward 47 Tuc is log $N$(H I) = 20.62
(Heiles & Cleary, 1979). To confirm this value, we fit the Ly$\alpha$ profile
of the star UIT 14 using a model from the OSTAR grid of Lanz & Hubeny (2003)
with $T_{\rm eff}$ = 55,000 K, $\log g$ = 4.75, and [Fe/H] = $-0.7$. Using a
Fitzpatrick (1999) extinction curve with $E(B-V)$ = 0.04 and an H I absorption
model with log $N$(H I) = 20.62, we are able to reproduce the observed line
profile. Why then must we assume a higher H I column density to reproduce the
Bright Star’s Ly$\alpha$ line?

We can learn something from Fig. 6, where the observed Ly$\alpha$ profile is
plotted, along with three model spectra. The green spectrum is our best-fit
model, which employs the high H I column. It provides a reasonable fit to the
data, despite its ad hoc construction. The blue spectrum is the same model,
but assuming the correct H I column. Its Ly$\alpha$ feature is far too narrow.
The orange spectrum is derived from the H-He non-LTE model discussed above,
again assuming the correct H I column. Its Ly$\alpha$ profile is broader than
that of the LTE model but is still too narrow. We suspect that additional non-
LTE effects broaden the Ly$\alpha$ line. For example, all of the models over-
predict the strength of the Si II lines at 1190, 1193, 1194, and 1197 Å and
under-predict the strength of the Si III line at 1206 Å, suggesting that these
transitions suffer from non-LTE effects.

Of the higher-order Lyman lines, only the region around Ly$\beta$ has
sufficient flux that we can study its shape. Fig. 7 shows the Ly$\beta$ line
profile. The data are plotted in black, and the blue and orange curves
represent our best-fit LTE model with the actual and enhanced H I columns,
respectively. The stellar Ly$\beta$ line is so much broader than the ISM
feature that the blue and orange curves are essentially identical. As shown by
Dixon & Chayer (2013), the broad troughs in the stellar continuum shortward of
Ly$\beta$ are sculpted by resonances in the bound-free absorption cross
section of the first excited state of N I. The broad absorption feature
present in the synthetic spectrum near 1045 Å is an absorption resonance in
the second excited state of N I. Both are stronger in our synthetic spectrum
than in the data. Apparently, our LTE model over-predicts the fraction of
neutral nitrogen in these excited states. Even so, our synthetic spectrum
reasonably reproduces the Ly$\beta$ line profile.

A more recent measure of the H I column density toward 47 Tuc is provided by
the H I 4-PI Survey (HI4PI Collaboration et al., 2016), which reports that log
$N$(H I) = 20.75. This column density yields a Ly$\alpha$ profile slightly
wider than that in the spectrum of UIT 14, so we use the lower value of Heiles
& Cleary (1979).

Figure 6: The Ly$\alpha$ line profile. The black spectrum represents the data,
and colored spectra are various models, as described in the text. Both data
and models are binned by 7 pixels, the approximate width of a resolution
element. The models are scaled to have the same flux at 1500 Å. Figure 7: Same
as Fig. 6, but for Ly$\beta$.

###  4.3 Helium Lines

The He I lines used to derive the star’s helium abundance are listed in Table
2. Strangely, many of these features did not appear in our initial set of
synthetic spectra. The lines were not weaker than expected; they were missing.
SYNSPEC produces an output file that lists all of the lines included in the
spectrum, and the He I features did not appear in that file. After some
experimentation, we found that, by adjusting the parameter IDSTD, we were able
to restore the missing helium lines. IDSTD is the index of the standard depth,
defined as the depth at which $T\approx$ (2/3)$T_{\rm eff}$. For a model with
70 depth points, the nominal value of IDSTD is 50. Raising IDSTD from 50 to 55
restored the helium features.

Table 4: Derived Parameters  Parameter | Value
---|---
$T_{\rm eff}$ (K) | $10,850\pm 250$
$\log g$ [cm s-2] | $2.20\pm 0.13$
$R_{*}/R_{\sun}$ | $9.63\pm 0.13$
$M_{*}/M_{\sun}$ | $0.54\pm 0.16$
$\log L_{*}/L_{\sun}$ | $3.06\pm 0.04$

##  5 Discussion

###  5.1 Cluster Membership

Our use of the Bright Star to explore theories of stellar evolution and the
chemistry of 47 Tuc is predicated on the assumption that the star is a cluster
member. That membership has been questioned in the literature. By fitting
Kurucz (1979) models to an IUE spectrum of the star, Thejll (1992) derived a
surface gravity ($\log g$ = 3.8–4.2, depending on the metallicity) that
implies a stellar mass between 71 and 200 $M_{\sun}$ if the star lies at the
distance of the cluster. He concluded that the star is a Population I
foreground star.

Dixon et al. (1995) derived a lower gravity, $\log g$ = 2.0 from the star’s
H$\beta$ Balmer line and $\log g$ = 2.5 from its HUT spectrum, corresponding
to a mass between 0.33 and 1.04 $M_{\sun}$. They also pointed out that, while
IUE spectra sample wavelengths longer than than $\sim$ 1150 Å, HUT spectra
extend to the Lyman limit. In the region between Lyman $\beta$ and Lyman
$\alpha$, solar-metallicity models under-predict the observed flux. The
authors concluded that the star does not have Pop I abundances.

We have measured the star’s radial velocity using 20 absorption features
scattered across the COS spectrum. We find a mean heliocentric velocity of
$-14.3\pm 2.8$ km s-1. The accuracy of the COS wavelength solution is about
7.5 km s-1 (Plesha et al., 2019), so the total uncertainty is about 8 km s-1.
The cluster has a mean heliocentric velocity of $-18.0$ km s-1, with a central
velocity dispersion of 11.0 km s-1 (Harris, 1996, 2010). The star’s radial
velocity is thus consistent with cluster membership. Finally, inspection of
Table 3 and Fig. 4 confirms that the Bright Star’s chemical abundances are
consistent with those of 47 Tuc. We conclude that the star is indeed a cluster
member.

###  5.2 Stellar Mass and Luminosity

We can derive a star’s radius, and from this its mass and luminosity, by
comparing its observed and predicted fluxes. The spectral irradiance of the
Bright Star, as reported in Gaia Early Data Release 3 (EDR3), is $G=10.952\pm
0.004$ mag (Gaia Collaboration et al., 2016, 2021). We generate a synthetic
spectrum using our best-fit Castelli-Kurucz model, scaled by a Fitzpatrick
(1999) extinction curve with $R_{V}=3.1$ and $E(B-V)$ = 0.04 (Harris, 1996,
2010). We compute synthetic stellar magnitudes using the recipe provided in
Riello et al. (2021). The ratio between the observed and model fluxes is
$\phi=(2.89\pm 0.01)\times 10^{-20}$.

In the synthetic spectra generated by SYNSPEC, the flux is expressed in terms
of the flux moment, $H_{\lambda}$. If the star’s radius and distance are
known, then the scale factor required to convert the model spectrum to the
flux at earth is $\phi=4\pi(R_{*}/d)^{2}$ (Kurucz, 1979). Using Gaia EDR3
data, Maíz Apellániz et al. (2021) derive a distance to 47 Tuc of $4.53\pm
0.06$ kpc, in excellent agreement with the eclipsing-binary result of Thompson
et al. (2020). Adopting this distance and our scale factor, we derive a
stellar radius $R_{*}/R_{\sun}$ of $9.63\pm 0.13$. Applying our adopted
surface gravity ($\log g$ = 2.20), we find that the stellar mass
$M_{*}/M_{\sun}$ is $0.54\pm 0.16$. Finally, combining the stellar radius with
our best-fit effective temperature ($T_{\rm eff}$ = 10,850 K), we derive a
stellar luminosity $\log L_{*}/L_{\sun}$ of $3.06\pm 0.04$. Dixon et al.
(1995) derived a similar value, $\log L_{*}/L_{\sun}$ of $3.07\pm 0.20$, from
fits to the star’s HUT spectrum.

Despite the star’s high luminosity, we have identified no wind features in its
spectrum.

Figure 8: Evolutionary tracks for stars similar to those of 47 Tuc during and
after the horizontal-branch stage. Darker colors represent more massive
objects. The Bright Star is indicated by the black cross. Blue parts of the
tracks indicate the regions where energy generation is dominated by H-burning,
while red parts indicate the regions where He-burning dominates.

###  5.3 Evolutionary Status

To understand the evolutionary history and current status of the Bright Star
in 47 Tuc, we have computed a small grid of stellar-evolution models
specifically tailored to this cluster. Following the work of Denissenkov et
al. (2017), models were computed assuming [Fe/H] = $-0.7$, [O/Fe] = 0.4, and a
helium abundance $Y=0.27$. Abundance ratios of metals are taken from Salasnich
et al. (2000). This leads to an initial composition of
$(X,Y,Z)=(0.72328,0.27,0.00676)$. The initial mass of the model is taken as
$M_{\rm ZAMS}$ = 0.87 $M_{\sun}$, which corresponds to a total age of 12.76
Gyr for the Bright Star, consistent with the age of the cluster. Microphysics
and macrophysics prescriptions are taken as in Miller Bertolami (2016), with
the obvious exception of the Rosseland opacities, which are chosen to reflect
the $\alpha$-enhanced abundances of the cluster. Winds on the RGB are adjusted
to populate the zero-age horizontal branch (ZAHB) as in Dixon et al. (2017,
2019).

The HB and post-HB evolution of the stellar models is illustrated in Fig. 8.
Masses of the tracks are, from red to blue on the HB ($M_{\rm ZAHB},M_{\rm
pAGB}$) = (0.87 $M_{\sun}$, 0.554 $M_{\sun}$), (0.80 $M_{\sun}$, 0.547
$M_{\sun}$), (0.75 $M_{\sun}$, 0.540 $M_{\sun}$), (0.70 $M_{\sun}$, 0.533
$M_{\sun}$), (0.65 $M_{\sun}$, 0.526 $M_{\sun}$), (0.625 $M_{\sun}$, 0.522
$M_{\sun}$), (0.60 $M_{\sun}$, 0.517 $M_{\sun}$), (0.575 $M_{\sun}$, 0.512
$M_{\sun}$), (0.55 $M_{\sun}$, 0.505 $M_{\sun}$), (0.525 $M_{\sun}$, 0.496
$M_{\sun}$), and (0.50 $M_{\sun}$, 0.494 $M_{\sun}$). The ZAHB of 47 Tuc
corresponds to tracks with masses between 0.65 and 0.80 $M_{\sun}$ ($T_{\rm
eff}$ between 5100 K and 5400 K). Darker colors show the typical evolution of
post-red horizontal branch (RHB) sequences ($M\gtrsim$ 0.60 $M_{\sun}$), while
lighter tracks show the typical evolution of post-blue and extreme horizontal
branch sequences (BHB and EHB).

To these evolutionary tracks we have added a sample of cluster stars from the
catalog of Stetson et al. (2019). Effective temperatures are estimated from
$B-V$ colors, and luminosities from $V$ magnitudes, using the color-color
relations and bolometric corrections of Worthey & Lee (2011). The cluster
distance is taken from Maíz Apellániz et al. (2021) and the reddening from
Harris (1996, 2010).

47 Tuc hosts a short, red horizontal branch and is essentially devoid of both
BHB and EHB stars. The Bright Star is thus a decendant of the RHB. Colors in
Fig. 8 indicate whether nuclear-energy generation is dominated by H-burning
(blue) or by He-burning (red). Whether the model departs the AGB as a H-burner
or a He-burner depends on the timing of the last thermal pulse. Models
departing the AGB during the quiescent H-burning burning phase (between
thermal pulses) will evolve as H-burners, while models departing during or
immediately after a thermal pulse will evolve as He-burners.

The timing of the last thermal pulse also affects the exact shape of the loops
traced by the star during its post-AGB evolution, as that shape is quite
sensitive to the remaining envelope mass. Given the uncertainties in the mass-
loss prescriptions and the various processes that may contribute to the actual
wind intensity of the AGB, the phase of the thermal pulse cycle in which the
star departs from the AGB can be considered effectively stochastic, and a star
of a given initial mass can depart at any phase. As a consequence, the red
loops in Fig. 8 should be considered representative of the possible
trajectories of He-burners within a given mass range rather than the exact
evolution a star with a given mass.

It is clear from Fig. 8 that the Bright Star is less bright (by about 0.3 dex)
than is predicted for post-AGB H-burners (dark-blue tracks). In fact, Bright
Star’s location in the HR diagram suggests that it is a He-burner.
Interestingly, the sequences that evolve through the observed HB of 47 Tuc
($0.65M_{\odot}\lesssim M_{\rm ZAHB}\lesssim 0.80M_{\odot}$) have post-AGB
masses ($0.526M_{\odot}\lesssim M_{\rm pAGB}\lesssim 0.547M_{\odot}$) that are
consistent with the mass of the Bright Star ($0.54\pm 0.16M_{\odot}$). It is
unlikely to be a post-RGB star, as it shows no sign of a binary companion, and
post-RGB stars with the luminosity of the Bright Star would have significantly
lower masses ($\sim 0.4M_{\odot}$).

The integrated mass loss between the ZAHB and the post-AGB phase in an old
cluster with $\alpha$-enhanced composition is of special interest, as it can
shed light on the initial-final mass function of these populations and on the
dependence of cold winds on chemical composition. Reliable determinations of
the final stellar masses in globular clusters (via spectroscopy) are difficult
and consequently rare (e.g., Moehler et al., 2004; Kalirai et al., 2009), and
no previous determination is available for 47 Tuc.111From fits to the cooling
sequence of 47 Tuc, Campos et al. (2016) estimate that young white dwarfs in
47 Tuc have masses around $0.52M_{\odot}$. In this context, the mass of the
Bright Star is highly valuable, as it effectively indicates the final mass of
single stars that are now evolving away from the main sequence. The main-
sequence turn-off mass for 47 Tuc, as inferred from theoretical isochrones, is
$\sim 0.9M_{\odot}$, while the masses of stars on its horizontal branch range
between $\sim 0.65M_{\odot}$ and $\sim 0.73M_{\odot}$ (Salaris et al., 2016).
The mass derived for the Bright Star indicates that the stars in 47 Tuc lose
$\Delta M_{\rm AGB}\sim 0.1M_{\odot}$–$0.2M_{\odot}$ on the AGB, a value that
is only slightly lower than the mass that the same stars lose on the RGB,
$\Delta M_{\rm AGB}\sim 0.17$–$0.21M_{\odot}$ (Salaris et al., 2016).

Low-mass models such as those presented in this section do not experience
third dredge up on the AGB, and no surface enrichment is expected at this
stage. The N enhancement observed in the Bright Star might be a hint of non-
convective mixing processes, such as rotation-induced mixing or fingering
convection, during previous evolutionary stages (Charbonnel & Lagarde, 2010).

###  5.4 Photospheric Abundances

Galactic globular clusters host multiple stellar populations. First-generation
(FG) stars display abundances typical of halo field stars, while second-
generation (SG) stars, which may have multiple subpopulations, are enriched in
He, N, Na, and Al and depleted in C, O and Mg. Models suggest that the SG
stars were formed from gas polluted by material expelled by massive FG stars.
For details, see the review by Bastian & Lardo (2018).

Milone et al. (2012) found that FG stars constitute $\sim$ 30% of the stars in
47 Tuc. The majority of the cluster consists of SG stars, which are enhanced
in He and N and depleted in C and O. An additional subpopulation of SG stars,
making up about 8% of the cluster’s stars, is visible only on the subgiant
branch. To which generation does the Bright Star belong? 47 Tuc exhibits the
Na–O anti-correlation characteristic of all globular clusters. Several authors
use this relationship to define the cluster’s subpopulations (e.g., Cordero et
al., 2014; Thygesen et al., 2014). Since we do not know the Na abundance of
the Bright Star, we employ the anti-correlation of the cluster’s C and N
abundances.

Figure 9: Carbon versus nitrogen abundance for stars in 47 Tuc. Cluster data
are from Mészáros et al. 2020. The Bright Star is indicated by the black
symbol.

In Fig. 9, we plot the C and N abundances of stars in 47 Tuc, along with those
of the Bright Star. The cluster sample includes both RGB and AGB stars from
the full length of the giant branch (Mészáros et al., 2020). We see that the
Bright Star lies among the most nitrogen-rich, carbon-poor stars in the
cluster. This analysis is complicated by the fact that, within globular
clusters, C and N abundances are determined by two astrophysical processes,
pollution from FG stars and deep mixing, a non-convective mixing process that
results in a steady depletion of carbon and an enhancement of nitrogen as
stars evolve up the RGB (Gratton et al., 2004). If deep mixing were the
dominant effect, we would expect all of the cluster’s AGB stars—which
presumably evolved to the tip of the RGB—to be enhanced in N; instead, they
show the same range of N abundance as the RGB stars. We conclude that, while
the Bright Star’s C/N ratio may have been reduced as it climbed the RGB, it
began life as a SG star.

Returning to Fig. 4, we see that abundances of the intermediate-mass elements
(Mg through Ga) generally scale with Fe. Where cluster values are available,
all but copper are slightly higher than those measured for the Bright Star.
Though the difference is small for any single element, the uniformity of this
offset is striking. It suggests that systematic effects, whether the use of
different model atmospheres, absorption features, or temperature regimes, may
be at play.

The heaviest elements, Pd through Pb, have roughly solar abundances. These
elements are generally associated with s-process nucleosynthesis on the AGB;
however, the star’s low carbon abundance ($N_{\rm C}\ll N_{\rm O}$) indicates
that it did not undergo third dredge-up while on the AGB and implies that the
heaviest elements were made by a previous generation of stars. If so, then we
would expect this pattern to be present throughout the cluster, not just in
this star.

Roederer et al. (2010) derived the Pb abundance of a sample of metal-poor
dwarf and giant stars in the Galactic halo and disk by fitting LTE models to
their Pb I $\lambda 4057$ lines. The resulting abundances are subject to LTE
effects, because Pb is mostly ionized in stars with $T_{\rm eff}$ $>4000$ K.
Mashonkina et al. (2012) derived non-LTE corrections for these stars. In Fig.
10, we plot the corrected value of [Pb/H] as a function of [Fe/H] from
Mashonkina et al. Yong et al. (2006) measured the Pb abundance in five giant
stars in the globular cluster NGC 6752 and four in M13, again fitting LTE
models to the Pb I $\lambda 4057$ line. We apply the Mashonkina et al.
corrections to their results and plot them as well. Finally, we plot the
Bright Star. Because its Pb abundance was derived from a Pb II line, non-LTE
effects should be small. We see that the stars in NGC 6752, M13, and 47 Tuc
follow the same trend as the halo and disk stars.

Figure 10: Lead abundance versus metallicity for metal-poor dwarf and giant
stars in the Galactic halo and disk and giant stars in NGC 6752 and M13. The
Bright Star is indicated by the black symbol.

##  6 Conclusions

We have analyzed FUV and optical spectra of the Bright Star in 47 Tuc. Its
high-order Balmer lines, H06 and above, are well-fit by LTE stellar-atmosphere
models, but the low-order lines are not. These features are predicted to be
broad and shallow, while the observed features have narrow wings and deep
cores. Comparisons with metal line-blanketed non-LTE models at higher
effective temperatures suggest that similar models at lower temperatures could
reproduce the observed features.

The horizontal branch of 47 Tuc is short and red. Stellar-evolution models
that reproduce the cluster’s ZAHB have post-AGB masses consistent with that of
the Bright Star. The star’s relatively low luminosity indicates that it is now
a He-burner, having left the AGB during or immediately after a thermal pulse.
Its derived mass suggests that single stars in 47 Tuc lose $\Delta M_{\rm
AGB}\sim 0.1M_{\odot}$–$0.2M_{\odot}$ on the AGB, a value that is only
slightly lower than the mass the same stars lose on the RGB.

The star’s low C/N ratio suggests that it belongs to the second generation of
cluster stars, while its low C/O ratio indicates that it did not undergo third
dredge-up on the AGB. Abundances of intermediate-mass elements are consistent
with cluster values, while heavy-element abundances are essentially solar. We
conclude that the star experienced no significant change in its photospheric
abundances as it climbed the AGB. If so, then its heavy-element abundances are
typical of the cluster values. Thus, the Bright Star provides a unique look at
the abundances of heavy elements in 47 Tuc.

The authors would like to think Christian Johnson for helpful comments on the
manuscript. P.C. is supported by the Canadian Space Agency under a contract
with NRC Herzberg Astronomy and Astrophysics. This work has made use of NASA’s
Astrophysics Data System (ADS); the SIMBAD database, operated at CDS,
Strasbourg, France; the Mikulski Archive for Space Telescopes (MAST), hosted
at the Space Telescope Science Institute, which is operated by the Association
of Universities for Research in Astronomy, Inc., under NASA contract
NAS5-26555; data, software and/or web tools obtained from the High Energy
Astrophysics Science Archive Research Center (HEASARC), a service of the
Astrophysics Science Division at NASA/GSFC and of the Smithsonian
Astrophysical Observatory’s High Energy Astrophysics Division; and data from
the European Space Agency (ESA) mission Gaia
(https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and
Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement. Publication of this work is
supported by the STScI Director’s Discretionary Research Fund.

##  Appendix A Selected Features in the Bright Star’s Spectrum

Table 5: Selected FUV Absorption Features

Ion | $\lambda_{\rm lab}$ | $\log gf$ | $E_{l}$ | Abundance | Grating
---|---|---|---|---|---
| (Å) |  | (cm-1) |  |
C I | 1328.834 | -1.20 | 0.000 | $-5.90\pm 0.19$ | G130M |
| 1329.085 | -1.20 | 16.400 |  $\cdots$ |  |
| 1329.100 | -1.11 | 16.400 |  $\cdots$ |  |
| 1329.123 | -1.32 | 16.400 |  $\cdots$ |  |
| 1329.577 | -0.62 | 43.400 | $-6.22\pm 0.19$ |  |
| 1329.600 | -1.10 | 43.400 |  $\cdots$ |  |
| 1656.267 | -0.75 | 16.400 | $-6.35\pm 0.20$ | G160M |
| 1657.907 | -0.85 | 16.400 | $-6.13\pm 0.18$ |  |
| 1658.121 | -0.75 | 43.400 | $-6.21\pm 0.18$ |  |
C II | 1323.862 | -1.30 | 74930.100 | $-5.26\pm 0.13$ | G130M |
| 1323.906 | -0.34 | 74932.620 |  $\cdots$ |  |
| 1323.951 | -0.15 | 74930.100 |  $\cdots$ |  |
| 1323.995 | -1.30 | 74932.620 |  $\cdots$ |  |
N I | 1165.594 | -2.56 | 19224.465 | $-3.93\pm 0.21$ | G130M |
| 1165.717 | -3.99 | 19224.465 |  $\cdots$ |  |
| 1165.836 | -3.61 | 19233.178 |  $\cdots$ |  |
| 1169–1173aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-3.62\pm 0.14$ |  |
| 1176.510 | -1.15 | 19224.465 | $-3.41\pm 0.21$ |  |
| 1177.695 | -1.37 | 19233.177 | $-4.57\pm 0.34$ |  |
| 1243.171 | -1.51 | 19224.465 | $-3.48\pm 0.22$ |  |
| 1243.179 | -0.35 | 19224.465 |  $\cdots$ |  |
| 1243.306 | -0.54 | 19233.177 |  $\cdots$ |  |
| 1243.313 | -1.49 | 19233.177 |  $\cdots$ |  |
| 1318.823 | -2.55 | 28838.919 | $-3.81\pm 0.25$ |  |
| 1318.998 | -1.53 | 28838.919 |  $\cdots$ |  |
| 1319.005 | -1.84 | 28839.307 |  $\cdots$ |  |
| 1319.669 | -1.68 | 28838.919 | $-4.16\pm 0.34$ |  |
| 1319.676 | -1.18 | 28839.307 |  $\cdots$ |  |
| 1327.917 | -2.42 | 28838.919 | $-3.27\pm 0.36$ |  |
| 1327.924 | -2.72 | 28839.307 |  $\cdots$ |  |
N II | 1300.035 | -5.25 | 15316.200 | $-3.43\pm 0.37$ | G130M |
O I | 1152.151 | -0.28 | 15867.862 | $-3.60\pm 0.13$ | FUSE |
|  | $-3.61\pm 0.16$ | G130M
Mg II | 1307.875 | -2.49 | 35669.309 | $-4.58\pm 0.23$ | G130M |
| 1365.544 | -2.33 | 35669.309 | $-5.25\pm 0.18$ |  |
| 1367.708 | -2.24 | 35669.309 | $-5.27\pm 0.20$ |  |
| 1369.423 | -1.94 | 35760.880 | $-4.96\pm 0.29$ |  |
| 1398.776 | -4.83 | 0.000 | $-4.73\pm 0.23$ |  |
| 1476.000 | -1.84 | 35669.309 | $-4.91\pm 0.32$ | G160M |
| 1482.890 | -1.62 | 35760.880 | $-5.02\pm 0.35$ |  |
| 1734.852 | -1.10 | 35669.309 | $-5.16\pm 0.48$ |  |
| 1737.613 | -1.81 | 35760.880 | $-5.46\pm 0.37$ |  |
| 1737.628 | -0.85 | 35760.880 |  $\cdots$ |  |
| 1753.474 | -1.15 | 35760.880 | $-4.67\pm 0.40$ |  |
Al II | 1670.787 | 0.32 | 0.000 | $-6.20\pm 0.09$ | G160M |
Si II | 1264.738 | +0.55 | 287.240 | $-5.09\pm 0.05$ | G130M |
| 1265.002 | -0.41 | 287.240 |  $\cdots$ |  |
| 1309.276 | -0.12 | 287.240 | $-5.15\pm 0.12$ |  |
| 1309.453 | +0.55 | 55309.353 |  $\cdots$ |  |
| 1346.884 | -0.55 | 42932.622 | $-4.23\pm 0.28$ |  |
| 1348.543 | -0.57 | 42824.289 | $-4.90\pm 0.34$ |  |
| 1350.072 | -0.17 | 43107.910 | $-4.87\pm 0.30$ |  |
| 1350.516 | -1.07 | 42932.622 | $-4.91\pm 0.27$ |  |
| 1350.656 | -1.27 | 42824.289 |  $\cdots$ |  |
| 1353.721 | -0.53 | 43107.910 | $-4.94\pm 0.27$ |  |
| 1533.431 | -0.29 | 287.240 | $-5.12\pm 0.08$ | G160M |
| 1710.836 | -0.50 | 55309.353 | $-5.41\pm 0.34$ |  |
| 1711.304 | -0.20 | 55325.181 | $-5.45\pm 0.34$ |  |
Si III | 1296.726 | -0.13 | 52724.691 | $-4.85\pm 0.12$ | G130M |
| 1301.149 | -0.13 | 52853.283 | $-4.90\pm 0.11$ |  |
| 1417.237 | -0.11 | 82884.404 | $-4.93\pm 0.22$ |  |
P II | 1149.958 | -0.51 | 164.900 | $-7.34\pm 0.45$ | G130M |
| 1155.014 | -0.74 | 164.900 | $-7.30\pm 0.64$ |  |
| 1310.703 | -1.38 | 469.120 | $-7.10\pm 0.36$ |  |
| 1542.304 | -1.45 | 469.120 | $-7.23\pm 0.36$ | G160M |
| 1543.133 | -2.24 | 469.120 | $-7.09\pm 0.32$ |  |
S I | 1425.188 | -0.87 | 0.000 | $-6.22\pm 0.27$ | G130M |
| 1425.219 | -2.05 | 0.000 |  $\cdots$ |  |
|  | $-6.26\pm 0.28$ | G160M
| 1436.967 | -0.72 | 573.640 | $-6.00\pm 0.28$ |  |
| 1448.229 | -0.19 | 9238.609 | $-5.88\pm 0.25$ |  |
| 1473.994 | -0.43 | 0.000 | $-6.18\pm 0.26$ |  |
Cl II | 1079.080 | -1.62 | 696.000 | $-7.13\pm 0.35$ | FUSE |
Ca II | 1288.210 | -1.29 | 13710.880 | $-6.37\pm 0.32$ | G130M |
| 1553.176 | -0.35 | 13650.190 | $-6.95\pm 0.28$ | G160M |
| 1554.642 | -0.20 | 13710.880 | $-7.42\pm 0.28$ |  |
| 1554.642 | -1.50 | 13710.880 |  |  |
Sc III | 1610.194 | -0.47 | 0.000 | $-9.76\pm 0.39$ | G160M |
Ti III | 1293.225 | -0.75 | 420.400 | $-7.99\pm 0.42$ | G130M |
| 1339.703 | -1.67 | 8473.500 | $-7.51\pm 0.31$ |  |
| 1421.641 | -0.92 | 10603.600 | $-7.93\pm 0.25$ |  |
| 1421.755 | -1.01 | 10603.600 | $-7.75\pm 0.27$ |  |
| 1491.975 | -1.02 | 10721.200 | $-7.96\pm 0.30$ | G160M |
| 1498.695 | -0.33 | 8473.500 | $-7.95\pm 0.40$ |  |
| 1499.177 | -0.94 | 10721.200 | $-8.06\pm 0.20$ |  |
| 1504.623 | -1.39 | 10538.400 | $-7.94\pm 0.31$ |  |
| 1506.101 | -1.63 | 10603.600 | $-7.62\pm 0.36$ |  |
Cr II | 1427.382 | -0.59 | 12303.860 | $-7.80\pm 0.28$ | G130M |
|  | $-7.75\pm 0.30$ | G160M
| 1431.859 | -0.05 | 12303.860 | $-7.78\pm 0.36$ | G130M |
|  | $-7.86\pm 0.27$ | G160M
| 1435.202 | -0.27 | 12147.821 | $-7.82\pm 0.16$ | G160M |
| 1435.229 | -0.85 | 12032.580 |  $\cdots$ |  |
| 1435.251 | -1.31 | 12303.860 |  $\cdots$ |  |
Cr III | 1426.985 | -1.27 | 49768.646 | $-7.36\pm 0.50$ | G130M |
|  | $-7.06\pm 0.29$ | G160M
Mn II | 1163.326 | -1.30 | 0.000 | $-7.23\pm 0.39$ | G130M |
| 1464.570 | -0.39 | 14325.860 | $-8.06\pm 0.35$ | G160M |
| 1732.700 | -0.72 | 14593.819 | $-7.41\pm 0.30$ |  |
| 1740.144 | -0.81 | 14781.190 | $-7.87\pm 0.37$ |  |
| 1741.977 | -0.86 | 14901.180 | $-7.68\pm 0.43$ |  |
Fe II | 1610–1619aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.25\pm 0.29$ | G160M |
| 1620–1629aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.25\pm 0.31$ |  |
| 1630–1640aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.36\pm 0.50$ |  |
| 1640–1650aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.25\pm 0.34$ |  |
| 1650–1660aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.36\pm 0.34$ |  |
| 1660–1669aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.43\pm 0.49$ |  |
| 1672–1680aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.46\pm 0.38$ |  |
| 1680–1690aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.34\pm 0.68$ |  |
| 1690–1700aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.85\pm 0.41$ |  |
| 1701–1709aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.71\pm 0.26$ |  |
| 1710–1720aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.75\pm 0.33$ |  |
| 1720–1730aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.40\pm 0.39$ |  |
| 1730–1740aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.35\pm 0.44$ |  |
| 1755–1760aaSpectral region with multiple features. See discussion in text. |  $\cdots$ |  $\cdots$ 00 | $-5.23\pm 0.37$ |  |
Co II | 1448.011 | -0.39 | 0.000 | $-7.97\pm 0.38$ | G160M |
| 1455.885 | -0.39 | 950.510 | $-7.88\pm 0.40$ |  |
| 1466.203 | 0.10 | 0.000 | $-8.16\pm 0.51$ |  |
| 1472.902 | 0.01 | 950.510 | $-8.05\pm 0.41$ |  |
| 1476.666 | -0.51 | 1597.320 | $-7.97\pm 0.35$ |  |
Ni II | 1500.434 | -0.72 | 1506.940 | $-7.04\pm 0.54$ | G160M |
| 1501.959 | -2.34 | 0.000 | $-6.52\pm 0.34$ |  |
| 1502.148 | -1.10 | 0.000 | $-6.60\pm 0.41$ |  |
| 1536.741 | -1.35 | 1506.940 | $-7.30\pm 0.38$ |  |
| 1536.939 | -1.64 | 1506.940 | $-6.56\pm 0.46$ |  |
| 1703.405 | -1.13 | 0.000 | $-7.35\pm 0.41$ |  |
| 1748.282 | -0.25 | 1506.940 | $-8.16\pm 0.44$ |  |
| 1754.809 | -0.92 | 1506.940 | $-7.58\pm 0.47$ |  |
Cu II | 1358.773 | -0.27 | 0.000 | $-8.98\pm 0.50$ | G130M |
| 1367.951 | -1.48 | 0.000 | $-8.29\pm 0.40$ |  |
| 1472.395 | -2.08 | 0.000 | $-8.08\pm 0.38$ | G160M |
Ga II | 1414.402 | 0.35 | 0.000 | $-10.53\pm 0.31$ | G130M |
| 1473.690 | -0.24 | 47367.548 | $-8.49\pm 0.40$ | G160M |
| 1483.453 | -0.37 | 47814.115 | $-8.54\pm 0.47$ |  |
| 1483.903 | 0.33 | 48749.852 | $-9.19\pm 0.57$ |  |
| 1495.185 | -0.25 | 47814.115 | $-8.85\pm 0.64$ |  |
| 1504.331 | -0.16 | 48749.852 | $-8.57\pm 0.41$ |  |
| 1504.926 | -0.10 | 47367.548 | $-8.19\pm 0.31$ |  |
| 1515.106 | -0.22 | 47814.115 | $-8.03\pm 0.32$ |  |
| 1535.309 | 0.53 | 48749.852 | $-8.90\pm 0.33$ |  |
Ga III | 1495.045 | 0.03 | 0.000 | $-8.79\pm 0.40$ |  |
| 1534.462 | -0.28 | 0.000 | $-9.15\pm 0.55$ |  |
Pd II | 1345.556 | 0.00 | 0.000 | $-9.89\pm 0.34$ | G130M |
| 1363.745 | 0.00 | 0.000 | $-10.15\pm 0.32$ |  |
| 1397.448 | 0.00 | 3539.000 | $-9.93\pm 0.35$ |  |
| 1403.613 | 0.00 | 0.000 | $-10.04\pm 0.32$ |  |
In II | 1586.450 | 0.39 | 0.000 | $-11.31\pm 0.43$ | G160M |
Sn II | 1475.011 | 0.47 | 4251.494 | $-9.88\pm 0.43$ | G160M |
Hg II | 1649.947 | 0.22 | 0.000 | $-10.73\pm 0.39$ | G160M |
Pb II | 1433.905 | 0.24 | 0.000 | $-10.48\pm 0.53$ | G160M |

Note. — Abundances relative to hydrogen: $\log N({\rm X})/N({\rm H})$.

Table 6: Selected Optical Absorption Features

Ion | $\lambda_{\rm lab}$ | $\log gf$ | $E_{l}$ | EW (mÅ) | Abundance
---|---|---|---|---|---
| (Å) |  | (cm-1) |  |
N I | 7468.312 | -0.27 | 83364.615 | $43.5\pm 1.3$ | $-3.47\pm 0.02$ |
| 8216.336 | -0.01 | 83364.615 | $63.9\pm 1.7$ | $-3.30\pm 0.02$ |
| 8223.128 | -0.39 | 83317.826 | $38.7\pm 1.7$ | $-3.42\pm 0.02$ |
| 8680.282 | 0.24 | 83364.615 | $74.2\pm 2.5$ | $-3.71\pm 0.02$ |
| 8683.403 | -0.04 | 83317.826 | $67.1\pm 2.3$ | $-3.59\pm 0.02$ |
| 8686.149 | -0.45 | 83284.068 | $51.9\pm 2.4$ | $-3.43\pm 0.02$ |
| 8703.247 | -0.41 | 83284.068 | $54.4\pm 2.2$ | $-3.46\pm 0.02$ |
| 8711.703 | -0.38 | 83317.826 | $52.8\pm 2.0$ | $-3.51\pm 0.02$ |
| 8718.837 | -0.43 | 83364.615 | $51.6\pm 2.0$ | $-3.47\pm 0.02$ |
N II | 3994.997 | 0.28 | 149187.793 | $18.0\pm 2.0$ | $-3.43\pm 0.05$ |
| 4630.539 | 0.13 | 149076.512 | $9.7\pm 0.9$ | $-3.45\pm 0.05$ |
O I | 6156.778 | -0.73 | 86627.782 | $8.8\pm 1.4$ | $-3.97\pm 0.04$ |
| 6158.187 | -0.44 | 86631.453 | $12.7\pm 1.3$ | $-3.97\pm 0.03$ |
| 7771.944 | 0.32 | 73768.202 | $79.1\pm 1.5$ | $-3.87\pm 0.02$ |
| 7774.166 | 0.17 | 73768.202 | $75.6\pm 1.5$ | $-3.78\pm 0.02$ |
| 7775.388 | -0.05 | 73768.202 | $64.2\pm 1.5$ | $-3.76\pm 0.02$ |
| 7950.803 | 0.34 | 101147.522 | $10.1\pm 1.5$ | $-3.60\pm 0.04$ |
| 7952.159 | 0.17 | 101155.417 | $8.6\pm 1.6$ | $-3.53\pm 0.05$ |
| 7981.942 | -0.47 | 88630.585 | $13.1\pm 2.1$ | $-3.93\pm 0.04$ |
| 7982.398 | -0.34 | 88631.307 |  $\cdots$ 00 |  $\cdots$ |
| 8446.247 | -0.52 | 76794.977 | $148.1\pm 2.4$ | $-3.72\pm 0.01$ |
| 8446.359 | 0.17 | 76794.977 |  $\cdots$ 00 |  $\cdots$ |
| 8446.758 | -0.05 | 76794.977 |  $\cdots$ 00 |  $\cdots$ |
Ne I | 6096.163 | -0.27 | 134459.279 | $7.7\pm 1.6$ | $-4.04\pm 0.10$ |
| 6334.428 | -0.31 | 134041.849 | $9.8\pm 1.4$ | $-3.80\pm 0.06$ |
| 6382.991 | -0.26 | 134459.279 | $8.4\pm 1.3$ | $-3.92\pm 0.06$ |
| 6402.246 | 0.36 | 134041.849 | $21.1\pm 1.4$ | $-3.91\pm 0.04$ |
| 6506.528 | 0.03 | 134459.279 | $13.6\pm 1.3$ | $-3.92\pm 0.04$ |
Mg I | 3832.299 | -0.36 | 21870.465 | $12.8\pm 1.6$ | $-4.89\pm 0.05$ |
| 3832.304 | 0.12 | 21870.465 |  $\cdots$ 00 |  $\cdots$ |
| 3838.290 | -1.53 | 21911.178 | $17.9\pm 1.5$ | $-4.89\pm 0.04$ |
| 3838.292 | 0.39 | 21911.178 |  $\cdots$ 00 |  $\cdots$ |
| 3838.295 | -0.35 | 21911.178 |  $\cdots$ 00 |  $\cdots$ |
Mg II | 3848.211 | -1.59 | 71490.190 | $12.1\pm 2.3$ | $-5.04\pm 0.06$ |
| 3848.341 | -2.54 | 71491.064 |  $\cdots$ 00 |  $\cdots$ |
| 3850.386 | -1.84 | 71491.064 | $7.6\pm 2.6$ | $-4.98\pm 0.11$ |
| 4390.572 | -0.53 | 80650.022 | $32.1\pm 1.0$ | $-5.03\pm 0.01$ |
| 4433.988 | -0.91 | 80650.022 | $22.6\pm 1.1$ | $-4.79\pm 0.02$ |
| 4481.126 | 0.74 | 71490.190 | $232.6\pm 1.7$ | $-5.08\pm 0.01$ |
| 4481.150 | -0.56 | 71490.190 |  $\cdots$ 00 |  $\cdots$ |
| 4481.325 | 0.59 | 71491.064 |  $\cdots$ 00 |  $\cdots$ |
| 5401.521 | -0.45 | 93799.630 | $13.2\pm 1.8$ | $-5.14\pm 0.05$ |
| 5401.556 | -0.34 | 93799.750 |  $\cdots$ 00 |  $\cdots$ |
| 7877.054 | 0.39 | 80619.500 | $66.5\pm 1.7$ | $-4.86\pm 0.02$ |
| 7896.042 | -0.31 | 80650.022 | $95.1\pm 2.9$ | $-5.03\pm 0.02$ |
| 7896.366 | 0.65 | 80650.022 |  $\cdots$ 00 |  $\cdots$ |
Al II | 4663.046 | -0.28 | 85481.354 | $25.4\pm 1.0$ | $-6.06\pm 0.02$ |
| 5593.300 | 0.41 | 106920.564 | $3.4\pm 2.2$ | $-6.46\pm 0.26$ |
| 6226.195 | 0.05 | 105427.522 | $6.4\pm 1.4$ | $-5.80\pm 0.08$ |
| 6231.750 | 0.40 | 105441.498 | $3.8\pm 1.3$ | $-6.54\pm 0.15$ |
| 6243.367 | 0.67 | 105470.928 | $20.2\pm 1.2$ | $-5.89\pm 0.03$ |
| 7042.083 | 0.35 | 91274.503 | $21.8\pm 1.4$ | $-6.11\pm 0.04$ |
| 7056.712 | 0.13 | 91274.503 | $7.4\pm 1.4$ | $-6.53\pm 0.09$ |
| 7063.682 | -0.35 | 91274.503 | $10.8\pm 1.5$ | $-5.85\pm 0.06$ |
Si II | 3853.665 | -1.52 | 55309.353 | $69.6\pm 3.2$ | $-4.99\pm 0.03$ |
| 3856.018 | -0.56 | 55325.181 | $102.1\pm 1.9$ | $-5.25\pm 0.03$ |
| 3862.595 | -0.82 | 55309.353 | $95.9\pm 2.5$ | $-5.15\pm 0.03$ |
| 4128.054 | 0.32 | 79338.502 | $82.2\pm 1.2$ | $-5.20\pm 0.02$ |
| 4130.872 | -0.82 | 79355.025 | $84.0\pm 1.1$ | $-5.33\pm 0.02$ |
| 4130.894 | 0.48 | 79355.025 |  $\cdots$ 00 |  $\cdots$ |
| 5041.024 | 0.29 | 81191.346 | $100.1\pm 1.9$ | $-4.59\pm 0.02$ |
| 5055.984 | 0.59 | 81251.320 | $112.1\pm 2.5$ | $-5.22\pm 0.02$ |
| 5056.317 | -0.36 | 81251.320 |  $\cdots$ 00 |  $\cdots$ |
| 5957.559 | -0.30 | 81191.346 | $27.3\pm 2.1$ | $-5.29\pm 0.03$ |
| 5978.930 | 0.00 | 81251.320 | $42.6\pm 1.7$ | $-5.22\pm 0.03$ |
| 6371.371 | 0.00 | 65500.467 | $99.4\pm 1.2$ | $-5.15\pm 0.02$ |
S II | 4162.665 | 0.83 | 128599.152 | $10.9\pm 1.9$ | $-5.33\pm 0.08$ |
| 5212.620 | 0.66 | 121530.021 | $13.4\pm 1.1$ | $-5.01\pm 0.07$ |
| 5428.655 | -0.01 | 109560.686 | $11.4\pm 1.7$ | $-5.03\pm 0.08$ |
| 5453.855 | 0.56 | 110268.606 | $21.4\pm 1.5$ | $-5.05\pm 0.05$ |
| 5606.151 | 0.16 | 110766.562 | $6.8\pm 1.6$ | $-5.40\pm 0.11$ |
| 5639.977 | 0.33 | 113461.537 | $9.0\pm 1.0$ | $-5.28\pm 0.06$ |
| 5640.346 | 0.15 | 110508.706 | $4.6\pm 0.9$ | $-5.64\pm 0.10$ |
Ca II | 3736.902 | -0.15 | 25414.399 | $22.8\pm 3.3$ | $-7.09\pm 0.06$ |
| 3933.663 | 0.13 | 0.000 | $124.8\pm 1.6$ | $-7.21\pm 0.04$ |
| 3968.469 | -0.17 | 0.000 | $52.2\pm 1.5$ | $-6.74\pm 0.04$ |
Ti II | 3741.635 | -0.11 | 12758.110 | $12.4\pm 2.2$ | $-8.16\pm 0.08$ |
| 3759.296 | 0.20 | 4897.650 | $12.3\pm 2.3$ | $-8.57\pm 0.04$ |
| 3761.323 | 0.10 | 4628.580 | $23.1\pm 2.4$ | $-8.57\pm 0.04$ |
Cr II | 4242.364 | -0.59 | 31219.350 | $14.2\pm 1.5$ | $-7.90\pm 0.04$ |
| 4558.650 | -0.66 | 32854.308 | $24.6\pm 0.9$ | $-7.23\pm 0.01$ |
| 4588.199 | -0.63 | 32836.680 | $21.9\pm 1.0$ | $-7.33\pm 0.02$ |
Fe II | 3779.576 | -3.78 | 20516.961 | $3.7\pm 2.3$ | $-5.48\pm 0.15$ |
| 3781.509 | -2.77 | 36252.917 | $5.4\pm 2.2$ | $-5.30\pm 0.15$ |
| 3783.347 | -3.16 | 18360.647 | $16.1\pm 2.1$ | $-5.52\pm 0.04$ |
| 3845.180 | -2.29 | 36126.386 | $11.6\pm 2.2$ | $-5.45\pm 0.06$ |
| 4173.461 | -2.18 | 20830.582 | $34.3\pm 1.3$ | $-5.93\pm 0.02$ |
| 4177.692 | -3.75 | 20516.961 | $9.3\pm 1.2$ | $-5.12\pm 0.03$ |
| 4178.862 | -2.48 | 20830.582 | $36.3\pm 1.2$ | $-5.59\pm 0.01$ |
| 4303.176 | -2.49 | 21812.055 | $35.7\pm 1.1$ | $-5.52\pm 0.01$ |
| 4351.769 | -2.10 | 21812.055 | $52.6\pm 0.9$ | $-5.53\pm 0.01$ |
| 4555.893 | -2.29 | 22810.356 | $43.0\pm 0.9$ | $-5.51\pm 0.01$ |
| 4583.837 | -2.02 | 22637.205 | $67.4\pm 0.9$ | $-5.35\pm 0.01$ |
| 4629.339 | -2.37 | 22637.205 | $45.2\pm 0.8$ | $-5.41\pm 0.01$ |
| 5018.440 | -1.22 | 23317.632 | $89.1\pm 2.7$ | $-5.61\pm 0.04$ |
| 5169.033 | -0.87 | 23317.632 | $112.2\pm 2.0$ | $-5.45\pm 0.03$ |
| 5234.625 | -2.05 | 25981.630 | $45.0\pm 2.6$ | $-5.51\pm 0.03$ |
| 5276.002 | -1.94 | 25805.329 | $68.1\pm 1.5$ | $-5.16\pm 0.02$ |
| 5316.225 | 0.34 | 84035.139 | $89.0\pm 2.5$ | $-5.33\pm 0.02$ |
| 5316.615 | -1.85 | 25428.783 |  $\cdots$ |  $\cdots$ |
| 5316.784 | -2.91 | 25981.630 |  $\cdots$ |  $\cdots$ |
Ni II | 3849.554 | -1.88 | 32523.540 | $18.9\pm 2.4$ | $-6.59\pm 0.05$ |
| 4067.031 | -1.29 | 32499.529 | $16.7\pm 2.0$ | $-7.25\pm 0.06$ |

Note. — Abundances relative to hydrogen: $\log N({\rm X})/N({\rm H})$.

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