11institutetext:  INAF-Osservatorio Astronomico di Padova, 36012 Asiago (VI),
Italy,
11email<EMAIL_ADDRESS>22institutetext: INAF-Osservatorio Astronomico
di Capodimonte, via Moiariello 16, 80131 Napoli, Italy 33institutetext: INAF-
Osservatorio Astrofisico di Catania, via S. Sofia 78, 95123 Catania, Italy
44institutetext: INAF-Osservatorio di Astrofisica e Scienza dello Spazio, via
Gobetti 101, 40129 Bologna, Italy 55institutetext: Departamento de Ciencias
Físicas, Universidad Andrés Bello, Fernández Concha 700, Las Condes, Santiago,
Chile 66institutetext: Lund Observatory, Department of Astronomy and
Theoretical Physics, Box 43, SE-221 00 Lund, Sweden 77institutetext: Institute
for Astronomy, Astrophysics, Space Applications and Remote Sensing, National
Observatory of Athens, Penteli GR 15236, Greece

# THA 15$-$31: Discovery with VLT/X-Shooter and Swift/UVOT of a new symbiotic
star of the accreting-only variety

U. Munari 11 J. K. Alcalá 22 A. Frasca 33 N. Masetti 4455 G. Traven 66 S.
Akras 77 L. Zampieri 11

(Received ; accepted )

We report the discovery and characterization of a new symbiotic star of the
accreting-only variety, which we observed in the optical/near-infrared (NIR)
with VLT/X-Shooter and in the X-rays/ultraviolet with Swift/UVOT+XRT. The new
symbiotic star, THA 15$-$31, was previously described as a pre-main sequence
star belonging to the Lupus 3 association. Our observations, ancillary data,
and Gaia EDR3 parallax indicate that THA 15$-$31 is a symbiotic star composed
of an M6III red giant and an accreting companion, is subject to $E_{B-V}$=0.38
reddening, and is located at a distance of $\sim$12 kpc and at 1.8 kpc above
the Galactic plane in the outskirts of the Bulge. The luminosity of the
accreting companion is $\sim$100 L⊙, placing THA 15$-$31 among the symbiotic
stars accreting at a high rate (2.5$\times$10-8 M⊙ yr-1 if the accretion is
occurring on a white dwarf of 1 M⊙). The observed emission lines originate
primarily from HI, HeI, and FeII, with no HeII or other high-excitation lines
observed; a sharp central absorption superimposed on the Balmer emission lines
is observed, while all other lines have a simple Gaussian-like profile. The
emission from the companion dominates over the M6III red giant at $U$ and
$B$-band wavelengths, and is consistent with an origin primarily in an
optically thick accretion disk. No significant photometric variability is
observed at optical or NIR wavelengths, suggesting either a face-on orbital
orientation and/or that the red giant is far from Roche-lobe filling
conditions. The profile of emission lines supports a low orbital inclination
if they form primarily in the accretion disk. An excess emission is present in
AllWISE W3 (12$\mu$m) and W4 (22$\mu$m) data, radiating a luminosity $\geq$35
L⊙, consistent with thermal emission from optically thin circumstellar dust.

###### Key Words.:

symbiotic stars

##  1 Introduction

THA 15$-$31 was first noted by The (1962) while compiling a list of faint
H-alpha emission stars in Lupus and Scorpius, and later rediscovered as Sz 105
by Schwartz (1977) during a deep red objective-prism survey of southern dark
clouds. This led to THA 15$-$31 being considered as a candidate young stellar
object (YSO) and as such it was included in many surveys of the Lupus 3 star
forming region, with observations ranging from X-rays to optical, infrared
(IR), and radio (Hughes et al. 1994; Chen et al. 1997; Ghez et al. 1997;
Krautter et al. 1997; Nuernberger, Chini, & Zinnecker 1997; Gómez & Persi
2002; Kirk & Myers 2011; Mužić et al. 2014). The classification of THA 15$-$31
as a classical T Tau star of M4 spectral type with 0.15 M⊙ mass and 0.03 L⊙
luminosity was generally copied from one survey to another.

The first hint of an erroneous YSO classification of THA 15$-$31 came from
López Martí et al. (2011) who investigated the distribution in proper motion
of members of Lupus 3 and found that THA 15$-$31 belongs to a heterogeneous
group of background objects that are highly distinct from true YSOs in the
foreground. The problematic absence of LiI 6708 absorption in the spectra of
THA 15$-$31 —a hallmark of YSOs— was then noted by Alcalá et al. (2014), while
Frasca et al. (2017) derived a surface gravity for THA 15$-$31 that is more in
line with a cool giant than a YSO. Based on these findings, THA 15$-$31 was
then finally rejected by Alcalá et al. (2017) from the list of YSO members of
Lupus 3. This is, in broad terms, the current and rather limited knowledge
about THA 15$-$31.

Our interest in THA 15$-$31 was triggered by the recent release of Gaia EDR3
(Gaia Collaboration 2020), which lists a 0.0422$\pm$0.0398 mas parallax for
the object. Although still affected by a large error, it significantly
improves upon the previous value of 0.0323$\pm$0.0765 mas given in Gaia DR2,
and definitively supports a large distance (at least several kpc) to THA
15$-$31, which in turn implies the presence of a cool giant in the system and
definitively rules out any link to the Lupus dark clouds. To investigate its
nature, we analyzed an archive VLT/X-Shooter optical-IR spectrum of THA
15$-$31 and obtained satellite ultraviolet (UV) observations with the Neil
Gehrels Swift Observatory (hereafter Swift). Supplemented by epoch photometry
from NeoWISE and ASASSN, as well as other ancillary data, we now show that THA
15$-$31 is indeed an accreting-only symbiotic star (SySt), with an accretion
luminosity ($L_{\rm acc}$) that is much higher than that characterizing the
sample of 33 new, accreting-only SySts that we recently identified and studied
(Munari et al. 2021) among the 0.6 million stars observed by the GALAH
spectroscopic survey of the southern sky (de Silva et al. 2015, Buder et al.
2021).

Symbiotic stars are interacting binaries where a red giant (RG) fuels a hot
companion (HC) via accretion either through Roche-lobe overflow or wind
intercept. In the vast majority of SySts the hot companion is believed to be a
white dwarf (WD), while the hard and pulsed X-ray emission observed in a small
subset suggests that the accretion takes place on a neutron star (NS). The
RG+WD symbiotic stars are broadly divided into two major groups (see the
recent review by Munari 2019 for details): those accreting-only (acc-SySts)
whose optical spectra are dominated by the RG with no or weak emission lines,
and the burning-type (burn-SySts) with optical spectra dominated by a strong
nebular continuum and a rich emission line spectrum, that originate from the
wind of the RG which is largely ionized by the very hot and luminous WD
undergoing surface nuclear burning of accreted material.

It is possible that SySts spend most of their time in the accreting-only
phase, quietly accumulating material on the surface of the WD. When enough
material has built up, nuclear burning begins. If the accreted matter is
electron degenerate, the burning proceeds explosively (Starrfield et al. 2020)
resulting in a nova outburst. If the accreted matter is instead not electron-
degenerate, the nuclear burning develops slowly and in thermal equilibrium
(Fujimoto et al. 1982).

There is a clear disproportion among cataloged SySts (Allen 1984, Belczyński
et al. 2000, Merc et al. 2019, Akras et al. 2019a) in favor of the burn-SySt
type which are much more easily discovered, and the known examples of the acc-
SySt variety could simply be “the tip of the iceberg” (Mukai et al. 2016). The
acc-SySts are well distributed in accretion luminosity, from the 10$-$100 L⊙
of some well-studied objects like EG And (Skopal 2005, Nuñez et al. 2016) to
the emerging large population of lower 1$-$10 L⊙ objects exemplified by the
first sample of 33 SySts discovered by Munari et al. (2021) in the course of
the GALAH survey. The higher $L_{\rm acc}$ systems are characterized by a
brighter HC (collectively including the accretion disk, its central star, and
some RG wind ionized by them both) and by much stronger emission lines
compared to lower $L_{\rm acc}$ systems, although generally missing the higher
ionization features like HeII, [FeVII], OVI, or [FeX] that adorn the spectra
of the burn-SySt type. SySts accreting on neutron stars, while generally
inconspicuous at optical and near-UV wavelengths, may radiate abundantly in
hard X-rays (e.g., Masetti et al. 2007a,b, Nucita et al. 2014).

Identification of new SySts of the accreting-only, high-$L_{\rm acc}$ variety
appears to be important for properly measuring their frequency relative to the
lower $L_{\rm acc}$ systems, and will help to address the fundamental question
of whether accretion luminosity can be used to distinguish between SySts of
different nature (e.g., in terms of mass, orbital separation, and filling of
the Roche lobe), or whether accretion simply proceeds in a highly sporadic
fashion, irrespective of the nature of the underlying SySt. The latter applies
to the symbiotic star and recurrent nova T CrB, as discussed in Munari et al.
(2016) and Luna et al. (2020): for 70 years following its nova eruption in
1946, the accretion luminosity was so low that only feeble H$\alpha$ emission
was generally visible superimposed on the spectrum of the M3III cool giant.
However, starting with 2015, the $L_{\rm acc}$ in T CrB hugely increased and
the optical spectra are now persistently dominated by a strong nebular
continuum and imposing emission lines, in a replica of the accretion surge
that preceded the 1946 thermonuclear outburst (Hachenberg & Wellmann 1939).

Figure 1: Optical spectrum of THA 15$-$31 as recorded with X-Shooter (black)
is fitted (red) by combining the spectrum of the cool giant (M6III, green
spectrum) and that of the hot companion (blue line). The latter is a
bremsstrahlung spectrum taken to approximate the mix of the accretion disk,
the central star, and possibly some RG wind ionized by them. The spectra of
the M6III cool giant and of the hot companion are reddened by $E_{B-V}$=0.38
following the Fitzpatrick (1999) reddening law for the standard $R_{V}=3.1$
case. The inset is a zoomed-in view of the head of the Balmer series and the
Balmer jump.

##  2 Observations

###  2.1 X-Shooter spectrum

The X-Shooter data used here were acquired on April 18, 2012, within the
context of the X-Shooter guaranteed-time observations (GTO) granted to INAF
(cf. Alcalá et al. 2011, 2014 and references therein). X-Shooter provides
simultaneous wavelength coverage from $\sim$3000 Å to $\sim$2.48 $\mu$m. For
the observations of THA 15-31, slits of
0$\aas@@fstack{\prime\prime}$5/0$\aas@@fstack{\prime\prime}$4/0$\aas@@fstack{\prime\prime}$4
were used in the UltraViolet-Blue/Visual/Near-InfraRed (UVB/VIS/NIR) arms,
respectively, yielding resolving powers of 9100/17400/10500.

The observations were performed at parallactic angle with one cycle using the
A-B nodding mode, with single-node exposure of 150 s, 100 s, and 100 s for a
total of 300 s, 200 s, and 200 s in the UVB, VIS, and NIR arms, respectively.
Telluric standard stars were observed with the same instrumental setup and at
very similar airmass to THA 15-31. Flux standards during the same night were
observed through the 5 arcsec slits for flux calibration purposes.

Data reduction was performed independently for each spectrograph arm using the
X-Shooter pipeline (version 1.3.7; Modigliani et al. 2010). The standard steps
of processing included bias and dark subtraction, flat-fielding, optimal
extraction, sky subtraction, and wavelength and flux calibration (the latter
carried out in IRAF111IRAF is distributed by the National Optical Astronomy
Observatories, which are operated by the Association of Universities for
Research in As- tronomy, Inc., under cooperative agreement with the National
Science Foun- dation.). Wavelength shifts due to instrumental flexure were
corrected for using the flexcomp package within the pipeline. The precision in
wavelength calibration is better than 0.01 pix in the UVB and VIS arms,
corresponding to 0.02 Å, but errors can be as large as $\sim$0.06 Å in the NIR
arm. The wavelength calibration allows an average precision in radial velocity
of $\sim$1.0 km/s in the VIS arm. The telluric correction was performed
independently in the VIS and NIR spectra, as explained in Appendix A of Alcalá
et al. (2014). By comparing the response function of different flux standards
observed during the same night, we estimate an intrinsic precision on the flux
calibration of $\sim$5% (cf. Alcalá et al. 2014, 2017).

Table 1: Swift-UVOT observations of THA 15$-$31. filter | $\lambda_{c}$ |  Fλ | percent
---|---|---|---
| (Å) |  ($\times$10-16 erg cm-2 s-1 Å-1) | error
UVW1 | 2600 | 2.9 | 10%
UVM2 | 2246 | 1.1 | 20%
UVW2 | 1928 | 1.5 | 14%

###  2.2 Swift observations

Our Swift observations of THA 15$-$31 were acquired on April 9, 2021, starting
at 21:43 UT with the on-board instruments X-Ray Telescope (hereafter XRT,
Burrows et al. 2005) and UltraViolet-Optical Telescope (hereafter UVOT, Roming
et al. 2005). The XRT allows coverage of the X–ray band between 0.3 and 10
keV, whereas UVOT data were collected using the three available UV filters
(cf. Poole et al. 2008 and Breeveld et al. 2011 for details). On-source
pointing was simultaneously performed with the two instruments; it lasted 1560
s for XRT, whereas for UVOT 444 s were spent using the UVW1 filter and 515 s
in both UVM2 and UVW2 filters.

Count rates on Level 2 (i.e. calibrated and containing astrometric
information) UVOT images at the position of our target were measured through
aperture photometry using 5′′ apertures, whereas the corresponding background
was evaluated using a combination of several circular regions in source-free
nearby areas. Magnitudes were measured with the uvotsource task. The data were
then calibrated using the UVOT photometric system described by Poole et al.
(2008), and we included the recent (November 2020) fixings recommended by the
UVOT team. The observed fluxes in the three UV filters are listed in Table 1
and the corresponding magnitudes in Table 4 (Vega scale).

THA 15$-$31 was not detected in X-rays with XRT. Data analysis was performed
using the xrtdas standard pipeline package (xrtpipeline v. 0.13.4) in order to
produce screened event files. All X-ray data were acquired in photon counting
(PC) mode (Hill et al. 2004) adopting the standard grade filtering (0-12 for
PC) according to the XRT nomenclature. For each source, scientific data were
extracted from the images using an extraction radius of 47′′ (20 pixels)
centered at the optical coordinates of the source, while the corresponding
background was evaluated in a source-free region of radius 94′′ (40 pixels).

We do not detect any X-ray emission down to a 3$\sigma$ limit of
2.2$\times$10-3 cts s-1 in the 0.3–10 keV band at the position of THA 15$-$31;
this estimate was obtained using the xspec package. Assuming a bremsstrahlung
spectrum, with temperature $kT$ = 2 keV, which gives a 0.3–10 keV flux limit
of $<$6.3$\times$10-14 erg cm-2 s-1.

##  3 Results

###  3.1 Spectral type, radial velocity, and reddening

The optical part of the X-Shooter spectrum shown in Figure 1 is dominated by
the molecular TiO absorption bands of the RG, and this offers the possibility
to derive some of the basic properties of THA 15$-$31: (1) the relative
intensity of molecular bands and their fine structure, which vary by huge
amounts among cool giants, thus allowing us to derive an accurate spectral
type for the RG; (2) the veiling by the emission from HC, which increases
toward shorter wavelengths; and (3) the overall slope, controlled by the
amount of reddening. The reddening and spectral type are determined without
significant degeneracy among them, contrary to the case of modeling the
spectral energy distribution (SED) based on broad-band photometric data only:
the key to success here is the the strong and sharp band-heads that develop
within individual TiO bands over a short wavelength interval to make their
interplay independent from reddening. The technique was thoroughly tested and
applied in Corradi et al. (2010) for the analysis of new SySts discovered
photometrically by the INT Photometric H$\alpha$ Survey (IPHAS) and re-
observed spectroscopically for confirmation and characterization.

The fitting to the optical part of the X-Shooter spectrum of THA 15$-$31 is
presented in Figure 1. To this aim we adopted the library of intrinsic spectra
of real M-type giants by Fluks et al. (1994)222We corrected an incorrect shape
in the Fluks et al. spectrum over the 7300$-$8200 Å range and replaced the
linear extrapolation below 3900 Å by substituting them with observations of HD
148783 that we obtained with the Asiago 1.22m+B&C telescope. This star is the
prime standard for spectral type M6III according to the compilation by
Yamashita et al. (1977). and the interstellar reddening law of Fitzpatrick
(1999) for the standard $R_{V}$=3.1 case. The emission from HC is a mixture of
different sources (accretion disk, central star, nebular continuum) present in
unknown proportions, and physical parameters that cannot be unambiguously
determined with available data. As a fair approximation of the overall shape
of the HC emission in the wavelength range covered by Figure 1, we adopted a
bremsstrahlung spectrum for gas at $T_{e}$=10,000 K. However, none of the
results presented in this paper depend on this assumption.

The fit of the X-Shooter spectrum displayed in Figure 1 is remarkably good,
fixing the reddening to $E_{B-V}$=0.38$\pm$0.04 and the spectral type to
M6III, with an uncertainty smaller than 1 spectral subclass, because the
structure of molecular bands is strongly different for stars with M5III and
M7III spectral types. A spectral type M6III is the one most frequently found
in Galactic SySts. The modest $E_{B-V}$=0.38$\pm$0.04 reddening agrees well
with the relatively high Galactic latitude of the object ($b$=+8.5 deg), and
with the Stilism 3D-map of the Galactic interstellar reddening (Lallement et
al. 2014, Capitanio et al. 2017): the latter indicates an exit of the line of
sight to THA 15$-$31 from the dust slab on the Galactic plane at a distance of
600 pc from Earth, where it reaches $E_{B-V}$=0.40$\pm$0.05. The agreement is
also good with the value $E_{B-V}$=0.34$\pm$0.03 derived from the 3D map of
Galactic extinction by Marshall et al. (2006). Unfortunately, the location on
the sky of THA 15$-$31 is too far south for other popular 3D maps of Galactic
extinction to be considered, such as those based on IPHAS (Sale et al. 2014)
or PanSTARRS (Green et al. 2019) sky surveys. The limited resolution of
X-Shooter is not sufficient to attempt to resolve the interstellar from the
stellar component of the NaI line, and thus derive the reddening from its
equivalent width (see also below, Sect. 3.6). The heliocentric radial velocity
of the M6III giant is determined as

${\rm RV}_{\odot}=+4\pm 2{\rm km\leavevmode\nobreak\ s^{-1}}$ (1)

by application of the rotfit code (Frasca et al. 2006), specifically modified
for X-Shooter data (cf Alcalá et al. 2014), and confirmed by cross-correlation
against the synthetic spectral library of Munari et al. (2005).

The inset in Figure 1 zooms onto the THA 15$-$31 spectrum at the head of the
Balmer series, highlighting two interesting features that are further
discussed in Sect. 4. First, the series of individual Balmer lines is visible
in absorption up to about H(29), corresponding to much higher quantum levels
than those observed in normal stars. Second, the Balmer continuum below
$\lambda\leq 3646$ Å is in absorption. This is a rare occurrence: inspection
of the spectral atlases of SySt by Allen (1984) and Munari and Zwitter (2002),
which survey most ($\sim$150) of the symbiotic stars known at the time,
reveals only about half a dozen objects showing a Balmer continuum in
absorption, whilst the majority of surveyed SySts present the Balmer continuum
in emission, sometimes spectacularly so. The SySts showing the Balmer
continuum in absorption include V4368 Sgr (a symbiotic nova at maximum
brightness when the burning shell around the WD expanded to mimic an A-F
supergiant, Munari 2019), TX CVn (the A/B spectral type of the companion is
explained by Kenyon and Webbink as originating from a low-mass white dwarf
($\sim$0.27 M⊙) with a rejuvenated hydrogen burning shell), and V694 Mon (the
A spectral type of the companion comes from a large accretion disk forming
around a massive ($\geq$0.9 M⊙) white dwarf; Lucy et al. 2020). We believe
that in THA 15-31 the Balmer continuum in absorption originates in the
accretion disk around the companion, as discussed below in sect. 3.7.

###  3.2 Distance

Figure 2: Epoch photometry of THA 15$-$31 in the infrared W1 and W2 bands as
recorded by the NeoWISE mission. The thick red horizontal mark is the
corresponding value listed in the AllWISE catalog (see Sect. 3.3 for details).
Figure 3: ASASSN $V$-band light curve of the blend composed by THA 15$-$31 and
the nearby field star 2MASS 16083631$-$4016120, recorded together as a single
unresolved source.

We mention above how the very small parallax given in Gaia EDR3 indicates a
large distance to THA 15$-$31 (at least several kpc). The associated error is
however too large for a safe inversion of the parallax to derive an accurate
value for the distance (cf. Bailer-Jones 2015; Luri et al. 2018). To estimate
the latter, we compare THA 15$-$31 with SU Lyn, a prototype for accreting-only
symbiotic stars that harbors a similar M6III which suffers from a low
$E_{B-V}$=0.05$\pm$0.02 (Mukai et al. 2016, Lallement et al. 2014, Gontcharov
& Mosenkov 2018). SU Lyn is nearby and therefore its Gaia EDR3 parallax is
accurate at 1.373$\pm$0.063 mas, which fixes the absolute 2MASS $K_{S}$
magnitude of SU Lyn to M($K_{S}$)=$-$7.72$\pm$0.11, where the error propagates
the uncertainties on $K_{S}$, the parallax and the reddening. Assuming that
the M6III giant in THA 15$-$31 is characterized by the same M($K_{S}$) as the
M6III giant in SU Lyn, the distance to THA 15$-$31 turns out to be 12$\pm$0.6
kpc, a value that we adopt in the remainder of this paper. The corresponding
height above the galactic plane is 1.8 kpc, which matches the median value for
known SySts believed to be old-population objects (Munari et al. 2021, their
Figure 7).

###  3.3 Photometric stability

THA 15$-$31 does not appear to be significantly variable at optical or IR
wavelengths. The RG of THA 15$-$31 can be probed with the epoch IR photometry
performed by the Wide-field Infrared Survey Explorer (WISE) over its four
photometric bands W1, W2, W3, and W4 (3.4, 4.6, 12 and 22 $\mu$m,
respectively). While observations in W3 and W4 bands could be performed only
during the initial cryogenic phase before the depletion of hydrogen coolant
(the first 10 months following launch in 2009), a four-month mission extension
was subsequently conducted in the remaining W1 and W2 bands. The AllWISE
catalog (Cutri et al. 2013) combines observations from these 2009-2010
cryogenic and post-cryogenic survey phases, while NeoWISE (Mainzer et al.
2011, 2014) refers to the data the satellite is collecting in W1 and W2 bands
since it was brought out of hibernation and resumed observations in 2014. The
long-term light curve of THA 15$-$31 in the W1 and W2 bands is plotted in
Figure 2. The IR brightness of THA 15$-$31 appears to have remained relatively
constant over the monitored 2009-2010 and 2014-2020 time intervals. A Fourier
analysis suggests the presence of a sinusoidal signal of $\sim$0.06 mag semi-
amplitude and $\sim$1400 day period in the W2 data, albeit at low statistical
significance. However, the currently available NeoWISE data do not yet cover
two full cycles of this possible long periodicity. The W1 data are affected by
errors larger than the amplitude of such a possible periodicity, and cannot be
used to confirm that it is real rather than an artefact.

If the RG dominates the flux at longer wavelengths, the relative contribution
of HC grows moving blueward and optical photometry is strongly influenced by
it (cf. spectral fitting in Figure 1 and the SED discussed in Sect. 3.4
below). We therefore inspected public databases of patrol all-sky surveys in
search for epoch optical photometry of THA 15$-$31\. The object is
unfortunately located too far south for the ZTF survey (Masci et al. 2019,
Bellm et al. 2019), which is conducted with the Palomar Schmidt telescope and
has a deep limiting magnitude and a sharp angular resolution. On the other
hand, the ASASSN survey (Shappee et al. 2014, Kochanek et al. 2017) covers the
whole sky nightly in $V$ or $g$ bands, but the large adopted pixel size (7.8
arcsec) means that THA 15$-$31 cannot be resolved from a nearby field star of
comparable magnitude located 11 arcsec away (2MASS 16083631$-$4016120), and
the two stars are measured together as an unresolved pair by ASASSN. Gaia EDR3
reports $B_{P}$=15.18 and 16.08 mag for THA 15$-$31 and the nearby star,
respectively. The ASASSN light curve for this blended source is presented in
Figure 3 and is characterized by a limited dispersion in brightness (rms=0.126
mag) around a mean value ($V$=14.76 mag) that could be somewhat fluctuating
over a long timescale (by $\sim$0.2 mag). How much of this rather limited
variability is real and intrinsic to THA 15$-$31 is hard to say, and it could
be due to the field star, or instead to the blending of two stellar images
into one. Certainly, variable seeing and focusing can be expected to affect
the result of the aperture-photometry performed by ASASSN on such an
unresolved but obviously not-single source.

In summary, recorded photometric data suggest that THA 15$-$31 does not appear
to vary significantly, to the point that combining multi-wavelength data that
have been gathered at different epochs into a single SED seems a reasonable
procedure for deriving some basic properties of the system.

###  3.4 Spectral energy distribution

Figure 4: Spectral energy distribution of THA 15$-$31 (black dots and
black/green line) corrected for $E_{B-V}$=0.38 reddening and compared to that
of a normal, single M6III star (red dots and line).

The reddening-corrected SED of THA 15$-$31 is presented in Figure 4, where it
is compared to that of an M6III giant built from the intrinsic colors listed
by Bessell (1990), Koornneef et al. (1983), Lee (1970), and Wu et al. (1982).
The absence of significant variability by either the RG or the HC, as
discussed in the previous section, allows us to build the SED of THA 15$-$31
by combining non-simultaneous data from different sources: AllWISE W1, W2, W3,
and W4 magnitudes (Cutri et al. 2013), 2MASS JHKs (Cutri et al. 2003), APASS
DR8 BVg’r’i’ (Henden & Munari 2014), SkyMapper z’ (Wolf et al. 2018), and the
Swift UVOT UVW1, UVM2, and UVW2 data from our observations as described in
Sect. 2 above.

Figure 5: THA 15$-$31 in the near infrared. Top and central panels: THA
15$-$31 (black line) is overplotted to Figure 4 from Akras et al. (2019a).
Bottom panel: Energy distribution of THA 15$-$31 fitted with a single
blackbody in the same fashion as Figure 3 of Akras et al. (2019a). See
discussion in sect. 3.4.

The SED in Figure 4 shows that THA 15$-$31 matches well a field M6III in the
red and in the NIR, confirming the results from the spectral fitting in Figure
1. Integrating the flux distribution under the red line in Figure 4 for the
M6III giant in THA 15$-$31 returns a luminosity of

$L_{\rm M6III}=6400\leavevmode\nobreak\ L_{\odot},$ (2)

corresponding to a bolometric magnitude Mbol=$-$4.75 which closely matches the
expectation for a M6III (cf. Straižys 1992, Sowell et al. 2007).

The contribution by HC to the SED of THA 15$-$31 is rather obvious in Figure
4, with the observed flux distribution floating well above that of the RG in
the optical and UV. The difference for $\lambda\leq$ 7000 Å between the
interpolating black and red lines in Figure 4 corresponds to a radiated
luminosity of

$L_{\rm HC}=100\leavevmode\nobreak\ L_{\odot}.$ (3)

A value frequently cited in the literature for other SySts is the $L_{\rm UV}$
luminosity radiated over the three UVOT photometric bands (cf Mukai et al.
2016): for THA 15$-$31 this value is 6.3$\times$1034 erg s-1 or

$L_{\rm UV}=17\leavevmode\nobreak\ L_{\odot},$ (4)

which places THA 15-31 among the acc-SySts accreting at high rates. Lopes de
Oliveira et al. (2018) found $L_{\rm UV}$ to vary between 0.3 and 2.8 L⊙ for
SU Lyn, and its median value for the accreting SySts studied by Luna et al.
(2013) is 5 L⊙, while Luna et al. (2018) derived $L_{\rm UV}$=25 L⊙ for RT
Cru. Flickering in the UV is ubiquitous among acc-SySts (Luna et al. 2013),
but there is far less flickering at optical wavelengths (Zamanov et al. 2017),
with its detection requiring the greatest care (Munari et al. 2021). The acc-
SySts with the largest-amplitude flickering are generally those with the
highest $\dot{M}_{\rm acc}$, and this bodes well for a search in THA 15-31.

The SED in Figure 4 shows an excess emission over the WISE W3 and W4 bands,
which is probably caused by dust. Integrating the difference between the red
and green lines in Figure 4 for $\lambda\geq$ 5 $\mu$m allows us to derive the
luminosity of this IR excess as

$L_{\rm dust}=35\leavevmode\nobreak\ L_{\odot}.$ (5)

Such dust does not add to the extinction of interstellar origin affecting THA
15$-$31 as defined in Sect. 3.1 above, meaning it either does not project over
the M6III star or, if it engulfs the central binary, it is located at a
sufficiently large distance to dilute the dust grains below the column density
required to impact the observed reddening.

An energy distribution fitted by a stellar photosphere over the entire 1$-$4
$\mu$m range (bands $J$ through $L$) is labeled $S$-type following the
terminology introduced by Allen & Glass (1974), as opposed to $D$-type when
emission from dust at 800-1000 K dominates instead. The majority of known
SySts are of the $S$-type, while SySts of the $D$-type usually harbor an RG
affected by Mira pulsating variability. Later, Allen (1982) introduced a third
class, the $D^{\prime}$-type, where emission from dust at 400-500 K causes an
excess in $L$-band (3.5 $\mu$m) while $J$$H$$K$ bands remain dominated by the
stellar photosphere.

A fourth class, termed S+IR, was recently introduced by Akras et al. (2019a,
hereafter Ak19) and is characterized by a large excess emission detectable
only over the 12$-$22 $\mu$m range covered by the WISE W3 and W4 bands. Ak19
compiled an updated catalog of known SySts and classified them into the four
types: $S$, S+IR, $D^{\prime}$, and $D$. Their criteria would have grouped THA
31-15 with the $S$-type SySts, as illustrated in Figure 5 where our program
star is overplotted on the original Ak19 Figure 4 and its IR energy
distribution is fitted in the same fashion as their Figure 3. Nonetheless, THA
15-31 presents a clear excess emission in W3 and W4 bands, and an association
with the pure $S$-type objects would under-represent this. To classify SySts
in IR, Ak19 fitted their energy distributions with simple blackbodies: one for
the stellar photosphere, another for the dust. The latter was only required
when the excess caused by dust was too large to be accommodated within a
stretching of the fit with just the stellar-photosphere blackbody. Thus, in
the Ak19 reference frame, the transition from pure $S$ to S+IR type occurs
only for really large excess emission at W3 and W4 bands: a moderate excess as
displayed by THA 15-31 is not distinguished from a pure stellar photosphere.
THA 15-31 is the first identified example of a symbiotic star with a moderate
amount of excess in W3 and W4 bands, a transition object between SySts of
really pure $S$-type and the S+IR ones. It would be fair to expect that other
SySts similar to THA 15-31 exist among those grouped into the $S$-type by
Ak19.

We conclude this section by noting that the X-ray nondetection of THA 15-31 is
easily explained by its inferred distance. Indeed, the $\sim$1500 s
integration we performed with Swift/XRT, which implies an X-ray luminosity
upper limit of $<$1.1$\times$1033 ($<$1.4$\times$1033 unabsorbed) in the
0.3–10 keV band, would not have led to the X-ray detection of SU Lyn (the
proto-type of the low-$L_{\rm acc}$ SySts; Mukai et al. 2016) or EG And
(representing the high-$L_{\rm acc}$ SySts; Nuñez et al. 2016) if they were at
the distance of THA 15$-$31, which is approximately 20 times more distant than
the two above examples, thus reducing the observed flux by a factor of 400 and
bringing it below the sensitivity threshold of XRT.

###  3.5 Emission lines

Figure 6: Line profiles for THA 15$-$31\. Top and central rows: Profiles for a
sample of the emission lines recorded on the X-Shooter spectrum of THA
15$-$31\. Bottom row: Some additional spectral features from the same spectrum
and discussed in the text. All panels are plotted over the same velocity
scale. The thin vertical lines mark the laboratory wavelength in the rest
frame of the red giant.

A sample of the recorded emission line profiles of THA 15-31 is presented in
Figure 6: Balmer lines show broad wings and a sharp absorption superimposed on
the emission component, while all other lines are well fitted by a single
Gaussian in emission. The P-Cyg profile for HeI 10830 is only apparent and is
not genuine: the absorption to the left is a normal photospheric line from the
M6III star (SiI 10827 Å; cf. Vaughan & Zirin 1968). Table 2 lists the radial
velocity and integrated flux of observed emission lines.

The profile of Balmer lines was de-convolved with a Voigt profile for the
emission and a Gaussian for the absorption, and the resulting fit is over-
plotted on the observed profile in Figure 6. The principal parameters of the
fitting are listed in Table 3. Given the weakness of the emission component of
H$\gamma$ and H$\delta$, the corresponding values in Table 3 are much more
uncertain than for H$\alpha$ and H$\beta$. There is an obvious progression
moving up along the Balmer series: the absorption component is on the blue
side of the emission in H$\alpha$ and progressively turns to the red side for
higher terms. This behavior is not generally observed in SySts, for which the
difference in velocity between emission and absorption components usually
remains the same for different Balmer lines. However, we sometimes observe a
pattern in SySts that is similar to that displayed by THA 15-31; for example
as displayed by CH Cyg and Hen 2-468 on the spectra we collected on July 24,
2016, with the Asiago 1.82m+Echelle telescope. We do not have a simple
explanation at hand for such behavior. In most SySts, the absorption component
is seen to follow the orbital motion of the RG with an offset of $-$10/$-$20
km s-1, suggesting it forms in the wind that blows off the RG and engulfs the
whole binary. In any case, the radial velocities reported in Table 2 for the
emission lines of THA 15-31 other than Balmer suggest that they form in a
stratified medium: the mean velocity for HeI singlets is +7.5$\pm$0.5 and for
triplets is +20.5$\pm$1 km s-1; it is $-$3$\pm$1 km s-1 for FeII multiplet
#42, and +12.5$\pm$1 km s-1 for multiplet #49.

Table 2: Heliocentric wavelength of line photo-center, corresponding heliocentric radial velocity, and integrated flux for emission lines identified in the X-Shooter spectrum of THA 15-31. Pa-$\alpha$ is overly affected by saturated telluric absorptions, preventing a more meaningful measurement. line |  | $\lambda\odot$ |  RV⊙ | flux
---|---|---|---|---
|  | (Å) |  (km s-1) |  (erg cm-2 s-1)
HeI | sing | 4922.068 | 8.5 | 1.33e-15
HeI | ” | 5015.770 | 5.7 | 1.42e-15
HeI | ” | 6678.328 | 8.0 | 1.81e-15
HeI | ” | 20581.80 | 7.6 | 1.54e-14
HeI | trip | 5876.053 | 20.5 | 5.53e-15
HeI | ” | 7065.802 | 22.3 | 8.75e-15
HeI | ” | 10830.85 | 18.8 | 9.58e-14
FeII | #42 | 4923.877 |  $-$2.7 | 1.67e-15
FeII | ” | 5018.407 |  $-$1.6 | 2.36e-15
FeII | ” | 5168.946 |  $-$4.9 | 2.02e-15
FeII | #49 | 5234.800 | 10.3 | 5.74e-16
FeII | ” | 5276.258 | 15.0 | 5.46e-16
FeII | ” | 5316.825 | 12.2 | 2.11e-15
H$\delta$ |  | 4101.289 |  $-$32.8 | 1.17e-15
H$\gamma$ |  | 4340.415 |  $-$3.7 | 5.84e-15
H$\beta$ |  | 4861.565 | 14.4 | 3.12e-14
H$\alpha$ |  | 6563.379 | 25.7 | 3.38e-13
Pa-$\gamma$ |  | 10938.65 | 15.2 | 9.17e-15
Pa-$\beta$ |  | 12818.63 | 12.9 | 3.17e-14
Pa-$\alpha$ |  | present |  |  $>$1.2e-13
Table 3: Parameters of the fitting to the Balmer lines in Figure 5 with one Gaussian for the absorption and one Voigt profile for the emission. For the latter, the FWHM of the Gaussian core and of the Lorentzian wings are listed. line | emission |  | absorption
---|---|---|---
|  RV⊙ | FWHM | FWHM |  |  RV⊙ | FWHM
|  | core | wings |  |  |
H$\alpha$ | 18 | 145 | 225 |  |  $-$1 | 54
H$\beta$ | 12 | 105 | 160 |  | 12 | 71
H$\gamma$ | 1: | 112: | 150 |  | 11 | 65
H$\delta$ | 4: | 86: | 145 |  | 13 | 69

The lack of essential information for this object, such as orbital period,
orbital inclination, or the presence of flickering, prevents us from
speculating as to where and how the emission lines form in THA 15-31, if in
the accretion disk, in the ionized wind of the M6III, or both. If they form
primarily in the main body of an accretion disk, the single-Gaussian shape of
the nonBalmer lines argues in favor of a disk seen at low inclination (face-
on). A high inclination disk (edge-on) could be accommodated if, at the time
of X-Shooter observation, the view to a significant fraction of the disk was
obstructed (as, e.g., during ingress or egress from an eclipse behind the
M6III giant). On the other hand, the presence of a fast-developing flickering
could argue in favor of the presence of a luminous hot spot at the outer rim
of the disk, and this could be a relevant site for emission-line formation (as
is inferred from the so called ”S-wave” orbital modulation observed in
cataclysmic variables; e.g., Hellier 2001). The presence of a hot spot would
in turn imply that the M6III likely fills its Roche lobe, and mass transfer
proceeds primarily via overflow at L1 with the resulting stream impacting the
disk at the location of the hot spot. The higher and harder the luminosity
radiated by the HC, the larger the fraction of the RG wind it can ionize; in
turn this would shift the primary source of emission-line formation from the
disk toward the ionized wind. The orbital modulation of the radial velocity
and of the line profiles would be different in the two cases, and long-term,
high-resolution spectroscopic monitoring would be clearly beneficial in
identifying the location of formation of the emission lines and in
constraining other basic parameters of THA 15-31.

As a final comment, we note that the H$\alpha$ and H$\beta$ emission lines in
Figure 6 would have well satisfied the selection criteria in Munari et al.
(2021) if THA 15-31 were in GALAH, in addition to the spectral signatures of
an M-giant and the absence of radial pulsations in the light curve. The
profile of THA 15-31 H$\alpha$ is a fine match to that of the 33 GALAH SySts,
and the condition on equivalent widths EW(H$\alpha$)$>$EW(H$\beta$) is met
within a large margin. The EW of H$\alpha$ and H$\beta$ is one order of
magnitude larger in THA 15-31 than in the 33 GALAH SySts, and the luminosity
radiated by the respective accreting stars scales by the same proportion.

###  3.6 Other spectral features

The bottom row of Figure 6 shows the X-Shooter spectrum at some interesting
wavelengths. The CaII line is well centered at the M6III radial velocity, with
a clean profile. Considering that CaII and FeII have a similar ionization
potential, this suggests that FeII emission lines do not form in or near the
M6III atmosphere. On the other hand, the NaI doublet is instead blueshifted
(by $-$9.7$\pm$0.6 and $-$10.4$\pm$0.6 km s-1 for D1 and D2 components,
respectively). Considering that (1) the equivalent width of the interstellar
component corresponding to $E_{B-V}$=0.38 reddening is rather large at 0.53 Å
(from the calibration by Munari and Zwitter 1997), and that (2) the general
radial velocity of the interstellar medium is expected to be negative along
the line of sight to THA 15-31 (Brand & Blitz 1993), the observed $-$10 km s-1
radial velocity for NaI could result from the unresolved blend between
interstellar and stellar components. The HeII panel in Figure 6 clearly shows
the absence of this line, the hallmark of high-ionization conditions in
symbiotic stars. The ZrO panel is centered at the head of the deepest
absorption band from this molecule, which is usually the strongest among
$s$-element oxides in optical spectra (Keenan 1954); its absence suggests the
M6III giant has not (yet) undergone the third dredge-up and has not been
polluted in $s$-elements by the progenitor of the current WD companion (cf.
Jorissen 2003 for extrinsic $s$-element stars among SySts). Finally, no LiI is
present in the spectrum of THA 15-31.

Table 4: Summary of the basic parameters of THA 15$-$31. | alt.name | Sz105 | long (Gal) | 338.724
---|---|---|---|---
RA | (J2000) | 16:08:36.90 | lat (Gal) | +08.516
DEC | (J2000) |  $-$40:16:20.4 | dist | 12 kpc
APASS | B | 16.35 |  $z$-height | 1.8 kpc
| V | 15.06 | $E_{B-V}$ | 0.38
Gaia | G | 12.85 | spec. type | M6 III
|  BP | 15.17 |  RV⊙(M6III) |  4 km s-1
|  RP | 11.46 |  LM6III |  6400 L⊙
2MASS | J | 9.02 |  LHC |  $\geq$100 L⊙
| H | 8.01 |  Ldust |  $\geq$35 L⊙
|  KS | 7.59 |  M(VS) |  $-$1.65
UVOT | UVW1 | 17.86 |  M(KS) |  $-$7.65
| UVM2 | 19.10 | types | acc-only
| UVW2 | 18.86 |  | dusty

###  3.7 The source of UV excess

Although quite relevant, the $L_{\rm HC}$=100 L⊙ luminosity estimated above
for the hot component of THA 15-31 is too low to be ascribed to stable nuclear
burning on the surface of a WD, a condition believed to hold for a significant
fraction of known SySts (Kenyon 1986, Sokoloski 2003). The conditions for
stable and steady H-burning at the surface of an accreting WD have been
explored, among others, by Nomoto et al. (2007), Shen & Bildsten (2007), and
Wolf et al. (2013). In these models, the lowest burning luminosities,
$L_{burn}$=1740 L⊙, pertain to the lightest WD models of 0.51 M⊙ accreting at
2.5$\times$10-8 M⊙ yr-1, while for a WD of 1.34 M⊙ accreting at
1.1$\times$10-6 M⊙ yr-1 the burning luminosity reaches the upper limit of
$L_{burn}$=81,000 L⊙. A similar argument would hold against residual shell-
burning from a previous nova outburst, a scenario also contrasting with the
stable brightness seen when comparing the present-day THA 15-31 to the Palomar
DSS-1 plates of the 1950s.

Alternative to nuclear burning on a WD, the UV excess observed in THA 15-31
may originate from ionized circumstellar material; a normal stellar companion;
or an accretion disk. We consider them in turn, and conclude in favor of the
latter.

The inset of Figure 1 shows a Balmer continuum in marked absorption. Ionized
gas glowing under recombination would inevitably present the Balmer continuum
in emission (Osterbrock & Ferland 2006), not in absorption: this suffices to
exclude nebular material —such as an ionized fraction of the wind from the
M6III— as the primary source of the UV excess in THA 15-31.

Several reasons also argue against a normal star as the source of UV excess.
The emission lines of HeI observed in THA 15-31 originate from upper levels
$\sim$23 eV above the ground state. This and the lack of HeII emission lines
suggests excitation temperatures in the range 2.5$-$5.0$\times$104 K. For
normal stars, such a range corresponds to O and early-B spectral types, for
which the radiated luminosity (103 to 104 $L_{\odot}$; Drilling & Landolt
2000) appears too large in comparison to the 100 L⊙ of the HC in THA 15-31.
Another argument against a normal star is the Balmer series visible in
absorption up to H(29) in the inset of Figure 1. It is well known (e.g.,
Inglis & Teller 1939) that the quantum number of the last term observed in the
Balmer series relates to the electronic density; in practical terms (cf.
Jaschek & Jaschek 1987; Munari et al. 2005), the last Balmer line observable
is $\sim$H(16) for hot main sequence stars and $\sim$H(19) for their high-
luminosity counterparts. H(18) or H(19) are also the last term visible in
emission in nonbinary Mira variables (A. Siviero, priv. comm.). The H(29)
recorded in THA 15-31 suggests an electronic density much lower than
encountered in the stellar atmospheres of normal stars, and even more so for
the high-surface-gravity OB subdwarfs.

Further evidence against the hypothesis that THA 15-31 is a normal star is the
energy distribution of the hot component in Figure 4, which cannot be
approximated by any stellar model for O- or B-type stars or by a black body at
a given temperature: the observed shape resembles instead a gray-body
distribution, of the type that may originate from the combined emission of the
anuli of an accretion disk, each emitting at a different temperature (cf. La
Dous 1989). Finally, binaries composed of cool giants and hot normal star
companions are named VV Cep variables, from the prototype (Pantaleoni González
et al. 2020); as discussed in Munari et al. (2021) such objects are rather
massive and very young, with all known members of the class laying close to
the plane of the Galaxy, with a median height of just 50pc. The much greater
height of THA 15-31, 1.8 kpc from the plane, is incompatible with the young
age of the O- and B-type stars there.

We are then left with an accretion disk as the source for the UV excess
observed in THA 15-31. In support of this, we note that on the Asiago Echelle
spectra that we regularly obtain on V694 Mon, the highest Balmer line visible
in absorption is either H(28) or H(29), similarly to THA 15-31. V694 Mon is a
well-known symbiotic star that attracted much attention in 1990 when, for some
months, it displayed short-living and distinct broad absorption lines
distributed in velocity up to 6000 km/s (Tomov 1990). There is a general
consensus that V694 harbors an accreting WD, with its massive disk dominating
at optical and UV wavelengths (Shore et al. 1994, Zamanov et al. 2011, Lucy
etal. 2020), and confirmed by its large flickering activity (Michalitsianos et
al. 1993, Tomov et al. 1996) and the characteristic hard X-ray emission from
the boundary layer (Stute & Sahai 2009).

Figure 7: Examples of fitting the SED of THA 15-31 combining the M6III with an
accretion disk (and its central object) around a neutron star (top), a white
dwarf (center), and a main sequence star (bottom). Details are discussed in
Sect. 3.8. The arrow on the left marks the upper limit in X-rays from Swift-
XRT observations. The reference distribution for the M6III is the same as that
used in Figure 1.

###  3.8 The accreting star

Having established an accretion disk as the likely source of the UV excess
emission, the next step is to constrain the nature of the central object. We
lack crucial observations at wavelengths intermediate between Swift UVOT and
XRT intervals, which would be of high diagnostic value here. Nonetheless,
available data allow us set useful constraints.

We computed accretion disk models to fit the UV excess observed in THA 15-31
for different choices of the central object, always of the nonmagnetic
variety: a neutron star, a white dwarf, and a main sequence star. The models
are adapted from the implementation presented in Mucciarelli et al. (2007) and
Patruno & Zampieri (2008), with the disk that extends from the surface of the
central object up to twice the circularization radius, and includes self-
irradiation. Examples of the resulting SED are presented in Figure 7.

The computed luminosity and energy distribution depend on a number of
parameters, including: the masses of the central and the donor stars, the
binary period, the accretion rate, the SED of the donor, in addition to the
inclination angle and the orbital phase of the binary system. For THA 15-31,
most of them are unknown, and reasonable ranges have been extensively
explored. The results presented in Figure 7 are typical of each class, and can
therefore be regarded as guidelines for these three configurations: their
energy distribution invariably peaks in the X-rays for a neutron star, in the
extreme UV for a white dwarf, and in the near-UV for a main sequence star.

At the top of Figure 7, we consider the case of a neutron star of 1.5 M⊙ and
accreting at 7.5e-10 M⊙/yr. Whatever the mass of the neutron star or the
accretion rate, fitting the flux in the Swift UVOT bands always produces an
X-ray emission far in excess of the upper limit observed in THA 15-31 with
Swift XRT. Introducing a magnetic field for the neutron star would only
exacerbate the problem, channeling the inner parts of the disk onto the polar
caps and increasing the X-ray flux with the emission from the accretion
columns. Therefore, we can confidently exclude a neutron star as the accreting
object in THA 15-31.

At the bottom of Figure 7, we explore a main sequence (MS) star as the central
object, with a mass of 0.8 M⊙ and accreting at 1e-05 M⊙/yr. As already noted
by Kenyon and Webbink (1984), workable MS models for SySts invariably require
furious mass-accretion rates, and our simulations certainly confirm this. The
very shallow potential well of an MS star, whilst easily accounting for the
absence of any X-ray emission, requires very high $\dot{M}_{\rm acc}$ to
produce enough flux to fit THA 15-31 in the UV. The temperature attained by
the main body of the disk is generally too low to support photoionization at
energies larger than for hydrogen. Harder photons can be produced at the
boundary layer, but the high optical thickness of such a disk suggests that
they are all absorbed locally with few escaping beyond the boundary layer
itself. The modeling of the boundary layer under the very large accreting
rates required by an MS is rather uncertain, and a simple scaled-up version of
the formalism generally adopted for WD or NS could be misleading. For such a
reason, we have not included a boundary layer in the disk model for the MS
case at the bottom of Figure 7.

At the center of Figure 7, we finally consider a WD as the accreting object.
In the plotted case, the WD has a mass of 1.0 M⊙ and accretes at 7.5e-08
M⊙/yr. The WD is slowly rotating, has a radius of 0.008 $R_{\odot}$, and the
accretion efficiency and the albedo of the disk and donor surface layers are
fixed at $2.6\times 10^{-4}$ and 0.9, respectively. The mass adopted for the
M6III is 1.2 M⊙, 1000 days for the orbital period, and $i=60^{\circ}$ for the
orbital inclination. With such a configuration, the Roche lobe of the donor is
comparable to the radius of the M6III giant ($\sim 200$-$300$ R⊙). The
bolometric luminosity of the disk is $150\,L_{\odot}$, consistent with the
estimate reported in equation (3). The emission from the disk well accounts
for the Swift/UVOT flux and peaks in the extreme UV, providing a significant
reservoir of ionizing photons for the formation of the hydrogen emission lines
(there are also enough $\lambda$$\leq$228 Å photons to expect detection of
HeII 4686. The spectra of THA 15-31 do not show this line, but they were taken
several years before the acquisition of the Swift observations modeled here,
and changes in the accretion rate between the two epochs may explain the
difference). At the same time, the spectrum cuts off well below the soft
X-rays, which is consistent with the lack of detection with Swift/XRT. We did
not perform a fit because it would not be particularly constraining for the
accretion disk (considering that most of the SED is dominated by the red
supergiant). However, varying ${\dot{M}}$, $i,$ and the albedo, we verified
that the observed fluxes can be satisfactorily reproduced in a range of values
close to those reported above, which shows that the result is relatively
robust. In particular, a value of $\dot{M}_{\rm acc}$=2e-08 M⊙ is also
acceptable for a face-on system ($i\sim 0^{\circ}$) with an albedo of 0.8. We
note that the accretion rate is above Eddington, but the bolometric disk
luminosity is not because the efficiency is low.

The mass-loss rate of isolated, nonpulsating M giants is generally assumed to
be a few 10-7 M⊙ yr-1 (e.g., Livio & Warner 1984). The above requirements on
$\dot{M}_{\rm acc}$ for the case of an accreting WD in THA 15-31 would be
easily met by conventional Roche-lobe overflow or by the wind Roche-lobe
overflow suggested by the hydrodynamical simulations of Mohamed &
Podsiadlowski (2012). If the orbital separation in THA 15-31 is too large for
the RG to approach its Roche-lobe, the accretion needs to proceed via wind-
intercept. Assuming the wind leaves the RG and travels distances longer than
the orbital separation in a spherically symmetric arrangement, the standard
Bondi-Hoyle treatment for wind-intercept (Livio & Warner 1984) shows that the
companion captures only a few percent of the material lost by the RG, too low
a fraction to account for the high $L_{\rm acc}$ observed in THA 15-31. This
fraction needs to be one order of magnitude higher than allowed by the Bondi-
Hoyle scenario, a condition frequently encountered in SySts (e.g., Skopal
2005). This can be explained by the companion’s gravitational focusing of the
wind leaving the RG toward the orbital plane (e.g., Decin et al. 2019). The
latter results in a strong density enhancement on the orbital plane at the
expense of the polar directions: the higher density favors a higher accretion
by the companion compared to the case of an unfocussed, spherically symmetric
wind outflowing from the RG. The presence of a strong density enhancement on
the equatorial plane of SySts is expected theoretically (Mohamed &
Podsiadlowski 2007, Booth et al. 2016), and was confirmed by the marked
bipolar shape observed for the spatially resolved nova ejecta of RS Oph
(Ribeiro et al. 2009) and V407 Cyg (Giroletti et al. 2020), and inferred by
modeling of line profiles in high-inclination systems like EG And by Shagatova
et al. (2016, 2021).

In conclusion, an accretion disk around a WD provides the simplest and easiest
explanation for the UV excess seen in THA 15-31, with a large space of
parameters able to reproduce observations, and with the required mass-loss
rate from the cool giant being a fine match to normal values. While there is
no room for a NS in THA 15-31, a disk around an MS star could still fit the
UVOT observations but only while invoking an accretion rate on the MS star of
at least two orders of magnitude larger than the mass-loss rate typical of
nonpulsating M giants. A rate of 1e-05 M⊙/yr is achieved only at the tip of
the AGB, where it may be supported by radial pulsations (Höfner & Olofsson
2018). These evolved giants are usually enshrouded in massive circumstellar
cocoons abundant in dust, and may show over-abundance in $s$-element following
third dredge-up episodes. Such an AGB is unlikely present in THA 15-31: no
radial pulsations, or large dust excess, or even $s$-element overabundance is
observed, and as noted in Figure 1 the spectral match with normal M6III giants
is instead perfect. Therefore, while the presence of an MS star in THA 15-31
cannot be excluded with absolute certainty, it would require that such an odd
series of conditions be met that it is (highly) unlikely.

The brightness of THA 15-31 on the blue plates of the first Palomar all-sky
photographic atlas was similar to current values, suggesting it was accreting
to a comparable rate during the 1950s. It would be interesting to reconstruct
its long-term light curve, for example from inspection of archive photographic
plates, and ascertain for how long the current $\dot{M}_{\rm acc}$ has been
maintained. The presence or absence of past, unnoticed outbursts could then be
compared with the stability of accretion disks in SySts over time, and the
amount of material piled-up on the companion with the amount necessary to
ignite nuclear burning on the surface of a WD.

Finally, the previous misclassification of THA 15-31 as a young stellar object
suggests that other accreting-only symbiotic stars may be lurking in old YSO
catalogs. Such misclassification is not unprecedented: a significant fraction
of symbiotic stars of the burning type were originally discovered and
classified as bona fide planetary nebulae: it was only the discovery at later
times and on IR wavelengths of the presence of a cool giant in the system that
led to rectification of their initial classification (e.g., Allen 1983, Acker
& Stenholm 1990). In regard to this, we find that the IR selection criteria
for SySts, derived by Akras et al. (2019b) following a machine learning
approach, nicely distinguish THA 15-31 from true YSOs in the list of H$\alpha$
emission stars for the Lupus and Scorpius regions compiled by The (1962).

###### Acknowledgements.

We would like to thank the anonymous Referee for constructive suggestions that
led to improve the paper. A. Skopal provided useful comments on an early
version of this paper, and A. Siviero let us to access unpublished data from
his atlas of high resolution spectra of Mira variables. We thank Jamie Kennea
and the Swift team for the quick approval and the rapid acquisition of the
observations we requested. NM acknowledges financial support through ASI-INAF
and ’Mainstream’ agreement 2017-14-H.0 (PI: T. Belloni). JMA and AF
acknowledge financial support from the project PRIN-INAF 2019
Spectroscopically Tracing the Disk Dispersal Evolution (STRADE). This work was
supported by the Swedish strategic research programme eSSENCE. GT was
supported by the project grant ‘The New Milky Way’ from the Knut and Alice
Wallenberg Foundation and by the grant 2016- 03412 from the Swedish Research
Council. GT also acknowledges financial support of the Slovenian Research
Agency (research core funding no. P1-0188 and project N1-0040). This
publication makes use of data products from the Near-Earth Object Wide-field
Infrared Survey Explorer (NEOWISE), which is a joint project of the Jet
Propulsion Laboratory/California Institute of Technology and the University of
Arizona. NEOWISE is funded by the National Aeronautics and Space
Administration.

## References

  * Acker & Stenholm (1990) Acker A., Stenholm B., 1990, A&AS, 86, 219
  * Akras et al. (2019a) Akras S., Guzman-Ramirez L., Leal-Ferreira M. L., Ramos-Larios G., 2019a, ApJS, 240, 21
  * Akras et al. (2019b) Akras S., Leal-Ferreira M. L., Guzman-Ramirez L., Ramos-Larios G., 2019b, MNRAS, 483, 5077. doi:10.1093/mnras/sty3359
  * Alcalá et al. (2011) Alcalá J. M., Stelzer B., Covino E., Cupani G., Natta A., Randich S., Rigliaco E., et al., 2011, AN, 332, 242. doi:10.1002/asna.201111526
  * Alcalá et al. (2014) Alcalá J. M., Natta A., Manara C. F., Spezzi L., Stelzer B., Frasca A., Biazzo K., et al., 2014, A&A, 561, A2. doi:10.1051/0004-6361/201322254
  * Alcalá et al. (2017) Alcalá J. M., Manara C. F., Natta A., Frasca A., Testi L., Nisini B., Stelzer B., et al., 2017, A&A, 600, A20. doi:10.1051/0004-6361/201629929
  * Allen & Glass (1974) Allen D. A., Glass I. S., 1974, MNRAS, 167, 337. doi:10.1093/mnras/167.2.337
  * Allen (1982) Allen D. A., 1982, ASSL, 27. doi:10.1007/978-94-009-7834-8_4
  * Allen (1983) Allen D. A., 1983, MNRAS, 204, 113. doi:10.1093/mnras/204.1.113
  * Allen (1984) Allen D. A., 1984, PASAu, 5, 369
  * Bailer-Jones (2015) Bailer-Jones C. A. L., 2015, PASP, 127, 994. doi:10.1086/683116
  * Belczyński et al. (2000) Belczyński K., Mikołajewska J., Munari U., Ivison R. J., Friedjung M., 2000, A&AS, 146, 407
  * Bellm et al. (2019) Bellm E. C., Kulkarni S. R., Graham M. J., Dekany R., Smith R. M., Riddle R., Masci F. J., et al., 2019, PASP, 131, 018002. doi:10.1088/1538-3873/aaecbe
  * Bessell (1990) Bessell M. S., 1990, PASP, 102, 1181. doi:10.1086/132749
  * Booth, Mohamed, & Podsiadlowski (2016) Booth R. A., Mohamed S., Podsiadlowski P., 2016, MNRAS, 457, 822. doi:10.1093/mnras/stw001
  * Brand & Blitz (1993) Brand J., Blitz L., 1993, A&A, 275, 67
  * Breeveld et al. (2011) Breeveld A. A., Landsman W., Holland S. T., Roming P., Kuin N. P. M., Page M. J., 2011, AIPC, 1358, 373, AIPC.1358
  * Buder et al. (2021) Buder S., Sharma S., Kos J., Amarsi A. M., Nordlander T., Lind K., Martell S. L., et al., 2021, MNRAS, 506, 150. doi:10.1093/mnras/stab1242
  * Burrows et al. (2005) Burrows D. N., et al., 2005, SSRv, 120, 165
  * Capitanio et al. (2017) Capitanio L., Lallement R., Vergely J. L., Elyajouri M., Monreal-Ibero A., 2017, A&A, 606, A65
  * Chen et al. (1997) Chen H., Grenfell T. G., Myers, P. C., Hughes J. D., 1997, ApJ, 478, 295. doi:10.1086/303769
  * Corradi et al. (2010) Corradi R. L. M., Valentini M., Munari U., Drew J. E., Rodríguez-Flores E. R., Viironen K., Greimel R., et al., 2010, A&A, 509, A41. doi:10.1051/0004-6361/200913231
  * Cutri et al. (2003) Cutri R. M., et al., 2003, yCat, 2246,
  * Cutri & et al. (2013) Cutri R. M., et al., 2013, yCat, 2328
  * Decin et al. (2019) Decin L., Homan W., Danilovich T., de Koter A., Engels D., Waters L. B. F. M., Muller S., et al., 2019, NatAs, 3, 408. doi:10.1038/s41550-019-0703-5
  * De Silva et al. (2015) De Silva G. M., et al., 2015, MNRAS, 449, 2604
  * Drilling & Landolt (2000) Drilling J. S., Landolt A. U., 2000, asqu.book, 381
  * Fitzpatrick (1999) Fitzpatrick E. L., 1999, PASP, 111, 63. doi:10.1086/316293
  * Fluks et al. (1994) Fluks M. A., Plez B., The P. S., de Winter D., Westerlund B. E., Steenman H. C., 1994, A&AS, 105, 311
  * Frasca et al. (2006) Frasca A., Guillout P., Marilli E., Freire Ferrero R., Biazzo K., Klutsch A., 2006, A&A, 454, 301. doi:10.1051/0004-6361:20054573
  * Frasca et al. (2017) Frasca A., Biazzo K., Alcalá J. M., Manara C. F., Stelzer B., Covino E., Antoniucci S., 2017, A&A, 602, A33. doi:10.1051/0004-6361/201630108
  * Fujimoto (1982) Fujimoto M. Y., 1982, ApJ, 257, 767
  * Gaia Collaboration (2020) Gaia Collaboration, 2020, yCat, I/350
  * Ghez et al. (1997) Ghez A. M., McCarthy D. W., Patience J. L., Beck T. L., 1997, ApJ, 481, 378. doi:10.1086/304031
  * Giroletti et al. (2020) Giroletti M., Munari U., Körding E., Mioduszewski A., Sokoloski J., Cheung C. C., Corbel S., et al., 2020, A&A, 638, A130. doi:10.1051/0004-6361/202038142
  * Gómez & Persi (2002) Gómez M., Persi P., 2002, A&A, 389, 494. doi:10.1051/0004-6361:20020601
  * Gontcharov & Mosenkov (2018) Gontcharov G. A., Mosenkov A. V., 2018, yCat, II/354
  * Green et al. (2019) Green G. M., Schlafly E., Zucker C., Speagle J. S., Finkbeiner D., 2019, ApJ, 887, 93. doi:10.3847/1538-4357/ab5362
  * Inglis & Teller (1939) Inglis D. R., Teller E., 1939, ApJ, 90, 439. doi:10.1086/144118
  * Jaschek & Jaschek (1990) Jaschek C., Jaschek M., 1990, The Classification of Stars, Cambridge Univ. Press
  * Jorissen (2003) Jorissen A., 2003, ASPC, 303, 25
  * Hachenberg & Wellmann (1939) Hachenberg O., Wellmann P., 1939, ZA, 17, 246
  * Hellier (2001) Hellier C., 2001, Cataclysmic Variable Stars, Springer
  * Henden & Munari (2014) Henden A., Munari U., 2014, CoSka, 43, 518
  * Hill et al. (2004) Hill J.E., et al., 2004, Proc. SPIE, 5165, 217
  * Höfner & Olofsson (2018) Höfner S., Olofsson H., 2018, A&ARv, 26, 1. doi:10.1007/s00159-017-0106-5
  * Hughes et al. (1994) Hughes J., Hartigan P., Krautter J., Kelemen J., 1994, AJ, 108, 1071. doi:10.1086/117135
  * Keenan (1954) Keenan P. C., 1954, ApJ, 120, 484. doi:10.1086/145937
  * Kenyon & Webbink (1984) Kenyon S. J., Webbink R. F., 1984, ApJ, 279, 252. doi:10.1086/161888
  * Kenyon (1986) Kenyon S. J., 1986, The Symbiotic Stars, Cambridge Univ. Press
  * Kirk & Myers (2011) Kirk H., Myers P. C., 2011, ApJ, 727, 64. doi:10.1088/0004-637X/727/2/64
  * Kochanek et al. (2017) Kochanek C. S., Shappee B. J., Stanek K. Z., Holoien T. W.-S., Thompson T. A., Prieto J. L., Dong S., et al., 2017, PASP, 129, 104502. doi:10.1088/1538-3873/aa80d9
  * Koornneef (1983) Koornneef J., 1983, A&A, 500, 247
  * Krautter et al. (1997) Krautter J., Wichmann R., Schmitt J. H. M. M., Alcala J. M., Neuhauser R., Terranegra L., 1997, A&AS, 123, 329. doi:10.1051/aas:1997163
  * La Dous (1989) La Dous C., 1989, A&A, 211, 131
  * Lallement et al. (2014) Lallement R., Vergely J.-L., Valette B., Puspitarini L., Eyer L., Casagrande L., 2014, A&A, 561, A91
  * Livio & Warner (1984) Livio M., Warner B., 1984, Obs, 104, 152
  * Lopes de Oliveira et al. (2018) Lopes de Oliveira R., Sokoloski J. L., Luna G. J. M., Mukai K., Nelson T., 2018, ApJ, 864, 46. doi:10.3847/1538-4357/aad2d5
  * Lee (1970) Lee T. A., 1970, ApJ, 162, 217. doi:10.1086/150648
  * Lucy et al. (2020) Lucy A. B., Sokoloski J. L., Munari U., Roy N., Kuin N. P. M., Rupen M. P., Knigge C., et al., 2020, MNRAS, 492, 3107. doi:10.1093/mnras/stz3595
  * López Martí, Jiménez-Esteban, & Solano (2011) López Martí B., Jiménez-Esteban F., Solano E., 2011, A&A, 529, A108. doi:10.1051/0004-6361/201016377
  * Luna et al. (2013) Luna G. J. M., Sokoloski J. L., Mukai K., Nelson T., 2013, A&A, 559, A6. doi:10.1051/0004-6361/201220792
  * Luna et al. (2018) Luna G. J. M., Mukai K., Sokoloski J. L., Lucy A. B., Cusumano G., Segreto A., Jaque Arancibia M., et al., 2018, A&A, 616, A53. doi:10.1051/0004-6361/201832592
  * Luna et al. (2020) Luna G. J. M., Sokoloski J. L., Mukai K., M. Kuin N. P., 2020, ApJL, 902, L14. doi:10.3847/2041-8213/abbb2c
  * Luri et al. (2018) Luri X., Brown A. G. A., Sarro L. M., Arenou F., Bailer-Jones C. A. L., Castro-Ginard A., de Bruijne J., et al., 2018, A&A, 616, A9. doi:10.1051/0004-6361/201832964
  * Mainzer et al. (2011) Mainzer A., et al., 2011, ApJ, 731, 53
  * Mainzer et al. (2014) Mainzer A., et al., 2014, ApJ, 792, 30
  * Marshall et al. (2006) Marshall D. J., Robin A. C., Reylé C., Schultheis M., Picaud S., 2006, A&A, 453, 635. doi:10.1051/0004-6361:20053842
  * Masci et al. (2019) Masci F. J., Laher R. R., Rusholme B., Shupe D. L., Groom S., Surace J., Jackson E., et al., 2019, PASP, 131, 018003. doi:10.1088/1538-3873/aae8ac
  * Masetti et al. (2007a) Masetti N., Rigon E., Maiorano E., Cusumano G., Palazzi E., Orlandini M., Amati L., et al., 2007a, A&A, 464, 277. doi:10.1051/0004-6361:20066517
  * Masetti et al. (2007b) Masetti N., Landi R., Pretorius M. L., Sguera V., Bird A. J., Perri M., Charles P. A., et al., 2007b, A&A, 470, 331. doi:10.1051/0004-6361:20077509
  * Merc et al. (2019) Merc J., Gális R., Wolf M., 2019, AN, 340, 598. doi:10.1002/asna.201913662
  * Michalitsianos et al. (1993) Michalitsianos A. G., Perez M., Shore S. N., Maran S. P., Karovska M., Sonneborn G., Webb J. R., et al., 1993, ApJL, 409, L53. doi:10.1086/186858
  * Modigliani et al. (2010) Modigliani A., Goldoni P., Royer F., Haigron R., Guglielmi L., François P., Horrobin M., et al., 2010, SPIE, 7737, 773728. doi:10.1117/12.857211
  * Mohamed & Podsiadlowski (2007) Mohamed S., Podsiadlowski P., 2007, ASPC, 372, 397
  * Mohamed & Podsiadlowski (2012) Mohamed S., Podsiadlowski P., 2012, BaltA, 21, 88. doi:10.1515/astro-2017-0362
  * Mucciarelli et al. (2007) Mucciarelli, P., Zampieri, L., Treves, A., et al. 2007, ApJ, 658, 999. doi:10.1086/512152
  * Mukai et al. (2016) Mukai K., et al., 2016, MNRAS, 461, L1
  * Munari & Zwitter (1997) Munari U., Zwitter T., 1997, A&A, 318, 269
  * Munari & Zwitter (2002) Munari U., Zwitter T., 2002, A&A, 383, 188. doi:10.1051/0004-6361:20011724
  * Munari et al. (2005) Munari U., Sordo R., Castelli F., Zwitter T., 2005, A&A, 442, 1127. doi:10.1051/0004-6361:20042490
  * Munari, Dallaporta, & Cherini (2016) Munari U., Dallaporta S., Cherini G., 2016, NewA, 47, 7
  * Munari (2019) Munari, U., in The Impact of Binary Stars on Stellar Evolution, G. Beccari and H.M.J. Boffin eds., 2019, Cambridge Univ. Press,p. 77 (arXiv:190901389)
  * Munari et al. (2021) Munari U., Traven G., Masetti N., Valisa P., Righetti G.-L., Hambsch F.-J., Frigo A., et al., 2021, MNRAS, 505, 6121. doi:10.1093/mnras/stab1620
  * Mužić et al. (2014) Mužić K., Scholz A., Geers V. C., Jayawardhana R., López Martí B., 2014, ApJ, 785, 159. doi:10.1088/0004-637X/785/2/159
  * Nomoto et al. (2007) Nomoto K., Saio H., Kato M., Hachisu I., 2007, ApJ, 663, 1269. doi:10.1086/518465
  * Nucita et al. (2014) Nucita A. A., Stefanelli S., De Paolis F., Masetti N., Ingrosso G., Del Santo M., Manni L., 2014, A&A, 562, A55. doi:10.1051/0004-6361/201322680
  * Nuernberger, Chini, & Zinnecker (1997) Nuernberger D., Chini R., Zinnecker H., 1997, A&A, 324, 1036
  * Nuñez et al. (2016) Nuñez N. E., Nelson T., Mukai K., Sokoloski J. L., Luna G. J. M., 2016, ApJ, 824, 23. doi:10.3847/0004-637X/824/1/23
  * Osterbrock & Ferland (2006) Osterbrock D. E., Ferland G. J., 2006, Astrophysics of gaseous nebulae and active galactic nuclei, University Science Books
  * Pantaleoni González et al. (2020) Pantaleoni González M., Maíz Apellániz J., Barbá R. H., Negueruela I., 2020, RNAAS, 4, 12. doi:10.3847/2515-5172/ab712b
  * Patruno & Zampieri (2008) Patruno, A. & Zampieri, L. 2008, MNRAS, 386, 543. doi:10.1111/j.1365-2966.2008.13063.x
  * Poole et al. (2008) Poole T. S., et al., 2008, MNRAS, 383, 627
  * Ribeiro et al. (2009) Ribeiro V. A. R. M., Bode M. F., Darnley M. J., Harman D. J., Newsam A. M., O’Brien T. J., Bohigas J., et al., 2009, ApJ, 703, 1955. doi:10.1088/0004-637X/703/2/1955
  * Roming et al. (2005) Roming P. W. A., et al., 2005, SSRv, 120, 95
  * Sale et al. (2014) Sale S. E., Drew J. E., Barentsen G., Farnhill H. J., Raddi R., Barlow M. J., Eislöffel J., et al., 2014, MNRAS, 443, 2907. doi:10.1093/mnras/stu1090
  * Schwartz (1977) Schwartz R. D., 1977, ApJS, 35, 161. doi:10.1086/190473
  * Shappee et al. (2014) Shappee B. J., Prieto J. L., Grupe D., Kochanek C. S., Stanek K. Z., De Rosa G., Mathur S., et al., 2014, ApJ, 788, 48. doi:10.1088/0004-637X/788/1/48
  * Shagatova, Skopal, & Cariková (2016) Shagatova N., Skopal A., Cariková Z., 2016, A&A, 588, A83. doi:10.1051/0004-6361/201525645
  * Shagatova et al. (2021) Shagatova N., Skopal A., Shugarov S. Y., Komžík R., Kundra E., Teyssier F., 2021, A&A, 646, A116. doi:10.1051/0004-6361/202039103
  * Shen & Bildsten (2007) Shen K. J., Bildsten L., 2007, ApJ, 660, 1444. doi:10.1086/513457
  * Shore, Aufdenberg, & Michalitsianos (1994) Shore S. N., Aufdenberg J. P., Michalitsianos A. G., 1994, AJ, 108, 671. doi:10.1086/117103
  * Skopal (2005) Skopal A., 2005, A&A, 440, 995. doi:10.1051/0004-6361:20034262
  * Sokoloski (2003) Sokoloski J. L., 2003, ASPC, 303, 202
  * Sowell et al. (2007) Sowell J. R., Trippe M., Caballero-Nieves S. M., Houk N., 2007, AJ, 134, 1089. doi:10.1086/520060
  * Starrfield et al. (2020) Starrfield S., Bose M., Iliadis C., Hix W. R., Woodward C. E., Wagner R. M., 2020, ApJ, 895, 70. doi:10.3847/1538-4357/ab8d23
  * Straižys (1992) Straižys V., 1992, Multicolor Stellar Photometry, Tucson: Pachart Pub. House
  * Stute & Sahai (2009) Stute M., Sahai R., 2009, A&A, 498, 209. doi:10.1051/0004-6361/200811176
  * The (1962) The P.-S., 1962, CoBos, 15, 0
  * Tomov et al. (1990) Tomov T., Kolev D., Georgiev L., Zamanov R., Antov A., Bellas Y., 1990, Natur, 346, 637. doi:10.1038/346637a0
  * Tomov et al. (1996) Tomov T., Kolev D., Ivanov M., Antov A., Jones A., Mikolajewski M., Lepardo A., et al., 1996, A&AS, 116, 1
  * Vaughan & Zirin (1968) Vaughan A. H., Zirin H., 1968, ApJ, 152, 123. doi:10.1086/149531
  * Yamashita, Nariai & Norimoto (1977) Yamashita Y., Nariai K., Norimoto Y., 1977, An Atlas of Representative Stellar Spectra, Univ. of Tokyo Press
  * Wolf et al. (2013) Wolf W. M., Bildsten L., Brooks J., Paxton B., 2013, ApJ, 777, 136. doi:10.1088/0004-637X/777/2/136
  * Wolf et al. (2018) Wolf C., Onken C. A., Luvaul L. C., Schmidt B. P., Bessell M. S., Chang S.-W., Da Costa G. S., et al., 2018, PASA, 35, e010. doi:10.1017/pasa.2018.5
  * Wu et al. (1982) Wu C. C., Boggess A., Bohlin R. C., Holm A. V., Schiffer F. H., Turnrose B. E., 1982, The IUE Spectral Atlas, ESASP, 182, 25
  * Zamanov et al. (2011) Zamanov R., Boeva S., Latev G., Stoyanov K., Bode M. F., Antov A., Bachev R., 2011, IBVS, 5995, 1
  * Zamanov et al. (2017) Zamanov R. K., Boeva S., Nikolov Y. M., Petrov B., Bachev R., Latev G. Y., Popov V. A., et al., 2017, AN, 338, 680. doi:10.1002/asna.201713362