# Observations of the Ultraviolet-Bright Star ZNG 1 in the Globular Cluster M5
(NGC 5904)

William V. Dixon Space Telescope Science Institute, 3700 San Martin Drive,
Baltimore, MD 21218, USA

###### Abstract

We have analyzed archival spectra of the hot UV-bright star ZNG 1 in the
globular cluster M5 (NGC 5904) obtained with the Far Ultraviolet Spectroscopic
Explorer (FUSE) and the Space Telescope Imaging Spectrograph (STIS). From
these data, we derive an effective temperature $T_{\rm eff}$ = $43{,}000\pm
1400$ K, a surface gravity $\log g$ = $4.47\pm 0.08$, a rotational velocity
$v\sin i$ = $157\pm 12$ km s-1, and a mass $M=0.92\pm 0.17$ $M_{\sun}$. The
atmosphere is helium-rich ($Y=0.99$) and enhanced in CNO (relative to the
cluster). The spectrum exhibits wind features with a terminal velocity near
1500 km s-1 and strong discrete absorption components (DACs). The high helium
abundance, stellar mass, and rotational velocity suggest that the star is a
merger remnant, and its parameters are consistent with models of a pair of
merging He-core white dwarfs.

stars: abundances — stars: atmospheres — stars: individual (NGC 5904 ZNG 1
(catalog Cl* NGC 5904 ZNG 1)) — ultraviolet: stars

††facilities: FUSE, HST(STIS)

##  1 Introduction

Zinn et al. (1972, hereafter ZNG) identified seven UV-bright stars—stars that
lie above the horizontal branch and to the left of the giant branch in the
Hertzsprung–Russell diagram—in the globular cluster M5 (NGC 5904). The
brightest of these, ZNG 1, is a puzzling object with an unusual chemistry and
a high rotational velocity.

The star was first studied by Bohlin et al. (1983), who obtained ultraviolet
images of the cluster with a rocket-borne telescope. These images revealed a
hot, luminous star some 20″ from the cluster center, which the authors
identified as ZNG 1. Follow-up observations with the International Ultraviolet
Explorer (IUE) showed strong resonance lines of N V $\lambda 1240$ and C IV
$\lambda 1550$ and emission in the N IV] $\lambda 1487$ line. The authors
estimated an effective temperature around 35,000 K. The N V $\lambda 1240$
doublet is particularly strong, suggesting that the star is enhanced in
nitrogen. The N IV] emission was interpreted as evidence that a nitrogen-rich
planetary nebula (PN) may be present.

Combining several IUE spectra and comparing them to solar-abundance Kurucz
(1979) models, de Boer (1985) derived $T_{\rm eff}$ = $28{,}000\pm 4000$ and
$\log g=4.5$ and argued that the apparent N IV] $\lambda 1487$ emission is
spurious.

To investigate the possibility of a PN around ZNG 1, Napiwotzki & Heber (1997)
obtained H$\alpha$ imaging and UV spectroscopy with the Wide Field and
Planetary Camera 2 (WFPC2) and the Goddard High Resolution Spectrograph
(GHRS), respectively, aboard the Hubble Space Telescope (HST). The star cannot
be observed from the ground, because it lies $0\farcs 5$ from a bright G-type
companion. The star is easily resolved from its companion in the WFPC2
H$\alpha$ image, and there is no evidence of a planetary nebula. The GHRS
spectrum exhibits strong wind features in both N V $\lambda 1240$ and C IV
$\lambda 1550$ and photospheric lines of C III, C IV, N III, N IV, O IV, O V,
and S V. The spectrum is similar to that of HD 128220 B, for which $T_{\rm
eff}$ = $40{,}600\pm 400$ K, $\log g$= $4.5\pm 0.1$ (Rauch, 1993).

Observations with the Far Ultraviolet Spectroscopic Explorer (FUSE) revealed
that the star is most unusual. By fitting its spectrum with non-LTE (local
thermodynamic equilibrium) models, Dixon et al. (2004) derived $T_{\rm eff}$ =
$44{,}300\pm 300$ K, $\log g$ = $4.3\pm 0.1$, and a rotational velocity $v\sin
i$ = $170\pm 20$ km s-1. The atmosphere is helium-rich and enhanced in carbon,
nitrogen, and oxygen. The spectrum shows evidence of a wind and blue-shifted,
highly-ionized absorption, either around the star or along the line of sight.

Zech et al. (2008) combined the FUSE data with E140M echelle spectra obtained
with the Space Telescope Imaging Spectrograph (STIS) on HST. They found high-
velocity absorption in C IV, Si IV, O VI, and lower-ionization species with
local standard of rest (LSR) velocities of roughly $-140$ and $-110$ km s-1.
Using photoionization models and path-length arguments, they conclude that the
gas is not circumstellar but represents absorption from highly-ionized high-
velocity clouds along the line of sight.

In this paper, we return to a consideration of the star itself. Using the full
set of FUSE and STIS data (Table 1) and the latest non-LTE models, we re-
derive the stellar parameters, as well as the abundances of C, N, O, Si, S,
and Fe. We model the wind features. We argue that the star is the product of a
stellar merger, perhaps of a pair of He-core white dwarfs.

In Section 2, we present our data. In Section 3, we discuss our atmospheric
models and use them to derive stellar parameters and abundances. In Section 4,
we model the wind profiles and estimate the mass-loss rate. In Section 5, we
consider possible formation scenarios. We summarize our conclusions in Section
6.

Table 1: Summary of FUSE and STIS 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 |  03,734 | 2000 July 15 | A1080303 | Dixon
|  |  |  | 26,761 | 2003 Apr 11-13 | D1570301-3 | Dixon
STIS | E140M | 1140–1729 | 45,800 |  04,950 | 2003 Jul 08 | O6N404010-20 | Howk
|  |  |  |  07,850 | 2003 Jul 18-19 | O6N403010-30 | Howk

##  2 Observations and Data Reduction

###  2.1 FUSE Data

FUSE provides medium-resolution spectroscopy from 1187 Å to the Lyman limit
(Moos et al., 2000; Sahnow et al., 2000). ZNG 1 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 FUSE
data used in this work are available at https://doi.org/10.17909/158e-8q57
(catalog 10.17909/158e-8q57). For each FUSE channel, the extracted spectra
from all exposures were shifted to a common wavelength scale, weighted by
exposure time, and combined into a single file.

The signal-to-noise (S/N) ratio of the resulting spectrum ranges from 10 to 20
per 0.05 Å resolution element in the SiC bands (900–1000 Å) and from 30 to 40
in the LiF bands (1000-1187 Å). In most of our model fits, we employ only the
spectrum from the channel with the highest S/N ratio. An exception is the
900–1000 Å region, for which we have combined data from the SiC 1B and SiC 2A
detectors, weighted by exposure time, into a single spectrum. The two channels
have different line-spread functions, so combining their spectra reduces the
spectral resolution; however, this effect is insignificant given the star’s
high rotational velocity.

The FUSE wavelength calibration is reasonably accurate, but small offsets of
the target within the $30\arcsec\times 30\arcsec$ aperture can introduce a
zero-point offset in the wavelength scale. To place the data on an absolute
wavelength scale, we shift the spectrum from each channel so that the
velocities of their interstellar lines match those of our model, described in
Section 3.1.

###  2.2 STIS Data

The design and construction of STIS are described by Woodgate et al. (1998),
and information about its on-orbit performance is provided by Kimble et al.
(1998). ZNG 1 was observed using the STIS far-ultraviolet MAMA detector with
the E140M grating and the $0\farcs 2\times 0\farcs 2$ aperture. The data were
retrieved from StarCAT (Ayres, 2009, 2010), a collection of high-resolution
ultraviolet STIS spectra hosted at MAST. The StarCAT project combined data
from all observations, exposures, and spectral orders of every star observed
with STIS during its first seven years on orbit into a single, fully-
calibrated spectrum. The STIS calibration pipeline yields spectra with a
heliocentric wavelength scale; StarCAT spectra have an absolute wavelength
uncertainty of approximately 1.4 km s-1.

In the course of our analysis, we discovered that the flux error bars in the
StarCAT spectrum are too large by about a factor of two. For example, in a
flat region of the spectrum near 1172 Å, the standard deviation of the flux
array is 0.08 (in arbitrary units), while the mean value of the error array is
0.15. Repeating this calculation across the STIS spectrum, we derive an
average correction factor of 0.55 for the error array. Applying this
correction, we find that the StarCAT spectrum of ZNG 1 has a S/N ratio ranging
from about 20 to 50 per two-pixel resolution element.

##  3 Analysis

###  3.1 Interstellar Medium

The FUV spectrum of ZNG 1 is punctuated by a variety of interstellar
absorption features. In the FUSE bandpass, molecular hydrogen is the dominant
species, while high-ionization species are prominant at longer wavelengths.
Zech et al. (2008) concluded that the high-velocity absorption from both high-
and intermediate-ionization species is due to high-velocity clouds along the
line of sight. We use the Zech et al. results to model the interstellar
absorption, with high-velocity components at $v_{\rm LSR}$ $\sim\;-140$ and
$-110$ km s-1 and low-velocity components at $v_{\rm LSR}$ $\sim\;-59$ and
$-6$ km s-1. Zech et al. report that a two-component model, with $v_{\rm LSR}$
= $-135$ and $-18$, fits the H I absorption as well as the four-component
model, so we adopt the simpler model for this species. We convert to a
heliocentric reference frame using the relation $v_{\rm LSR}$ $-$ $v_{\rm
Helio}$ = $+13.25$ km s-1. For low-velocity species not fit by Zech et al., we
fit the absorption by eye, which is sufficient for our needs.

Synthetic interstellar absorption spectra are computed using software written
at the University of California, Berkeley, by M. Hurwitz and V. Saba. Given
the column density, Doppler broadening parameter, and velocity of each
component, the program computes a Voigt profile for each absorption feature
and produces a high-resolution spectrum of optical depth versus wavelength.
Wavelengths, oscillator strengths, and other atomic data are taken from Morton
(2003).

###  3.2 Model Atmospheres

We compute non-LTE stellar-atmosphere models using version 208 of the program
TLUSTY (Hubeny & Lanz, 1995). We employ atomic models similar to those used by
Lanz & Hubeny (2003) to compute their grid of O-type stars. Within TLUSTY, one
normally specifies all chemical abundances relative to hydrogen; however, for
hydrogen-deficient objects, Hubeny & Lanz (2017) recommend the use of helium
as the reference atom, and we have taken this advice. Since CNO are the most
abundant metals in the star’s photosphere, we begin by generating models with
H+He+C+N+O. To determine the carbon abundance, we fix N and O at the levels
determined by Dixon et al. (2004) and generate a grid of models that vary in
C/He. We repeat the process for N and O. For all other species, we generate a
grid of models with H+He+C+N+O+X, where CNO are fixed at their best-fit
values.

Given a model atmosphere, we compute a synthetic spectrum using version 54 of
the program SYNSPEC (Hubeny, 1988). For the FUSE data, the synthetic spectra
are convolved with a Gaussian of FWHM = 0.06 Å to match the FUSE line-spread
function. For the STIS spectrum, we use a Gaussian with FWHM = 0.025 Å. We
multiply the synthetic spectrum by a Fitzpatrick (1999) extinction curve
assuming $E(B-V)$ = 0.03 (Harris, 1996, 2010), extrapolated to the Lyman
limit, and the ISM absorption model described above. Finally, we scale the
model to reproduce the continuum in a nearby (apparently) line-free region in
the observed spectrum.

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.

Continuum placement is the dominant uncertainty in our fits. Because of its
high rotational velocity, only the strongest features in the spectrum of ZNG 1
can be identified unambiguously. Weak features from additional species may
depress the apparent continuum. Allowing our fitting routines to scale the
model to the mean level of the “pseudo-continuum” thus underestimates the true
continuum level. To estimate the uncertainty inherent in our continuum
estimate, we perform each fit twice, once with the model continuum fixed as
described and again with the model scaled by a factor of 0.97. The difference
in the two abundances is an estimate of the systematic error in our abundance
estimates. We add this term and the statistical error in quadrature to compute
our final error for a single absorption feature. In most cases, the continuum
uncertainty is the dominant contributor to the final error.

Dixon et al. (2004) derived a surface gravity $\log g$ = 4.3 and a helium
abundance of 85% by number. Preliminary analyses of these new data suggest
that the actual values are somewhat higher, so we assume $\log g$ = 4.5 and
$N$(H)/$N$(He) = 0.03 in our initial models.

###  3.3 Radial and Rotational Velocities

To determine the star’s radial and rotational velocities, we use the C III*
multiplet at 1175 Å, a strong feature uncontaminated by interstellar
absorption. The line appears in both the FUSE and STIS spectra. We begin with
the STIS spectrum, which has higher S/N and spectral resolution. We generate
models with carbon abundance 1, 3, and 6 times some fiducial value. For each
model, we generate synthetic spectra with rotational velocities between 120
and 200 km s-1 in steps of 10 km s-1. Assuming a radial velocity $v_{\rm
Helio}$ = 30 km s-1, we fit the spectrum, allowing the carbon abundance and
rotational velocity to vary freely, then compute $\chi^{2}$ for the best-fit
model. We repeat the process for integral values of $v_{\rm Helio}$ between 30
and 60 km s-1. We find that the STIS spectrum is best reproduced by a model
with $v_{\rm Helio}$ = $42\pm 2$ km s-1 and $v\sin i$ = $150\pm 15$ km s-1.
The error bar on the heliocentric velocity represents the range of velocities
whose $\chi^{2}$ values lie within $\Delta\chi^{2}=2.3$ of the minimum (the
range appropriate for two parameters of interest; Press et al. 1989); it is
thus purely statistical. The error bar on the rotational velocity is computed
as described above. We repeat the process for the FUSE spectrum, deriving
$v_{\rm Helio}$ = $45\pm 3$ km s-1 and $v\sin i$= $168\pm 19$.

We adopt the weighted mean of the FUSE and STIS values, $v_{\rm Helio}$ =
$43\pm 2$ km s-1 and $v\sin i$ = $157\pm 12$ km s-1 (where we weight by the
measurement uncertainties), as the radial and rotational velocities of ZNG 1.
The cluster has a mean heliocentric velocity of $53.2$ km s-1, with a central
velocity dispersion of 5.5 km s-1 (Harris, 1996, 2010), so the velocity of ZNG
1 is within $\sim 2\sigma$ of the cluster mean.

Lanz et al. (2004) derived a rotational velocity $v\sin i$ = 100 km s-1 from
the FUSE spectrum of the He-rich sdB star PG 1544+488. From a high-resolution
optical spectrum, Ahmad et al. (2004) determined that the star is actually a
spectroscopic binary. Because the FUSE observation was obtained in time-tag
mode, the authors were able to divide the data into short sub-exposures,
extract a spectrum from each, and show that the best-fit radial velocity
varies over the course of several hours. To see if a similar effect is at work
in ZNG 1, we divide the three FUSE D15703 observations into sub-exposures
roughly 1000 s long, place the resulting spectra on a heliocentric wavelength
scale by shifting them to match our ISM model between 1120 and 1160 Å, and fit
a Gaussian to the C III* multiplet at 1175 Å. The resulting radial velocities
are plotted in Fig. 1. Error bars are computed from the covariance matrix and
are thus purely statistical. The radial velocities derived from observations
D1570301 and D1570303 show some scatter, but are generally consistent with a
constant value. The velocities derived from observation D1570302 show more
scatter, but no significant trend. We conclude that the FUSE data provide no
evidence that the star is a spectroscopic binary and proceed under the
assumption that its broad absorption features reflect its high rotational
velocity.

The STIS spectra have both higher resolution and a higher S/N ratio;
unfortunately, they were obtained in ACCUM mode, so are not amenable to this
analysis.

Figure 1: Radial velocities derived from the FUSE observations of ZNG 1 in M5.
Figure 2: Deriving the effective temperature and abundance of C, N, and O.
Points with error bars represent the abundance derived from model fits to each
absorption feature. Solid and dashed lines are low-order polynomial fits to
the measured points. Dotted lines represent the weighted mean of the measured
abundances, computed at 10 K intervals. The black box in each panel denotes
the allowed ranges of temperature and abundance, as described in the text.

###  3.4 Effective Temperature and CNO Abundances

Following Dixon et al. (2019), we use the absorption features of multiple
ionization states of CNO to derive the star’s effective temperature. Consider
the middle panel of Fig. 2. Using a series of models with $T_{\rm eff}$ =
41,000 K and $\log N({\rm N})/N({\rm He})$ between $-5.0$ and $-3.1$, we fit
the N III $\lambda 979$ feature to determine the nitrogen abundance. We repeat
using models with temperatures increasing in steps of 1000 K to 45,000 K. The
resulting nitrogen abundances are plotted as dark blue points and connected by
a low-order polynomial (evaluated at 10 K intervals). Vertical bars represent
the uncertainties returned by the fitting routine. As the temperature rises,
the fraction of nitrogen in the form of N III falls, requiring a higher
nitrogen abundance to reproduce the observed feature. We repeat this procedure
for the N III $\lambda 991$ line, the N III $\lambda\lambda 1183,1184$
doublet, and the strong N IV feature at 955 Å. Interestingly, the strength of
the N IV line is constant in this temperature regime.

We repeat the process for C and O (top and bottom panels of Fig. 2). Note that
we have excluded from consideration features that yield wildly discrepant
abundances, as well as those that yield abundances with error bars so large
that they provide no useful constraints. This filter eliminates all of the C
IV lines111For example, for models with $T_{\rm eff}$ = 43,000 K, the C IV
$\lambda\lambda 1351,1353$ doublet yields $\log N({\rm C})/N({\rm He})$ =
$-2.49\pm 0.38$; the C IV $\lambda\lambda 1548,1550$ resonance doublet, which
is blended with a strong wind feature, yields an upper limit of $\log N({\rm
C})/N({\rm He})$ $<-3.30$; and a broad, weak C IV absorption feature centered
at 1654 Å yields $\log N({\rm C})/N({\rm He})$ = $-3.40\pm 1.38$., leaving
only C III and C III*. The C III* $\lambda 1175$ multiplet appears in both the
FUSE and STIS spectra, so we fit it twice. We see that the curves for these
two species cross at temperatures near 41,000 K, while the nitrogen curves
cross at temperatures near 44,000 K. The O III and O IV curves remain
stubbornly separate. To combine these results in a quantifiable way, we
compute the error-weighted mean abundance and the error-weighted standard
deviation as a function of temperature for each element. The mean abundance is
plotted as a dotted line in each panel. At each temperature step, we compute

$\chi^{2}=\sum\Bigl{\\{}[y_{i}-y(x_{i})]^{2}/\sigma_{i}^{2}\Bigr{\\}},$

where $y_{i}$ is the abundance derived from a single feature (or group of
features; solid or dashed line), $y(x_{i})$ is the mean abundance (dotted
line), $\sigma_{i}$ is the uncertainty in the derived abundance (vertical
bars), and the summation is taken over the 20 abundance curves plotted in Fig.
2. $\chi^{2}$ has a minimum at $T_{\rm eff}$ = 43,000 K. Our upper limit to
the temperature is $T_{\rm eff}$ = 44,400 K, set by the point at which
$\chi^{2}$ rises by 4.72 relative to its minimum. This choice of
$\Delta\chi^{2}$ is strictly correct for a model with four interesting
parameters (one temperature and three abundances) if all errors are normally
distributed (Press et al., 1989). For each element, our best-fit abundance is
the error-weighted mean value (computed above) at the best-fit temperature.
The abundance uncertainty is the larger of the error-weighted standard
deviation or the uncertainty in the weighted mean. The black box in each panel
denotes the allowed range of temperature and abundance for that element.
Derived abundances are presented in Table 2.

###  3.5 Surface Gravity and Hydrogen Abundance

Figure 3: He II $\lambda 958$ (left) and $\lambda 1640$ (right) features in
the spectrum of ZNG 1 in M5. Synthetic spectra with surface gravity $\log g$ =
4.3, 4.5, and 4.7 are overplotted. Blue bars indicate the spectral regions
included in our fit.

As shown in Fig. 3, the strong He II features at 958 and 1640 Å provide a
sensitive probe of the surface gravity. Assuming an effective temperature of
$T_{\rm eff}$ = 43,000 K, we fit each line separately. The 958 Å line in the
FUSE spectrum yields a surface gravity $\log g$ = $4.51\pm 0.09$, while the
1640 Å line in the STIS spectrum yields $\log g$ = $4.37\pm 0.16$. The
weighted mean of these two values is $\log g$ = $4.48\pm 0.08$.

Figure 4: Hydrogen and helium lines in the FUSE spectrum of ZNG 1 in M5. The narrow lines are interstellar. Synthetic spectra with $N$(H)/$N$(He) ratios of 0.05 (green), 0.01 (orange) and 0.001 (blue) are overplotted. An ISM model, which is incorporated into our models, is omitted from this figure for clarity. Blue bars indicate the spectral regions included in our fits. Table 2: Photospheric Abundances  Species | $\log N({\rm X})/N({\rm He})$ |  $\log\beta_{\mathrm{X}}$aaLogarithm of mass fraction.
---|---|---
H | $-1.70\pm 0.30$ | $-2.30\pm 0.30$
He |  $\cdots$ | $-0.00371$
C | $-3.77\pm 0.09$ | $-3.30\pm 0.09$
N | $-4.03\pm 0.09$ | $-3.49\pm 0.09$
O | $-3.18\pm 0.51$ | $-2.58\pm 0.51$
Si | $-5.81\pm 0.24$ | $-4.97\pm 0.24$
S | $-7.65\pm 0.86$ | $-6.75\pm 0.86$
Fe | $-5.84\pm 0.73$ | $-4.70\pm 0.73$

Fig. 4 shows a series of H I Lyman lines, each blended with a He II Balmer
line. The flat-bottomed features are due to interstellar H I; they are shifted
to the blue by 31 km s-1, while the stellar spectrum is shifted to the red by
43 km s-1. Most of the other ISM features are due to molecular hydrogen. The
colored curves are synthetic spectra with $N$(H)/$N$(He) ratios between 0.05
(green) and 0.001 (blue). An ISM model, included in our fits, is omitted from
this figure for clarity. Though only the red wing of each line is available,
it is clear from the figure that the stellar hydrogen abundance is quite low.
Our fits yield a ratio $N$(H)/$N$(He) = $0.020\pm 0.015$.

Because a change in the hydrogen fraction (which is effectively a change in
the helium abundance) could affect the strength of the helium lines and thus
the derived surface gravity, we repeat our fits of the He II features at 958
and 1640 Å using models with a hydrogen abundance $N$(H)/$N$(He) = 0.02. The
derived surface gravity falls by only 0.01 dex, to $\log g$ = $4.47\pm 0.08$,
so further iterations are unnecessary. Because the surface gravity and
hydrogen abundance have changed so little from our initial values, we need not
adjust the effective temperature and CNO abundances that we derived in Section
3.4. Our final stellar parameters are presented in Table 3.

Table 3: Stellar Parameters  Parameter | Value
---|---
$T_{\rm eff}$ (K) | $43{,}000\pm 1400$
$\log g$ [cm s-2] | $4.47\pm 0.08$
$R_{*}/R_{\sun}$ | $0.92\pm 0.02$
$M_{*}/M_{\sun}$ | $0.92\pm 0.17$
$\log(L_{*}/L_{\sun})$ | $3.42\pm 0.06$
$v_{\rm Helio}$ (km s-1) | $43\pm 2$
$v\sin i$ (km s-1) | $157\pm 12$

###  3.6 Metal Abundances

Four other metals, Si, P, S, and Fe, exhibit absorption lines in the spectrum
of ZNG 1, but only of a single ionization state. We compute their abundances
as for the lighter elements. A list of absorption features and the abundance
derived from each is presented in Table 4. The full set of abundances for ZNG
1 is presented in Table 2. Notes on individual species follow.

Silicon: While there are several strong silicon features in our spectrum, most
are blended with other photospheric lines. An exception is the Si
IV$\lambda\lambda 1393,1402$ resonance doublet. The blue wing of each is
contaminated by ISM features, so we fit only the red wing.

Phosphorus: Our FUSE spectrum includes the P V resonance transitions at 1117
and 1128 Å. The 1117 Å line exhibits what could be a P Cygni profile, but the
1128 Å line does not. Meanwhile, the 1128.0 Å line is blended with strong Si
IV $\lambda 1128.3$ and S V $\lambda 1128.7$ features, which are not well
reproduced by our models. We are thus unable to constrain the star’s
phosphorus abundance.

Sulfur: Most of the strong sulfur lines in our spectra are blended with other
photospheric features. Only S IV $\lambda 1072$ appears to be uncontaminated;
it is the source of our abundance determination.

Iron: There are many Fe V lines in these spectra, but none are very strong.
Instead of fitting individual features, we fit the region between 1373 and
1381 Å, where our models predict that iron is the dominant source of
absorption.

The large uncertainties in the quoted abundances of S and Fe reflect the
change in the derived abundance when the stellar continuum is scaled by a
factor of 0.97.

Table 4: Selected Metal Lines

Ion | $\lambda_{\rm lab}$ | $\log gf$ | $E_{l}$ | Abundance | Grating
---|---|---|---|---|---
| (Å) |  | (cm-1) |  |
Si IV | 1393.755 | 0.03 | 0.000 | $-5.82\pm 0.30$ | E140M |
Si IV | 1402.770 | -0.28 | 0.000 | $-5.80\pm 0.40$ |  |
S IV | 1072.996 | -0.83 | 951.100 | $-7.65\pm 0.86$ | FUSE |
Fe V | 1373.589 | 0.31 | 187719.001 | $-5.84\pm 0.73$ | E140M |
| 1373.679 | 0.19 | 187157.505 |  $\cdots$ |  |
| 1373.964 | -0.18 | 213649.212 |  $\cdots$ |  |
| 1374.119 | 0.40 | 233848.909 |  $\cdots$ |  |
| 1374.253 | -1.27 | 216779.102 |  $\cdots$ |  |
| 1374.788 | 0.30 | 258628.499 |  $\cdots$ |  |
| 1375.416 | -1.00 | 258628.499 |  $\cdots$ |  |
| 1375.790 | -0.73 | 216860.399 |  $\cdots$ |  |
| 1376.337 | 0.44 | 188395.286 |  $\cdots$ |  |
| 1376.451 | 0.05 | 186725.501 |  $\cdots$ |  |
| 1377.718 | -0.66 | 214525.796 |  $\cdots$ |  |
| 1378.088 | 0.05 | 258769.503 |  $\cdots$ |  |
| 1378.561 | 0.43 | 205536.398 |  $\cdots$ |  |
| 1379.040 | -0.15 | 209523.911 |  $\cdots$ |  |
| 1379.202 | -1.12 | 213649.212 |  $\cdots$ |  |
| 1380.112 | -0.14 | 186433.588 |  $\cdots$ |  |

Note. — Abundance relative to helium: $\log N({\rm X})/N({\rm He})$.

Figure 5: Resonance doublets of O VI, N V, and C IV in the FUSE and STIS
spectra of ZNG 1 in M5, which show strong P Cygni profiles and discrete
absorption components (DACs) blueshifted by about 1000 km s-1. Velocities are
relative to the short-wavelength component of each doublet. Horizontal bars
mark the positions of the photospheric features. Vertical bars mark the blue
edges of the DACs. The narrow spectral features are interstellar.

##  4 Stellar Wind

Several absorption features in the FUV spectrum of ZNG 1 in M5 show evidence
of a stellar wind. The P Cygni profile of the N V $\lambda\lambda 1238,1242$
doublet was first seen in the star’s IUE spectrum by de Boer (1985), and the
star’s GHRS spectrum reveals C IV $\lambda\lambda 1448,1450$ features that are
blueshifted by $\sim 900$ km s-1 (Napiwotzki & Heber, 1997). In our data,
spectacular wind features are present in the resonance doublet of O VI
$\lambda\lambda 1032,1038$ as well as those of N V and C IV. These profiles
also exhibit discrete absorption components (DACs) blueshifted by about 1000
km s-1. O V $\lambda 1371$, which is not a resonance line, also exhibits a P
Cygni profile.

Fig. 5 compares the O VI, N V, and C IV profiles. The three profiles are
aligned in velocity such that the short-wavelenth photospheric component is at
rest. The DACs have sharp blue edges at $-1080$ km s-1, but the absorption
extends an additional 400 km s-1 to the blue; this extension is particularly
obvious in the C IV profile.

Figure 6: Model fits to the O V $\lambda 1371$ feature, which shows a P Cygni
profile. The data are overplotted by the wind model described in the text
(thick solid line). Also shown are the transmitted (dotted line) and scattered
(dashed line) components of the model profile. The spectra are plotted in the
stellar rest frame. The transmitted spectrum includes our photospheric
absorption model.

To constrain the properties of the wind, we fit a simple model to the O V
$\lambda 1371$ feature, which shows a P Cygni profile, but no DACs (Fig. 6).
We use the line-fitting routines of D. Massa, which combine the Sobolev with
Exact Integration (SEI) algorithm described by Lamers et al. (1987) with the
wind optical depth law introduced by Massa et al. (1995). This approach
employs the standard beta-law parametrization of the velocity, which has the
form

$w(x)=w_{0}+(1-w_{0})(1-1/x)^{\beta},$ (1)

where $w(x)={v(x)}/{v_{\infty}}$, $x=r/R_{*}$, $R_{*}$ is the stellar radius,
and $w_{0}$ is the initial velocity at the base of the wind. (In this
formalism, all wind velocities are expressed as a fraction of the terminal
velocity $v_{\infty}$.) We follow Massa et al. (1995) in adopting
$w_{0}=0.01$. The parameter $\beta$, which defines the steepness of the
spatial velocity gradient in the wind, is constrained by the shape of the
emission component of the P Cygni line. Other inputs to the code are the
turbulent velocity $w_{\rm D}$, which simulates macroscopic velocity fields in
the wind by smoothing the wind profile; the synthetic photospheric spectrum;
and $\tau_{\rm r}(w)$, the radial component of the optical depth of the
absorbing species as a function of the wind velocity.

The line profile is fit by adjusting $v_{\infty}$, $\beta$, $w_{D}$ and the
elements of $\tau_{\rm r}(w)$. The O V profile is well fit by the parameters
$v_{\infty}=600$ km s-1, $\beta=0.3$, and $w_{\rm D}=0.15$. The optical depth
falls smoothly from $\tau=0.4$ at the base of the wind to zero at
$v_{\infty}$. Our best-fit model is presented in Fig. 6. A complete
explanation of this technique, its applicability, and its limitations may be
found in Massa et al. (2003).

From the stellar abundances and our wind model, we can estimate $\dot{M}q(w)$,
the product of the stellar mass-loss rate and the fraction of oxygen in the
form of O+4 as a function of velocity in the wind. Over the velocity range
$0.1<w<0.9$, where the model is best constrained, $\dot{M}q(w)$ falls from
roughly $2\times 10^{-12}$ to $3\times 10^{-14}$ $M_{\sun}$ yr-1. Our model
predicts that $q=10^{-3}$ at the outer limit of the stellar atmosphere. If
this value is appropriate for the base of the wind, then $\dot{M}\sim 2\times
10^{-9}$ $M_{\sun}$ yr-1.

The blue-shifted absorption component of the O V $\lambda 1371$ profile
extends to about $-600$ km s-1, but the corresponding features of the
resonance lines extend to $-1500$ km s-1. The O VI absorption features are
weak in the photospheric spectrum, but the O VI wind feature is quite
prominent. Indeed, our models predict that the fraction of oxygen in the form
of O+5 is less than $10^{-13}$ at the edge of the stellar atmosphere. We
conclude that most of O+4 that leaves the star is ionized to O+5 before the
wind velocity exceeds 600 km s-1.

Discrete absorption components (DACs) are ubiquitous in the spectra of
luminous OB stars. They are localized (in velocity) absorption features that
generally appear at intermediate velocities and accelerate toward higher
velocities (that is, to the blue), with a corresponding decrease in velocity
width. Attributed to slowly-moving density enhancements in the wind, DACs in
O-type stars can represent as much as 50% of the wind material, although 20%
is more common. A strong correlation between the acceleration and recurrence
of DACs and the projected rotation velocity, $v\sin i$, has been observed,
though the nature and origin of this connection remain unclear. (See Massa et
al. 1995, Prinja et al. 2002, and references therein.) Given the high
rotational velocity of ZNG 1, DACs may be a common feature of its spectrum.
Indeed, the two FUSE observations of the star were taken more than three years
apart, yet the DACs in their spectra are identical.

##  5 Discussion

###  5.1 Cluster Membership

Before comparing ZNG 1 with other members of M5, we should confirm that it is
a cluster member. The HST UV Globular Cluster Survey (HUGS; Piotto et al.,
2015; Nardiello et al., 2018) provides astrometry, photometry, and stellar
membership probabilities for 56 globular clusters. They assign a membership
probability of 96.1% to ZNG 1.

Bond (2021) extracted the coordinates, parallaxes, and proper motions of the
entire ZNG sample from the Gaia Early Data Release 3 (EDR3; Gaia Collaboration
et al., 2016, 2021). He concluded that the proper motion of ZNG 1 is
consistent with cluster membership, but its parallax is discordant. The Gaia
color for this star, $BP-RP=1.096$, is too red for a star as hot as ZNG 1. The
Gaia Data Release 3 (Gaia Collaboration et al., 2022) entry for the star is
identical to that of EDR3, but includes estimates of the star’s effective
temperature and surface gravity. These values, $T_{\rm eff}$ = 5443 K and
$\log g$ = 2.8, suggest that the star is blended with its G-type companion in
the Gaia data.

In Section 3.3, we derive a radial velocity $v_{\rm Helio}$ = $43\pm 2$ km
s-1, a value about 2$\sigma$ from the cluster mean. Based on its proper motion
and radial velocity, we conclude that ZNG 1 is a cluster member.

###  5.2 Photospheric Abundances

The chemical abundances derived from our line fits are expressed as number
fractions relative to helium, which is the dominant species in the star.
Published abundances for the other stars in M5 are expressed in terms of
hydrogen. To facilitate comparison, we assume that iron was neither created
nor destroyed in the evolutionary process that produced ZNG 1 and plot $\log
N$(X)/$N$(Fe) for each element in Fig. 7. We see that C, N, and O have roughly
equal abundances. C and O are enhanced by about 1.5 dex relative to the
cluster, and N is slightly enhanced. The Si abundance is close to the cluster
value, while the S/Fe ratio is subsolar. Any evolutionary scenario proposed
for ZNG 1 must be able to explain this abundance pattern.

Figure 7: Photospheric abundances relative to iron of ZNG 1 (stars), M5 (red
lines), and the solar photosphere (black lines). Cluster values from Mészáros
et al. (2020). Solar values from Asplund et al. (2009).

###  5.3 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. For ZNG 1, the HUGS catalog
quotes Vega magnitudes of 11.97, 12.58, and 14.30 through the F275W, F336W,
and F438W filters, respectively. (We omit the F814W measurement, as it may be
contaminated by the star’s companion.) Uncertainties are not provided for
these values. We generate a synthetic spectrum using our best-fit model scaled
by a Fitzpatrick (1999) extinction curve with $R_{V}=3.1$ and $E(B-V)$ = 0.03
(Harris, 1996, 2010). For each filter, we compute synthetic stellar magnitudes
using the recipe provided in Riello et al. (2021). The ratio between the
observed and model fluxes is $\phi$ = (9.20, 9.93, and $10.15)\times 10^{-23}$
for the three filters. We adopt the mean of these values, $\phi=(9.76\pm
0.40)\times 10^{-23}$, as our scale factor.

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 a combination
of Gaia EDR3, HST, and literature data, Baumgardt & Vasiliev (2021) derive a
distance to M5 of $7.479\pm 0.060$ kpc. Adopting this distance and our scale
factor, we derive a stellar radius $R_{*}/R_{\sun}$ of $0.92\pm 0.02$.
Applying our best-fit surface gravity ($\log g$ = 4.47), we find that the
stellar mass $M_{*}/M_{\sun}$ is $0.92\pm 0.17$. Finally, combining the
stellar radius with our best-fit effective temperature ($T_{\rm eff}$ = 43,000
K), we derive a stellar luminosity $\log(L_{*}/L_{\sun})$ of $3.42\pm 0.06$.

###  5.4 Evolutionary Status

AGB stars in globular clusters consist of a C-O core surrounded by concentric
shells of He and H. During most of its time on the AGB, a star is powered by
the outer, H-burning shell, the ashes of which rain onto the He shell. When
the mass of the He shell becomes great enough, it ignites in a flash or
thermal pulse. As the He-shell flash dies out, conditions become favorable for
the convective dredge-up of material. This so-called third dredge-up brings
primary nucleosynthesis products from both H- and He-shell burning to the
surface; among these are C, the $s$-process elements, and He (Herwig, 2005).

The effective temperature and luminosity of ZNG 1 place it on the evolutionary
tracks of stars evolving from the blue horizontal branch (BHB; Moehler et al.,
2019), yet ZNG 1 is unlikely to be the descendent of an ordinary BHB star.
First, relative to the cluster, the star is equally enhanced in both C and O
(Fig. 7). Third dredge-up brings 5 to 10 times more C than O to the stellar
surface (Herwig, 2005), so it cannot explain the observed abundance pattern.
Second, the star’s mass considerably exceeds that of white dwarfs forming in
globular clusters today ($\sim 0.53$ $M_{\sun}$; Kalirai et al., 2009). Third,
its high rotational velocity cannot be explained by single-star evolution.
Indeed, the high mass and rotational velocity of ZNG 1 strongly suggest that
the star is some sort of merger remnant.

While single stars in globular clusters evolve into white dwarfs with C-O
cores, binary stars may follow an alternative evolutionary channel: as the
more massive star ascends the RGB, it can overflow its Roche lobe, shed its
outer envelope, and become a He-core white dwarf (WD) without passing through
the HB or AGB phases. Later, the secondary can undergo the same process,
resulting in a pair of He-core WDs. Given sufficient time, the stars will
merge, generally resulting in a helium-rich hot subdwarf of spectral type He-
sdO or He-sdB (Jeffery & Zhang, 2020).

Figure 8: Evolutionary tracks from Zhang & Jeffery (2012) for binary He-WD
merger remnants with final masses of 0.8 $M_{\sun}$ (orange) and 0.9
$M_{\sun}$ (blue). Red dots represent the Geier (2020) sample of helium-rich
subdwarf O and B stars. A representative error bar is plotted in the upper
right-hand corner. The black star marks the effective temperature and surface
gravity of ZNG 1.

In Fig. 8, we plot evolutionary tracks from Zhang & Jeffery (2012) for binary
He-WD merger remnants with abundances $Z=0.001$ and masses of 0.8 $M_{\sun}$
(orange) and 0.9 $M_{\sun}$ (blue). These are composite models, in which most
of the material of the secondary is transferred directly to the surface of the
primary, where it heats up and expands to form a hot corona. The remainder of
the material (0.1 $M_{\sun}$ in these models) forms a disk around the primary,
from which it is transferred to the stellar surface at a rate of $10^{-5}$
$M_{\sun}$ yr-1. For the 0.9 $M_{\sun}$ model, the surface gravity rises to
$\log g$ = 4.5 approximately 65,000 years after the merger. At an accretion
rate of $10^{-5}$ $M_{\sun}$ yr-1, a 0.1 $M_{\sun}$ disk would last 10,000
years, so would have long since disappeared. Overplotted are ZNG 1 and a
sample of He-sdO and He-sdB stars (Geier, 2020). ZNG 1 lies in a region of the
diagram consistent with a merger product that is evolving toward the helium-
burning main sequence.

Figure 9: The distribution of surface abundance for He-sdO (Hirsch 2009,
circles), He-sdB stars (Naslim et al. 2010, diamonds), and ZNG 1 (star),
ordered by carbon abundance. Elements are identified by color: black for C,
green for N, and red for He. Horizontal lines indicate the solar abundance of
each element.

In Fig. 9, derived from similar figures in Hirsch (2009) and Zhang & Jeffery
(2012), we plot the distribution of He, C, and N mass fractions in a sample of
He-sdO and He-sdB stars. The C abundance exhibits a bimodal distribution:
either C is enhanced above the solar value ($\log\beta_{\mathrm{C}}>-2.6$) or
it is strongly depleted (relative to solar). The C-rich objects exhibit a wide
range of N values ($-3.5<\log\beta_{\mathrm{N}}<-2.2$), while the C-poor stars
are all N-rich, with N in a narrow range ($-2.8<\log\beta_{\mathrm{N}}<-2.0$).
ZNG 1 lies between the C-poor and C-rich groups, though closer to the former,
while its N fraction is consistent with those of the C-rich sample. Zhang &
Jeffery (2012) were able to reproduce the abundance pattern displayed in Fig.
9 using composite models with a range of final masses. The C and N fractions
of ZNG 1 correspond to a model with $M\sim 0.67$ $M_{\sun}$, a value slightly
lower than our mass estimate. While the field stars in this figure have solar
metallicities, post-merger evolution is driven primarily by a helium-burning
shell in the merger remnant, so the tracks and nucleosynthesis should be
little affected by the stars’ initial metallicity.

Let us consider the high rotational velocity of ZNG 1. Gourgouliatos & Jeffery
(2006) model the angular-momentum evolution of a merging He-WD binary. In this
model, helium burning in a shell around the core of the primary causes the
merger product to expand into a giant. As the star evolves toward the helium
main sequence, it contracts, and its rotational velocity rises dramatically.
If no angular momentum were lost during the merger process, then a 0.7
$M_{\sun}$ remnant would have a rotational velocity of $\sim 1000$ km s-1 by
the time it reached the helium main sequence. It is likely that mass—and with
it angular momentum—is lost during the merger process, probably at the point
when the secondary fills its Roche lobe.

If high rotational velocities are possible, why are they so unusual? The
rotational velocities of the helium subdwarfs in Fig. 9 are less than 35 km
s-1 (Hirsch, 2009; Naslim et al., 2010). Were they more efficient at ejecting
angular momentum than ZNG 1? Perhaps the amount of angular momentum lost to
opacity-driven outflows is a function of metallicity. Alternatively, the mass
of ZNG 1 is relatively high. Perhaps it inherited a higher angular momentum
from its more massive progenitors.

##  6 Conclusions

We have analyzed archival FUV spectra of the hot UV-bright star ZNG 1 in the
globular cluster M5 (NGC 5904). From these data, we derive the star’s
heliocentric and rotational velocity, photospheric parameters, and chemical
abundances. The atmosphere is helium-rich and enhanced in CNO (relative to the
cluster). The spectrum exhibits wind features with a terminal velocity near
1500 km s-1 and strong discrete absorptions components (DACs). The high helium
abundance, stellar mass, and rotational velocity suggest that the star is a
merger remnant, and its parameters are consistent with models of a pair of
merging He-core white dwarfs.

The author wishes to thank M. M. Miller Bertolami for thoughtful discussions,
including the suggestion that ZNG 1 is the remnant of a binary WD merger, and
C. S. Jeffery for providing both his evolutionary tracks and helpful comments
on the manuscript. 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; 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.

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