# The GALAH Survey and Symbiotic Stars. I. Discovery and follow-up of 33
candidate accreting-only systems

U. Munari,1 G. Traven,2 N. Masetti,3,4 P. Valisa,5 G.-L. Righetti,5 F.-J.
Hambsch,5 A. Frigo,5 K. Čotar,6 G. M. De Silva,7,8 K. C. Freeman,9 G. F.
Lewis,10 S. L. Martell,11,12 S. Sharma,10,12 J. D. Simpson,11,12 Y.-S. Ting,13
R. A. Wittenmyer14 D. B. Zucker,8,12,15
1INAF Astronomical Observatory of Padova, 36012 Asiago (VI), Italy
2Lund Observatory, Department of Astronomy and Theoretical Physics, Box 43,
SE-221 00 Lund, Sweden
3INAF Osservatorio di Astrofisica e Scienza dello Spazio, via Gobetti 93/3,
40129 Bologna, Italy
4Departamento de Ciencias Físicas, Universidad Andrés Bello, Fernández Concha
700, Las Condes, Santiago, Chile
5ANS Collaboration, c/o Astronomical Observatory, 36012 Asiago (VI), Italy
6Faculty of Mathematics and Physics, University of Ljubljana, Jadranska 19,
1000 Ljubljana, Slovenia
7Australian Astronomical Optics, Macquarie University 105 Delhi Rd, North
Ryde, NSW 2113, Australia
8Macquarie University Research Centre for Astronomy, Astrophysics &
Astrophotonics, Sydney, NSW 2109, Australia
9Research School of Astronomy & Astrophysics, Australian National University,
ACT 2611, Australia
10Sydney Institute for Astronomy, School of Physics, A28, The University of
Sydney, NSW 2006, Australia
11School of Physics, UNSW, Sydney, NSW 2052, Australia
12Centre of Excellence for Astrophysics in Three Dimensions (ASTRO-3D),
Australia
13Department of Astrophysical Sciences, Princeton University, Princeton, NJ
08544, USA
14University of Southern Queensland, Centre for Astrophysics, USQ Toowoomba,
QLD 4350 Australia
15Department of Physics and Astronomy, Macquarie University, Sydney, NSW 2109,
Australia
E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

We have identified a first group of 33 new candidates for symbiotic stars
(SySt) of the accreting-only variety among the 600 255 stars so far observed
by the GALAH high-resolution spectroscopic survey of the Southern Hemisphere,
more than doubling the number of those previously known. GALAH aims to high
latitudes and this offers the possibility to sound the Galaxy for new SySt
away from the usual Plane and Bulge hunting regions. In this paper we focus on
SySt of the M spectral type, showing an H$\alpha$ emission with a peak in
excess of 0.5 above the adjacent continuum level, and not affected by coherent
radial pulsations. These constraints will be relaxed in future studies. The 33
new candidate SySt were subjected to a vast array of follow-up confirmatory
observations (X–ray/UV observations with the Swift satellite, search for
optical flickering, presence of a near-UV upturn in ground-based photometric
and spectroscopic data, radial velocity changes suggestive of orbital motion,
variability of the emission line profiles). According to Gaia $e$DR3
parallaxes, the new SySt are located at the tip of the Giant Branch, sharing
the same distribution in M(KS) of the well established SySt. The accretion
luminosities of the new SySt are in the range 1$-$10 L⊙, corresponding to
mass-accretion rates of 0.1-1 10-9 M⊙ yr-1 for WDs of 1 M⊙. The M giant of one
of the new SySt presents a large Lithium over-abundance.

###### keywords:

binaries: symbiotic stars – methods: data analysis – Galaxy: stellar content

††pubyear: 2020††pagerange: The GALAH Survey and Symbiotic Stars. I. Discovery
and follow-up of 33 candidate accreting-only systems–A

##  1 Introduction

Symbiotic stars (SySt) are interacting binaries where a red giant (RG) fuels a
white dwarf (WD) or a neutron star (NS) companion via accretion (either
through Roche-lobe overflow or wind intercept). Systems harboring NSs are a
recent addition to this class of celestial objects (Masetti et al., 2006,
2007a, 2007b, 2011; Bozzo et al., 2018; Yungelson, Kuranov, & Postnov, 2019)
and constitute a few percent of the known total (Merc, Gális, & Wolf, 2019a),
the vast majority containing WDs (Mukai et al., 2016; Sokoloski et al., 2017;
Akras et al., 2019). For some time, a main sequence star accreting from the RG
at very high rates (10-6 to 10-4 M⊙ yr-1) was considered a viable scenario for
SySt in general (Kenyon & Webbink, 1984), but this has been progressively
abandoned in the light of expanding observational data, especially in the far
UV and X-rays (Sokoloski et al., 2006; Skopal et al., 2006; Mohamed &
Podsiadlowski, 2012; Luna et al., 2013; Skopal & Cariková, 2015; Lutovinov et
al., 2020). Even if we will focus on RG+WD systems, the results of this paper
are independent on the actual nature of the accreting star.

The RG+WD symbiotic stars are broadly divided into two major groups (see the
recent review by Munari et al. 2019 for details): those accreting-only (acc-
SySt) whose optical spectra are dominated by the RG with no or weak emission
lines, and the burning-type (burn-SySt) displaying a strong nebular continuum
and a rich emission line spectrum: they 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.

SySt are believed to spend most of their time in the accreting-only phase,
quietly accumulating material on the surface of the WD. When enough has been
piled up, nuclear burning ignites. If the accreted matter is electron
degenerate, the burning proceeds explosively. It quickly reaches the Fermi
temperature (a matter of minutes, Starrfield et al. 2008) and ejects most of
the accreted envelope at high velocity (thousands of km s-1). Radio-
interferometry has nicely resolved the expansion of ejecta following, for ex.,
the 2006 outburst of RS Oph (O’Brien et al., 2006) and that of 2010 for V407
Cyg (Giroletti et al., 2020), the latter being also the first nova event ever
detected in GeV $\gamma$-rays by the Fermi satellite (Cheung et al., 2010).
The resulting nova outburst can repeat on time-scales as short as
years/decades if the WD mass is close to the Chandrasekhar limit, like in the
case for T CrB, RS Oph, V407 Cyg or V3890 Sgr. If the matter accreted on the
WD is instead not electron-degenerate, the nuclear burning proceeds in thermal
equilibrium, and no mass is ejected. The event takes a few years to reach peak
brightness (Fujimoto, 1982), and requires many decades to a few centuries to
burn the accreted envelope and let the system return to low luminosity. Some
of the best examples of this type are AG Peg, HM Sge, V1016 Cyg and V4368 Sgr.

There is a clear disproportion among catalogued SySt in favor of the burn-SySt
type, and the known examples of the acc-SySt variety are believed to be just
the tip of the iceberg (Mukai et al. 2016, hereafter Mk16). Between 15 and 20
are currently known of the acc-SySt type. Symbiotic stars were originally
proposed by Munari & Renzini (1992) as possible progenitors of type Ia
supernovae, and the viability of this single-degenerate channel obviously
relies on the total number of SySt spread throughout the Galaxy, which in turn
heavily depends on the number of acc-SySt. The latter have been usually
discovered as counterparts of satellite UV/X-ray sources, while the burn-SySt
can be easily spotted at optical wavelengths through the whole Galaxy and the
Local Group thanks to their outstanding emission-line spectrum (with
ionization up to [FeX] and beyond).

The subtle way the optically-quiet acc-SySt have been usually discovered is
well epitomized by SU Lyn, a $V$$\sim$8 mag, M6III giant at 650$\pm$35 pc
distance (Gaia Collaboration et al., 2018). For decades it was a completely
unnoticed field star, with just an old report about semi-regular photometric
variability which granted it a variable star name (Kippenhahn, 1955). While
looking for optical counterparts of hard X-ray sources newly discovered by the
Swift satellite, it was noted that SU Lyn lay within the error box of one of
them: 4PBC J0642.9+5528. Follow-up observations were organized with Swift (to
refine the position and the properties of the X-ray source) and with Asiago
telescopes to investigate if the optical spectra of SU Lyn could betray
peculiarities supporting a physical association with the Swift hard X–ray
source. The association was proved by a large flux excess observed at bluest
optical wavelengths ($\lambda$$\leq$4000 Å) and by the presence on high-
resolution spectra of weak and variable emission in H$\alpha$ and [NeIII]
lines (cf. Mk16).

The limited sensitivity of current X-ray satellites (soon to change thanks to
$e$-ROSITA sky survey, Merloni et al. 2012) restricts the serendipitous
discovery of acc-SySt to those that lie within $\sim$1 kpc from the Sun,
implying that many of their red giants rank among naked-eye objects (like 4
Dra, HR 1105 or $o$ Cet). To sample a much larger fraction of the RG in our
Galaxy, we have devised a reverse strategy, and first exploratory results are
presented in this paper: digging through large spectroscopic surveys in search
for the signatures of accretion onto a companion to the RG, primarily the
presence of Balmer emission lines having a flux and a profile compatible with
an origin in an accretion disk around a degenerate companion. The candidate
acc-SySt identified this way can then be subjected to follow-up confirmatory
observations, including pointing with X–ray/UV satellites. The GALAH survey of
the southern hemisphere (De Silva et al., 2015) offers an ideal hunting
opportunity: a large number (aiming at $>$1 million) of randomly selected
southern stars in the $12\leq V_{JK}\leq 14$ mag range are observed at high
S/N and high spectral resolution with the AAT 3.9m telescope + fibre-fed
spectrograph HERMES, over four distinct wavelength ranges that include
H$\alpha$ and H$\beta$ lines.

In this first paper we outline the methodologies guiding our search among
GALAH spectra for candidate acc-SySt, report on the discovery of a first batch
of 33 new candidates (about doubling the total known so far), and describe the
results of a wide range of follow-up observations aiming to confirm their acc-
SySt nature. We encourage further observations by the community in order to
validate their classification.

##  2 The GALAH survey

###  2.1 Motivation

The GAlactic Archaeology with HERMES (GALAH, De Silva et al. 2015) is an
ongoing spectroscopic survey whose ambitious goal is to unveil the formation
and history of the the Milky Way. This is the focus of the field of Galactic
Archaeology, which tries to determine how the fossil remnants of star-forming
regions and effects of ancient mergers paint the picture of our Galaxy that we
observe and measure today. This complex endeavor can be accomplished by
studying the detailed chemical composition and other properties of stars in
distinct regions of the Milky Way.

The Galactic archaeology community argues that the complete Galaxy formed
gradually over time. This formation history can be traced back by
investigating remnants of initial building blocks or subsequent additions
through mergers, which have been disrupted in the course of evolution and are
now dispersed around the Galaxy. The theoretical concept of chemical tagging
(Freeman & Bland-Hawthorn, 2002) demonstrates that individual galactic
components should have preserved their original chemical signature over time.
It is therefore essential to disentangle their formation site from migration
history in order to explain the current mixture of stellar populations.

GALAH aims to achieve this by measuring the abundance of up to 31 chemical
elements coming from seven independent major groups with different
nucleosynthetic origin: light proton-capture elements Li, C, O;
$\alpha$-elements Mg, Si, Ca, Ti; odd-Z elements Na, Al, K; iron-peak elements
Sc,V, Cr, Mn, Fe, Co, Ni, Cu, Zn; light and heavy slow neutron capture
elements Rb, Sr, Y, Zr, Ba, La; and rapid neutron capture elements Ru, Ce, Nd
and Eu (De Silva et al., 2015).

###  2.2 Instrument

The goals of the GALAH survey were the main driver for the construction of the
High Efficiency and Resolution Multi-Element Spectrograph (HERMES, Barden, et
al. (2010); Sheinis et al. (2015)), a multi-fibre spectrograph working in
tandem with the $3.9$-metre Anglo-Australian Telescope (AAT) situated at the
Siding Spring Observatory, Australia. The spectrograph has a resolving power
of R $\sim 28,000$ (or R $\sim 45,000$ when slit mask is used) and records
spectra in four separate wavelength ranges given in Table 1. The total
spectral coverage is therefore approximately 1000 Å, including the important
diagnostic lines H$\alpha$ and H$\beta$, together with lines of numerous
chemical elements that are necessary to fulfil the ambitions of the GALAH
survey.

The AAT uses the Two Degree Field (2dF) robotic positioning system with two
identical plates that are used to precisely position fibres at designated
locations. This configuration allows HERMES to simultaneously record spectra
from up to 392 fibres distributed over a $2^{\circ}$ diameter field of the
night sky, with an additional 8 fibres used for the telescope guiding. During
the exposure with the first plate, the robotic positioner places fibres on the
second plate, where each complete process of fibre allocation takes about half
an hour per plate. HERMES can typically achieve a signal to noise ratio (SNR)
$\sim 100$ per resolution element at magnitude V=14 in the red arm during a
1-hour long exposure. To achieve as high SNR as possible, minimize atmospheric
diffraction, and in order not to loose light due to the fibres’ field of view
of only 2″, all observations are ideally carried out close to the meridian.

###  2.3 Selection function

The main selection function of the GALAH survey is relatively simple: it
avoids the crowded regions of the Galactic plane ($|b|$ > $10^{\circ}$),
accounts for the nominal telescope operations
($-80^{\circ}\leq\delta\leq+10^{\circ}$) and restricts observations by
apparent magnitude ($12<V_{JK}<14$), while being colour independent. There is
an additional requirement for field selection: the density of stars has to be
at least 400 per $\pi$ square degrees to match the number of fibres and field
of view of the fibre positioner. The input catalogue for observations is based
on the Two Micron All-Sky Survey (2MASS Skrutskie et al., 2006), with $V_{JK}$
approximating standard $V$ magnitude and being computed as (for more details
see Martell et al., 2017):

$V_{JK}=K+2(J-K+0.14)+0.382\mathrm{e}^{(J-K-0.2)/0.5}$ (1)

The majority of observations are performed with the above described selection
function. Additionally, GALAH employs a bright mode ($9<V_{JK}<12$) during
twilight or poor observing conditions, and a faint mode ($12<V_{JK}<14.5$)
when fields with normal or bright configuration are not available to be
observed.

The spectroscopic data used in this work is further complemented by
observations of the K2-HERMES survey Wittenmyer et al. (2018) and the TESS-
HERMES survey Sharma et al. (2018), which are carried out with the same
observational and data reduction setup as the GALAH survey, albeit with their
peculiar selection functions focusing on K2 (Howell et al., 2014) and TESS
(Ricker et al., 2015) targets, respectively.

Together, the above described observations yield a dataset of 625 757
successfully reduced stellar spectra, of which a small fraction belongs to
repeated observations of 600 255 unique stars observed by GALAH at the time of
writing of this paper. They represent dwarf as well as giant stars, the former
consisting of mostly nearby stars (closer than 2 kpc) while the latter probing
distances as far as the Galactic center.

Table 1: Short and long wavelength limits (Å) of the four spectral intervals recorded by the GALAH survey. interval | start | end
---|---|---
blue | 4718 | 4903
green | 5649 | 5873
red | 6481 | 6739
far red | 7590 | 7890
Figure 1: Two spectra of SU Lyn, a prototype of the accreting-only symbiotic
stars, to illustrate the differences between low (black) and high (red)
accretion rates (Asiago 1.22m telescope). A higher rate causes a brightening
of the accretion disk around the white dwarf, which manifests as an excess
near-UV brightness ($\lambda$$\leq$4000 Å) and detectable emission lines
(primarily hydrogen Balmer lines).

###  2.4 Data reduction and analysis

Every stellar spectrum recorded by the HERMES spectrograph and used in this
work is homogeneously reduced by an automatic reduction pipeline, thoroughly
described by Kos et al. (2017). We hereby provide only a brief explanation of
this reduction procedure.

The path from a 2D image recorded in each of the HERMES arms to a one-
dimensional, continuum-normalized and zero-RV-shifted spectrum starts with the
following steps: raw image cosmetic corrections, spectral tracing, optical
aberrations correction, scattered light and apertures cross-talk removal,
wavelength calibration, sky subtraction, and telluric absorption removal.
Afterwards, spectra are continuum-normalized and shifted into their rest-frame
by cross-correlating them with a set of 15 AMBRE model spectra (de Laverny et
al., 2012). This procedure also yields initial radial velocities, while
initial stellar parameters ($T_{\rm eff}$, $\log g$, $\mathrm{[Fe/H]}$) are
determined by a (larger) grid of $16\,783$ AMBRE spectra. These, internally
also called GUESS parameters, were reported in the first GALAH data release
(Martell et al., 2017). The spectra are then directed to the analysis pipeline
with the aim of delivering precise and accurate stellar atmospheric parameters
and individual elemental abundances, which constitute the core of results
reported in subsequent GALAH data releases.

###  2.5 The new SySt probed by GALAH

Most of the known SySt have been discovered during photographic objective-
prism surveys of the Galactic plane and Bulge (e.g. Merrill & Burwell 1950;
Sanduleak & Stephenson 1973; Henize 1976), many of which were initially
classified and studied as planetary nebulae. Their true nature was generally
unveiled by infrared observations catching the presence of a cool giant in the
system (Allen, 1982, 1984; Sabbadin, Falomo, & Ortolani, 1987). The same
objective-prism technique was applied also to the Magellanic Clouds and
several SySt were discovered there too (Walker, 1983; Morgan, 1992). Searches
for SySt in external galaxies has steadily progressed over the years (Angeloni
et al., 2014; Gonçalves et al., 2008, 2012, 2015; Mikołajewska, Caldwell, &
Shara, 2014; Mikołajewska et al., 2017; Roth et al., 2018; Kniazev et al.,
2009; Drozd et al., 2019), and the current total of known extra-galactic SySt
amounts to $\sim$70 systems (Akras et al., 2019; Chen, Liu, & Shan, 2019;
Merc, Gális, & Wolf, 2019b). New surveys aiming to SySt within the Local Group
(Angeloni et al., 2019) promise to increase such numbers, especially for the
burn-SySt type.

Systematic searches for new SySt have also been performed in recent years
through our Galaxy. The IPHAS survey (Drew et al., 2005) targeted the plane of
the Milky Way for emission-line objects, among which many new candidate SySt
were identified by Corradi et al. (2008, 2010, 2011) and Rodríguez-Flores et
al. (2014). SySt were also hunted toward the Bulge by Miszalski, Mikołajewska,
& Udalski (2013) and Miszalski & Mikołajewska (2014). The current total of
confirmed Galactic SySt is about 250 (Akras et al., 2019; Merc, Gális, & Wolf,
2019b).

No recent survey has extensively probed the Galaxy for SySt toward directions
other than low on the Plane or toward the Bulge. This is a gap we intend to
fill with the help of the GALAH survey, that is exploring the Galaxy primarily
at high latitudes. In this way a more complete census of SySt will be possible
toward all stellar populations, allowing a more informed comparison with the
predictions about the total number of SySt spread throughout the Galaxy. Such
estimates vary over a wide range: from 1,200$-$15,000 of Lü, Yungelson, & Han
(2006), or 3,000$-$30,000 of Yungelson et al. (1995) and Zhu et al. (2010), to
3,000 of Allen (1984), 30,000 of Kenyon et al. (1993), and 300,000 of Munari &
Renzini (1992). The large spread is in part accounted for by differences in
the adopted binary evolution codes and stellar population synthesis, or by the
type of SySt considered (burning or accreting only), and also by the ratio
between SySt currently in an active state (and therefore detectable to
observations) and those dormant (and indistinguishable from single, field cool
giants). The populations of SySt containing a NS is generally estimated at
100$-$1,000 (eg. Lü et al., 2012), or about 1% of those harboring a WD.

Figure 2: Temporal sequence from Asiago 1.82m and Varese 0.84m Echelle spectra
(aligned in heliocentric wavelength) illustrating the great changes observed
in the H$\alpha$ profile of prototype SU Lyn since its recognition in late
2015 as an accreting-only symbiotic star. The vertical line marks the
H$\alpha$ rest wavelength at the stellar systemic velocity ($-$24 km s-1). The
epochs $a$, $b$ and $c$ are discussed in the text (Section 3). On the right
the corresponding $B$,$V$ lightcurve of SU Lyn as built from ANS Collaboration
observations.

##  3 SU Lyncis, a path-finder

As remarked earlier, several acc-SySt have been known for quite a while. A few
of them present emission lines bright enough to be recognizable even on
objective prism plates or low resolution optical spectra. They are well
epitomized by EG And, an M2III giant orbited every 481 days by a WD, whose
accretion luminosity is a few 102 L⊙ (Nuñez et al., 2016).

The discovery of SU Lyn by Mk16 revealed a ”hidden and potentially large
population" of acc-SySt accreting at much lower rates, with accretion
luminosities a couple of orders of magnitude lower than in EG And, and
therefore not detectable by conventional observations at optical wavelengths.
This paper aims to specifically discover more of these SySt characterized by
low accretion luminosities (Lacc$\sim$1-10 L⊙). As the properties of SU Lyn
will guide us as a path-finder throughout this paper, and since no paper has
so far reviewed this object at optical wavelengths, we will start with a quick
overview of spectroscopic and photometric properties of SU Lyn. These are
based on the very intensive monitoring we are keeping of the object since the
preparatory phase leading to the Mk16 discovery paper, and whose results will
be presented and discussed in detail elsewhere. The properties of SU Lyn that
will be reviewed here are those that we will look for in the dataset of GALAH
stars with the aim to discover the new SySt.

The two most obvious distinctive characteristics of optical spectra of SU Lyn
(and acc-SySt of that type) are shown in Fig. 1: ($a$) a near-UV ultraviolet
excess ($\lambda\leq 4000$ Å), and ($b$) weak emission lines, primarily from
the hydrogen Balmer series. They both originate in the accretion disk forming
around the degenerate companion, fed by the material accreted from the mass-
losing RG companion.

Both can vary greatly in time, as the two epochs compared in Fig. 1 illustrate
well: in just over one year SU Lyn passed from an active state (Dec 19, 2015),
when its SySt nature was quite obvious, to a quiet phase (Feb 16, 2017) during
which an optical spectrum would not have betrayed its binary nature. This
needs to be kept in mind when comparing GALAH spectra with ancillary data
taken at very different epochs.

The most important role of SU Lyn in guiding our search among GALAH stars is
however the observed profile for emission lines, H$\alpha$ in particular. Fig.
2 presents a temporal sequence from Asiago 1.82m and Varese 0.84m telescopes +
Echelle spectrographs illustrating the great variety of profiles observed for
H$\alpha$ in this star from 2015 to 2020 (for a more traditional presentation
of some of these same profiles for key epochs see fig. 3 of Munari 2019). The
orbital plane of SU Lyn is probably close to face-on conditions and this
accounts for the negligible changes in radial velocity observed for the M6III
absorption lines ($\leq$2 km s-1). The observed spectroscopic and photometric
changes are therefore not related to changes in the aspect angle of the
binary.

The epoch marked b in Fig. 2 corresponds to the lowest accretion rate and it
is noteworthy for the absence of emission in H$\alpha$ and a null near-UV
excess: at optical wavelengths SU Lyn appears like a normal and single M6III,
not as an interacting binary. Only the reduced equivalent width of the
H$\alpha$ absorption (only about half the photospheric value of typical M6III
giants) and its narrow FWHM (again half the value for normal M6III giants)
could cause doubt. The white vertical line in Fig. 2 represents the
photospheric velocity of the cool giant: on epoch b, the velocity of the
H$\alpha$ absorption almost equals that of the M6III giant. This is not the
case for most of the time, however: as clearly shown by Fig. 2, the H$\alpha$
absorption is normally blue-shifted compared to photospheric absorptions as if
originating in a gentle wind blowing off the inner regions of the accretion
disk (and thus appearing superimposed onto its emission), or forming in the
expanding wind of the RG that engulfs the whole binary system.

Epoch c in Fig. 2 marks the sudden appearance of a second, distinct and faster
moving absorption component, that gradually reduced its velocity during the
following months and finally merged with the slower, pre-existing absorption
component. At the time of maximum equivalent width of this second absorption
component, the emission in H$\alpha$ briefly vanished.

Finally, epoch a represents the condition of highest accretion rate, with a
strong near-UV excess and a prominent and structured emission in H$\alpha$,
with still superimposed the blue-shifted absorption component mentioned above.
It is a lucky circumstance that when we first observed SU Lyn in preparation
of Mk16 paper, the star was exhibiting a clear emission in H$\alpha$ and the
strongest near-UV up-turn: if at the time, it had appeared as in epoch b of
Fig. 2, we could not have recognized it for what it really was.

Figure 3: Definition of TiO classification indexes. The shadowed regions show
the wavelength ranges covered by the GALAH blue, green and red channels. The
yellow and red bands mark the A-B and C-D wavelength intervals, respectively,
used to compute the $b$, $g$ and $r$ ratios according to Eq.(2) and Table 3.
In background some Asiago 1.22m spectra illustrating the progress in strength
of TiO bands moving along the sequence of M giants (pay attention to the log-
scale of the ordinates).

Another important point illustrated by Fig. 2 concerns the photometric
activity of SU Lyn. The star is too bright to be recorded unsaturated by
surveys which monitor the heaven on a nightly basis in search for transients
(eg. ASAS-SN, MASTER, ZTF, etc.), therefore we monitored it ourselves with
small instruments, and fully transformed the observations to the Landolt
(1992) standard UBVRI system. As illustrated by the $B$,$V$ lightcurve on the
right panel of Fig. 2 (synchronized to the same temporal ordinates as the
spectra to the left), SU Lyn varies by up to 0.7 mag amplitude in a random
fashion, with no persistent periodicity. There is no associated change in
colour and no correspondence with the behavior of the accretion disk as traced
by H$\alpha$ intensity and profile. Thus, the variability is not caused by
brightening and fading of the accretion disk (hotter than the M6III), but must
originate with the cool giant. All cool giants exhibit variability to some
extent (Hoffmeister, Richter, & Wenzel, 1985; Sterken & Jaschek, 2005), given
the unstable nature of their convective outer layers (Pugh, Gray, & Griffin,
2015), and their tenuous atmosphere that extends across the temperatures for
molecule formation and through those for dust condensation. The absence of
radial velocity variability on the left panel of Fig. 2 precludes an origin of
the variability seen in SU Lyn with the coherent, large scale, persistent
radial pulsations of the type observed in Mira long-period variables. The
latter show a sinusoid-like radial velocity curve of $\sim$10-15 km s-1 in
amplitude, in phase with the equally sinusoid-like lightcurve that can span up
to 10 magnitude in amplitude, and characterized by bluer colour at maximum and
redder at minimum (see fig. 27 in Hoffmeister, Richter, & Wenzel, 1985).

The distinction with radially-pulsating giants is important: their outer
layers experience large scale motions that lead to shocks, which in turn power
emission lines. In our search for acc-SySt we aim to avoid the contamination
from radially-pulsating giants and their deceiving emission lines. The
selection criteria for removing false positives from the final acc-SySt sample
are outlined in the rest of the paper.

##  4 Selection steps

Our search for acc-SySt among GALAH targets started with isolating the cool
giants from the rest. In this paper we focus on giants of the M spectral
types, those most abundant ($\sim$70%) among the known SySt (see the catalogs
by Allen, 1984, Belczyński et al., 2000, Akras et al., 2019, Merc, Gális, &
Wolf, 2019b). SySt containing other types of cool giants will be explored in
follow-up papers.

###  4.1 Colour

The intrinsic colour of M0 giants in the Solar neighborhood is
$(J-K)_{o}$=0.97 (Lee, 1970; Koornneef, 1983), so we first apply a colour
selection of $(J-K_{s})$$\geq$0.90 to GALAH targets to both account for the
natural spread with the breadth of the M0 spectral type and the slight
difference in wavelength-baseline between the Johnson’s $(J-K)$ and 2MASS
$(J-K_{s})$ indices.

###  4.2 Parallax

On the HR diagram, stars with an M spectral type are either very low-mass main
sequence objects (down to the limit for stable H-burning in the core), or
giants resulting from the evolution of more massive progenitors. The stars in
our Milky Way have not lived long enough to significantly populate the range
in between. M-dwarfs are intrinsically very faint and over the $12<V_{JK}<14$
mag range covered by GALAH targets, they represent the stars closest to the
Sun, while the opposite is certainly true for the giants. In Table 2 we
compare (under negligible reddening) the expected parallax of dwarfs and
giants with an M spectral type at both ends of the GALAH magnitude range. No
giant has a parallax larger than 0.4 mas, or a dwarf smaller than 10mas.
Therefore a second criterion was set imposing a Gaia DR2 (Gaia Collaboration
et al., 2018) parallax $\pi$$\leq$3 mas. This selection criterium is fully
confirmed by the improved Gaia $e$DR3 parallaxes (Gaia Collaboration, 2020),
that became available only after this paper was originally submitted.

Table 2: Parallax for dwarf and giant M stars (under negligible reddening) at the bright ($V$=12) and faint ($V$=14) end of the GALAH magnitude range. | MV | parallax (mas)
---|---|---
|  | V=12 | V=14
M0V | 8.90 | 24 | 10
M5V | 12.30 | 115 | 46
M0III |  $-$0.10 | 0.38 | 0.15
M5III |  $-$0.90 | 0.26 | 0.10

###  4.3 TiO absorption bands

The above colour and parallax criteria combined together returned an initial
sample of 29 514 GALAH stars. This sample could contain hotter giants of G and
K spectral types affected by a large reddening or cool giants of a spectral
type other than M (like the Carbon giants). To prune them out and retain only
those with an M-type spectrum, we looked for signatures of their distinctive
TiO absorption bands in the GALAH spectra. This may have picked-up also some
S-type giants whose spectra, in addition to TiO bands, show also molecules
involving $s$-type elements (ZrO in particular; Turnshek et al., 1985). Some
percentage of known SySt do indeed possess S-type giants (Jorissen, 2003).

Fig. 3 presents a progression of M giant spectra with over-plotted in green
the wavelength range of the blue, green and red GALAH channels. Within their
span lie strong TiO bands, and the later the spectral type, the deeper such
bands become. To confirm the presence of the TiO bands and objectively derive
their depth, we defined eight narrow TiO wavelength bins (marked in red in
Fig. 3) and an equal number of control wavelength bins (marked in yellow in
Fig. 3). The control bins are all placed to the blue of their respective TiO
bins given the fact that absorption bands of the TiO molecule all degrade to
the red and present a steep band-head to the blue (Pearse & Gaydon, 1976). We
then defined eight ratios (b1, b2, b3 in the blue channel, g1, g2 in the
green, and r1 , r2 and r3 in the red) by integrating the flux within the
control bin and dividing by it the TiO bin as:

Table 3: Short and long wavelength limits (Å) of the intervals used to compute the ratios in Eq. (2) and illustrated in Fig. 3. ratio | $\lambda_{A}$ | $\lambda_{B}$ | $\lambda_{C}$ | $\lambda_{D}$
---|---|---|---|---
b1 | 4751 | 4756 | 4765 | 4770
b2 | 4798 | 4803 | 4807 | 4812
b3 | 4835 | 4840 | 4850 | 4855
g1 | 5801 | 5806 | 5812 | 5817
g2 | 5829 | 5834 | 5850 | 5855
r1 | 6641 | 6646 | 6653 | 6658
r2 | 6670 | 6675 | 6684 | 6689
r3 | 6706 | 6711 | 6716 | 6721

${\rm
ratio}=\frac{\int_{\lambda_{A}}^{\lambda_{B}}f(\lambda)d\lambda}{\int_{\lambda_{C}}^{\lambda_{D}}f(\lambda)d\lambda}$
(2)

where the respective wavelength intervals A$-$B and C$-$D are listed in Table
3. These ratios were computed on spectra before continuum normalization but
shifted to rest frame by the radial velocity listed in GALAH DR3 as rv_guess
(Buder et al., 2020). Absence of absorption by TiO returns a value around 1.0
for any of the ratios, and a progressively larger value with increasing M
spectral subtype (M0 $\rightarrow$ M10). Finally, we have combined the eight
ratios into a single one computed as

$f=\frac{[0.3(<b>-1)+1]+<g>+<r>}{3}$ (3)

where $<$b$>$ stands for the arithmetic average of b1, b2 and b3, and
similarly for the rest. The $<$b$>$ average ratio stretches over 3.3$\times$
the range of the other two, so we have consequently reduced its impact to make
it equal to the others. This is justified by the lower S/N of the spectra
recorded in the blue channel compared to green and red ones.

The depth of the TiO bands has been already used in the past to classify the M
giants of symbiotic stars, e.g. by Kenyon & Fernandez-Castro (1987) that used
to this aim the bands at 6180 and 7100 Å (both outside the GALAH wavelength
range). In these studies, equivalent widths of TiO bands were measured over
their entire wavelength span (hundreds of Å), too wide compared to the breadth
of GALAH channels, and therefore inapplicable in the present paper (and for
that matter also to more conventional, high-resolution Echelle spectra
covering too small a wavelength interval over a single order).

Table 4: Intervals in combined $f$-ratio (Eq. 3) defining the template bins for the selected 15 824 M giants. $(J-K_{s})_{\circ}$ is the median value within the bin for the reddening corrected $(J-K_{s})$ from 2MASS, and M($K_{s}$) is the median value of the absolute magnitude computed from Gaia $e$DR3 parallaxes satisfy the condition $\sigma(\pi)$/$\pi$$<$0.1. The last column lists the number of stars in each bin. The template spectra for each bin are plotted in Fig. 16. bin | left limit | right limit | $(J-K_{s})_{\circ}$ | M($K_{s}$) | N. stars
---|---|---|---|---|---
00 | 1.0200 | 1.0303 | 0.904 |  $-$3.78 | 3419
01 | 1.0304 | 1.0419 | 0.935 |  $-$4.01 | 2309
02 | 1.0420 | 1.0561 | 0.968 |  $-$4.27 | 1732
03 | 1.0562 | 1.0736 | 0.994 |  $-$4.58 | 1424
04 | 1.0737 | 1.0950 | 1.017 |  $-$4.81 | 1282
05 | 1.0951 | 1.1214 | 1.043 |  $-$4.99 | 1200
06 | 1.1215 | 1.1539 | 1.065 |  $-$5.18 | 1048
07 | 1.1540 | 1.1938 | 1.088 |  $-$5.42 | 919
08 | 1.1939 | 1.2428 | 1.110 |  $-$5.69 | 917
09 | 1.2429 | 1.3030 | 1.141 |  $-$5.94 | 849
10 | 1.3031 | 1.3770 | 1.166 |  $-$6.12 | 638
11 | 1.3771 | 1.4679 | 1.182 |  $-$6.29: | 77
12 | 1.4680 | 3.6559 | 1.139 |  | 10
Figure 4: Relation between the reddening-corrected 2MASS $(J-K_{s})_{\circ}$
colour index and the combined $f$-ratio from Eq.(2) for the 29 514 GALAH stars
in our initial sample. The coloured bands mark the approximate location of the
MKK spectral subtypes (M0III, M1III, …). The finer, log-scale subdivision
according to Table 4 is given by the horizontal lines and it is marked by the
numbers to the right (00, 01, …).

The way our ratios are defined, makes them insensitive to spectral fluxing,
continuum normalization, or reddening. The relation between $f$-ratio and
reddening corrected $(J-K_{s})_{\circ}$ index for the initial sample of 29 514
GALAH stars is illustrated in Fig. 4. The latter shows a number of highly
reddened stars contaminating the sample. To filter them out we took all stars
laying above the knee at $f$$>$1.02, and considered the remaining 15 824
objects (the 15k-sample for short) as representing the true M-giants present
among the 600 255 stars observed by GALAH at the time of writing of this
paper. The $(J-K_{s})_{\circ}$ colour index has been calculated from 2MASS
photometry and the full extinction as given by Schlafly & Finkbeiner (2011).
The choice of taking the full value of the tabulated extinction is justified
by considering the great distance to our targets (typically several kpc,
beyond which Gaia DR2 or $e$DR3 parallaxes become inaccurate) and their large
galactic latitude, which makes their sight-lines exit the Galactic dust slab
well before reaching them.

###  4.4 H$\alpha$ in emission

The next step has been searching the 15k-sample of true M giants for those
showing H$\alpha$ (and H$\beta$) in emission. All the symbiotic stars
considered in this paper show H$\alpha$ emission directly on the recorded
spectrum, well before any reference spectrum is subtracted (cf. Fig. 5). To
properly reconstruct the profile of the emission line, it is however necessary
to subtract the underlying spectrum of the M giant. To define the latter, we
grouped the M giants by similar strength of TiO bands, and computed the median
of the continuum normalized spectra within each bin. The spectral progression
of the templates around H$\alpha$ and H$\beta$ is compared in Fig. 16 of the
Appendix.

We initially selected to define a unique template for each of the eleven M0
III$\leftrightarrow$M10 III spectral subtypes of the original MKK
classification scheme (Morgan, Keenan & Kellman, 1943; Morgan & Keenan, 1973),
but quickly realized that at the high S/N and high resolution of GALAH
spectra, any such sub-type spans too wide a range in stellar properties for a
single template to adequately represent them all. Similarly, equally spaced
bins in $f$-ratio are not satisfactory, as they oversample at low $f$ values.
Therefore, we decided for a logarithmic progression, with Table 4 listing the
13 final bins, their interval in terms of $f$-ratio, and the number of true M
giants contained in that bin. Their subdivision of the giant branch is shown
in Fig. 4 (horizontal lines). The median of all the spectra in a given bin
defines the template spectrum for that bin. Taking the median value eliminates
any disturbance from cosmic ray hits or rare presence of emission components.
In addition, the median computed over thousands of individual spectra has an
extremely high S/N and it is completely insensitive to chemical peculiarities
of individual stars.

The template is finally subtracted from continuum normalized spectra of the
true M giants within the given bin, and the residuals inspected for emission
in H$\alpha$. In this first paper we limit ourselves to consider only the M
giants showing an H$\alpha$ emission profile rising above a threshold of 0.5
on the subtracted spectrum. This returned a total of 223 stars.

###  4.5 Visual inspection

The spectra of these 223 stars were visually inspected for inconsistencies
that could cause the H$\alpha$ to appear in emission on the subtracted
spectrum.

Some spectra were found to suffer from a mismatch in radial velocity with the
template. In these few cases, the radial velocity rv_guess from GALAH DR3,
which is used to shift spectra in all four GALAH channels to rest frame,
provides correct results for the blue, green, and infrared, but not for the
red channel, resulting in an excess flux around H$\alpha$ in the difference
spectrum, which is obviously spurious.

Another, more frequent cause for false positives was the presence of an
emission component probably originating from the sky background that has not
been fully cancelled-out by the sky-subtraction procedure. This may happen
e.g. when the target star is located in a region of the sky affected by
diffuse background emission, like emission nebulae and HII regions. Such cases
are easy to recognize because the H$\alpha$ and H$\beta$ emission lines are
($a$) typically quite sharp, single-component and symmetric, ($b$) frequently
accompanied by $[$NII$]$ 6548, 6584 nebular emission lines (which do not form
in the high density environment of an accretion disk), and ($c$) usually
displaced in velocity with respect to the M-giant spectrum by an amount larger
than expected from orbital motion in a symbiotic binary. In multi-fibre
spectroscopic surveys like GALAH, only a limited number of fibres are assigned
to record the sky background. They are distributed over the field of view away
from the position of known stars. This is perfectly fine for a great fraction
of the sky away from the Galactic plane. There are however, here and there,
regions of the sky affected by background emission due to diffuse nebulosity.
The intensity of such emission may vary appreciably over limited angular
distances, and the same may be the case for the internal nebular gas dynamics
and therefore the resulting wavelength for H$\alpha$ and H$\beta$ emitted
lines. Any sparse mapping realized by a limited number of sky fibres of such a
complex 2D pattern can naturally result in non-null residuals of the sky-
background subtraction, in addition to other caveats of sky subtraction as
explained in Kos et al. (2017).

There was also a group of stars which presented an equal excess emission in
H$\alpha$ and H$\beta$ on the subtracted spectrum, always with the same
characteristic profile (steep red and blue ends, flat and rounded top), as if
originating from a hydrogen deficiency compared with the template spectrum.
These stars were also disregarded in the current selection process.

After a further few rejections for miscellaneous reasons, we were left with a
final selection of 83 M giants showing genuine emission in H$\alpha$ with an
intensity in excess of 0.5 in the subtracted spectrum (the 83-sample
hereafter). Stars with an emission in H$\alpha$ weaker than 0.5, which are
inherently more subtle to treat, will be investigated in a follow-up paper.

###  4.6 Filtering out the radial pulsators

Not all the objects in the 83-sample are necessarily valid acc-SySt. The main
false-positives are expected to be the (large-amplitude) radial pulsators
(e.g. Miras or SRa variables), that may give origin to emission lines
(primarily hydrogen Balmer) deep in their atmosphere where shocks form.

It is of course not precluded that a Mira pairs with a compact companion to
form a symbiotic star. For example V407 Cyg, prior to its 2010 nova outburst,
was primarily studied as a Mira suffering from possible dust-obscuration
episodes (Munari, Margoni, & Stagni, 1990; Kolotilov et al., 1998) with little
(or none) spectroscopic evidences for binarity (Hinkle et al., 2013). Mira
itself (= $o$ Cet) is a well known acc-SySt: Hubble in the ultraviolet and
Chandra at X-rays have spatially resolved its WD companion, the accretion disk
and the stream fueling it (Karovska, 2006). Such a plethora of multi-
wavelength information is however not available for a typical GALAH anonymous
star, and to be on the safe side we decided - at the initial stage of our
project represented by this paper - to filter out the radial pulsators. Such a
pruning is based on converging photometric (lightcurve) and spectroscopic
(H$\beta$/H$\alpha$ emission ratio) criteria, described below.

Figure 5: Left-most panels: example H$\alpha$ and H$\beta$ profiles (in black)
from RV-zeroed GALAH spectra of a symbiotic star (top row) and of a radial
pulsator (bottom row), compared with those of the template for the same
$f$-ratio bin (in red, cf. Table 4). Center panels: the result of subtracting
the (red) template spectrum from the (black) object spectrum from the left
panels. Right-most panel: the corresponding $V$-filter lightcurve from ASAS-SN
sky patrol data.

We decided not to rely on existing compilations of known variables (as for
example collected in the VSX111http://vizier.u-strasbg.fr/viz-
bin/VizieR-3?-source=B/vsx/ catalog), because the adopted classification
criteria (as well as amplitude and periods) are sometimes in disagreement
among themselves and also in conflict with the naming conventions listed in
the IAU General Catalog of Variable Stars222http://www.sai.msu.su/gcvs/. We
preferred instead to analyze in a homogeneous way their photometric behavior
ourselves. To this aim we downloaded the available photometry for the
83-sample from the ASAS-SN sky patrol survey (Shappee et al., 2014; Kochanek
et al., 2017) and examine ourselves the resulting lightcurves, which are
usually composed of hundreds of individual observations in either $g$ or $V$
bands distributed over a time-interval of a few years, and therefore adequate
to reveal the presence of long and stable periodicities related to radial
pulsation. ASAS-SN sky patrol covers the whole accessible sky on both
hemispheres every night, and the $12<V_{JK}<14$ mag range of our targets is
ideally placed at the center of ASAS-SN dynamic range, away from the bright
saturation limit or the noisy detection threshold.

The lightcurve of a Mira shows regular, long period (from several months to
longer than a year) and large amplitude (from a few to $>$10 magnitudes)
sinusoid-like variations (e.g. Hoffmeister, Richter, & Wenzel, 1985; Sterken &
Jaschek, 2005). The large scale radial pulsation is betrayed by the in-phase
variation of the radial velocity of the absorption spectrum, with amplitudes
of the order of $\sim$10-15 km/s (Joy, 1926; Hoffmeister, Richter, & Wenzel,
1985). Our path-finder SU Lyn shows a significant variability (0.6 mag
amplitude in $V$), but it lacks the larger amplitude and the regular beat of a
Mira. As clearly illustrated by Fig. 2, the M giant in SU Lyn does not
radially pulsate in any significant way (the absorption lines are stable in
radial velocity to better than 2 km s-1), and therefore no shocks form in the
atmosphere which could lead to emission in the Balmer lines. Armed with these
considerations, and to stay on the cautionary side, we selected to flag (in
the 83-sample) all the stars showing ASAS-SN lightcurves with an amplitude in
excess of 0.7 mag and a stable and clean sinusoid-like shape (cf. Fig. 5).
Such a lightcurve could however arise also from orbital motion, for example
from ellipsoidal distortion when the cool giant fills its Roche lobe and/or
its side facing the companion is irradiated and heated up (e.g. as in the
accreting-only and recurrent nova T CrB; Munari, Dallaporta, & Cherini, 2016).
A spectroscopic confirmation of the radially pulsating nature is therefore
required, and this needs to be accommodated within the single-epoch spectra
(i.e. no revisit) of the GALAH survey.

In radially-pulsating cool giants, outward moving material collides against
the gas lifted during the previous cycle and now falling back inward. The
resulting shock is hot enough to excite emission in the Balmer lines of
hydrogen. Such shocks develop deep within the atmosphere, interior to the
outer layers where absorption by TiO molecules occurs. The absorption by TiO
begins around 4200 Å and, by superposition of successive bands (they all
decline toward the red), their combined absorption grows rapidly stronger with
wavelength (Fluks et al. 1994). As a consequence, emission in H$\delta$ (4101
Å) and higher Balmer terms exits the atmosphere unscathed by TiO, and their
flux declines with increasing upper $n$-quantum number in the usual way. What
happens to the red of H$\delta$? Going from H$\gamma$ to H$\beta$ and then
H$\alpha$, we move into deeper and deeper overlapping absorption by TiO bands
and the observed emission in the line becomes more and more absorbed. Saying
it a different way (Joy, 1926; Yamashita & Nariai, 1977), in pulsating M
giants the strongest line is H$\delta$ and the intensity of successive Balmer
emission lines declines either going to the blue or to the red (see exemplary
spectrum for LQ Sgr presented by Bragaglia et al., 1995). In most
astrophysical environments, the observed emission in H$\alpha$ is invariably
stronger than in H$\beta$, while in Miras the opposite is the case (Merrill,
1940), because of the TiO absorption in the outer atmospheric layers. This
offers a clear distinction between the emission originating from the accretion
disk in an acc-SySt (in terms of equivalent widths: EW(H$\alpha$) $>$
EW(H$\beta$)) from that expected to come from the internal shock regions of an
M-type radial pulsator (EW(H$\alpha$) $\leq$ EW(H$\beta$)).

GALAH high-resolution spectra show a further distinction between the two
types. In acc-SySt systems, a sharp absorption component appears superimposed
to the broad H$\alpha$ emission profile, generally blue-shifted by about
5$-$25 km s-1 and half as wide as the photospheric H$\alpha$ absorption line
(FWHM$\sim$15$-$25 km s-1 as opposed to 45$-$50 km s-1), that - by analogy
with the prototype SU Lyn or other acc-SySt discussed by Jorissen et al.
(2012) and Gorlova, Van Winckel, & Jorissen (2012) \- is believed to originate
in the gentle wind blowing off the accretion disk or in the out-flowing wind
of the RG that engulfs the whole binary system. In the GALAH spectra of
radially pulsating stars, the narrow absorption superimposed to the emission
is instead generally red-shifted as if coming from cooler material falling
back toward the star after being lifted higher up during the previous
pulsation cycle.

These photometric and spectroscopic criteria to separate acc-SySt and radially
pulsating objects in the 83-sample are illustrated in Fig. 5, which presents
line profiles and lightcurves well typical of the rest of the sample: high-
amplitude and regular beating for the lightcurve of a pulsator, with strong
H$\beta$ and red-shifted narrow-absorption; lower amplitude and less regular
lightcurve for the acc-SySt, an H$\beta$ much weaker than H$\alpha$ and blue-
shifted narrow-absorption. By applying these criteria, 30 radial pulsators are
pruned from the 83-sample, reducing the candidates acc-SySt to 53. Hereafter,
we refer to them as the 53-sample.

###  4.7 Avoiding contamination by T Tau stars

T Tau pre-main sequence variables have been considered as possible
contaminants in previous large-scale searches for new SySt, especially those
based on IPHAS $r$,$i$,H$\alpha$ photometric survey (Corradi et al., 2008,
2010, 2011; Rodríguez-Flores et al., 2014). T Tau stars are distributed over a
wide range of spectral types, from mid-F down to the coolest M-types, may be
heavily reddened, and present emission lines originating from the
circumstellar disk, the spotted surface and the magnetic-confined accretion
columns on the magnetic poles (Hartmann, Herczeg, & Calvet, 2016). T Tau stars
can therefore mimic the red $r$$-$$i$ colour of SySt and similarly stand out
in the $r$$-$H$\alpha$ photometric index. T Tau stars are less prone to be
confused with SySt if spectra are available, especially if they cover a broad
wavelength range (as those we have collected in Section 5.1 below for our
program stars): the spectra of T Tau stars show the absorption features
typical of a main sequence star, not of a giant, and the emission lines are
rather different both in assortment (CaII H & K doublet usually quite strong
in T Tau stars, while generally absent in SySt) and in their profile (in T Tau
stars sporting narrow jet components superimposed on a broad pedestal coming
from the rotating disk; Giannini et al., 2019). In addition, LiI 6707 Å is in
strong absorption in most T Tau stars, whereas it is generally absent in SySt.

In the context of our search for new SySt, T Tau stars are however of no
concern, because they cannot be confused with cool giants when the segregation
is carried out on Gaia parallaxes, as we did. Of relevance to the present
paper are only T Tau stars with an M spectral type (M-TTau for short). The
intrinsically brightest among M-TTau are those with the earliest M-types, M0V
and M1V. The absolute magnitude of M0V and M1V stars are M(V)=8.8 and 10.8,
respectively, while that of M0III and M1III giants is M(V)=$-$0.5 (Drilling &
Landolt, 2000; Sowell et al., 2007). The inflated photospheric radius (not yet
settled on the main sequence) and the contribution by the circumstellar disk
may raise the intrinsic brightness of M-TTau above that of single and isolated
dwarfs of the same spectral type, but they will never fill the huge
$\geq$10mag gap between main-sequence and giant stars of the M spectral type.

In confirmation of this, we have retrieved the Gaia $e$DR3 parallaxes for the
brightest known M-TTau (selected from the compilation by Yang et al., 2012)
and computed their M(V) magnitudes by adopting the photometric calibration of
the Gaia Consortium (Evans et al., 2018). The absolute M(V) magnitudes of
M-TTau turned out to be all fainter than +6.5, confirming that in no way an
M-TTau may have slipped into our sample of M giants segregated on the base of
Gaia parallaxes (see also Section 5.9 below). As a further check, we took the
compilation of T Tau stars listed in SIMBAD and VSX, and checked if any
matched the M-giants of the initial 15k-sample, and verified that none did.

###  4.8 Comparison with VVCep binaries

There are binaries which host cool supergiants and massive hot companions, of
which a comprehensive catalog of 108 Galactic objects has been recently
published by Pantaleoni González et al. (2020). The 23 entries from this
catalog with an M spectral type for the cool giant are summarized in Table 14,
where we add distances, absolute magnitudes and height above the galactic
plane as derived from Gaia $e$DR3 parallaxes and the 2MASS JHKS survey. Such
objects are usually named VVCep stars, from the best known member of the
group, and while not related to SySt are sometimes discussed in parallel. We
have not included $\delta$ Sge in Table 14 as it seems more related to other
types of objects like TX CVn or 17 Lep than to VVCep stars (Ginestet &
Carquillat, 2002; Pugh, Gray, & Griffin, 2015).

Figure 6: Distribution in M(KS) for the program new SySt compared to that of
VVCep stars. Figure 7: Distribution in height over the plane of the Galaxy for
the program new SySt and for VVCep stars.

The VVCep stars all have very massive components on very wide orbits. The
orbital solution for VV Cep itself gives an orbital period of 20.3 yrs, an
orbital separation of 24.8 AU, and individual masses of 18.2 M⊙ and 18.6 M⊙
for the M supergiant and the B-type companion, respectively (Bennett et al.,
2004). Of interest to us is the fact that in VVCep stars, the Balmer lines,
and H$\alpha$ in particular, may be seen in emission, with a line profile
reminiscent of that seen in Be stars (in which the emission comes from an
equatorial ring-like shell formed by material leaving the hot star which is
rotating close to break-up velocity). While the hot component in VVCep
dominates the spectrum in the blue, that of the M supergiant takes over at red
wavelength. With an H$\alpha$ in emission superimposed, it could be questioned
whether any of our new SySt may indeed belong to VVCep stars. This, however,
can be safely excluded on two independent grounds.

Figure 8: Left-most panels: comparison between the average $B$ and $V$
magnitudes from APASS all-sky survey, and the corresponding values derived
from Asiago 1.22m flux-calibrated spectra for a subset of the 53-sample. The
solid lines mark the 1:1 relation. Right-most panels: histogram distribution
of differences between Asiago and APASS derived $B$$-$$V$ colours, and
distribution of $U$$-$$B$ colours derived from Asiago spectra. The line in
blue shows the distribution of a hundred M giants from the Solar neighborhood
as measured by Fluks et al. (1994). Red colour refers to symbiotic stars
discovered in this paper, black to other GALAH M giants observed along with
them (including also some radial pulsators).

First, M supergiants of VVCep stars are intrinsically much brighter than the M
giants characterizing our new SySt. The two are compared in Fig. 6, with
values taken from Table 11 for the new SySt, and Table 14 for the VVCep stars.
The median absolute M(KS) magnitude for VVCep stars is $-$10.0, that of the
new SySt only $-$5.9. Secondly, their distance over the Galactic plane is
radically different. Being very massive and therefore very young, VVCep are
examples of extreme Pop I stars, laying close to the Galactic plane, while our
new SySt belong to a much older population as discussed in Section 5.9 below.
Their distributions in terms of height above the Galactic plane are compared
in Fig. 7, with input data taken from Table 11 and Table 14: the median value
is just 50pc for the VVCep stars and a much larger 1.8kpc for the new SySt.

##  5 The finally selected 33 new symbiotic stars and their properties

We have finally subjected the 53-sample to a series of follow-up observations
from the ground and the space to pick-up the most promising objects. Based on
the ensuing results, our final list of 33 proposed acc-SySt candidates is
presented in Table 5 (where GaSS in their names stands for Galah Symbiotic
Star); they are divided in two groups, a primary and a supplementary sample.
The division is based primarily on available follow-up observations (less for
the supplementary sample) and amplitude of photometric variability (larger for
the supplementary sample). A plot similar to Fig. 5, presenting the H$\alpha$,
H$\beta$ profiles and the lightcurve, is provided in the appendix for all 33
objects (Fig. 17 to 23).

Gaia $e$DR3 lists parallaxes ($\pi$) for all the 33 new SySt, and they are
reported with their uncertainties ($\sigma(\pi)$) in Table 11. In Table 5 we
provide distances, by direct inversion of parallax, for only 20 of them,
namely those satifying the condition $\sigma(\pi)$/$\pi$$<$0.2. This is the
limit reccommended by the Gaia team (Bailer-Jones, 2015; Luri et al., 2018).
The parallaxes of the remaining program stars are too small in comparison to
their formal uncertainties and even generally smaller than the mean 0.019mas
bias affecting Gaia $e$DR3 parallaxes as a whole (Lindegren et al., 2020).
Such bias depends in a non-trivial way on (at least) the magnitude, colour,
and ecliptic latitude of the sources, with no firm correcting recipe available
yet. Improved parallaxes in future Gaia data releases will allow to derive
valuable distances to (most of) the remaining program stars.

We now describe the follow-up observations that have led us to the compilation
of Table 5. A complete description of such extensive follow-up observations is
far beyond the scope and breadth of the present paper. Full details, augmented
by the results of further follow-up observations currently in progress, will
be discussed elsewhere.

###  5.1 Near UV up-turn

The presence of a bright accretion disk is betrayed by an emission excess at
wavelengths shorter than $\sim$4000 Å, as illustrated for SU Lyn in Fig. 1.
With the long-slit B&C spectrograph attached to the Asiago 1.22m telescope,
highly efficient down to the atmospheric cut-off ($\sim$3200 Å at the 1000m
above sea level location) we have observed most of the objects in the
53-sample located north of $-$29∘ in declination (Table 10 in the Appendix
provides a log-book of these observations). At the Asiago $+$46∘ latitude,
this corresponds to a horizon distance of 15 degrees. The near-UV faintness of
the program stars and the high atmospheric extinction at such large air-masses
conspired to make such observations quite demanding to execute. A consistent
number of spectro-photometric standards were observed in parallel each night
at similarly large air-masses. Suitably high S/N at 3400-3600 Å was obtained
for only a fraction of the attempted objects, and sometimes only after adding
spectra collected at different revisits. The objects marked "spc" in the
$U_{\rm exc}$ column of Table 5 are those showing excess emission on spectra
over the range 3400$\leq$$\lambda$$\leq$4000 Å. There are clear hints for some
of them that the amount of near-UV excess changes when spectra taken at
different epochs are compared.

Table 5: The candidate accreting-only symbiotic stars (acc-SySt) discovered in
this paper based on spectra from the GALAH survey. E(B-V) is from Schlafly &
Finkbeiner (2011). The spectral type is on the MKK scale and $f$ is the
template label/bin (see Table 4 and Fig. 16). $V$ is from the APASS catalog.
$\Delta V$ is our estimate of the amplitude of variability on ASAS-SN sky
patrol data (Section 4.6). ’P orb’ lists possible orbital periods (for
2$\times$P alternative see Section 5.7). ’dist’ is the distance (in kpc)
derived from Gaia $e$DR3 parallax for objects satisfying the conditions
$\sigma(\pi)$/$\pi$$<$0.2 (cf. Section 5). RV⊙ is from GALAH DR3 (rv_guess).
Objects marked ’var’ in the ’Asiago RV’ column show a variability in radial
velocity (cf. Table 6 and Section 5.2). The column ’Swift’ marks with ’X’ the
objects counterpart to X-ray sources in the 2SXPS catalog (Evans et al.,
2019), and with ’UV’ those found emitting in the ultraviolet UVM2 band in our
Swift follow-up observations (cf. Table 8 and Section 5.5). $U_{\rm exc}$
lists the presence of emission excess at $U$-band wavelengths: ’pht’ from
photometric UBV observations (Massey, 2002), ’spc’ from our spectra (Section
5.1). An ’F’ in the ’Flk’ column marks the presence of flickering (Section
5.4). The final two columns provide the velocity (w.r.t. M giant) and the FWHM
of the narrow absorption superimposed to the broader H$\alpha$ emission
(Section 4.6).

.

Table 6: Heliocentric radial velocity of candidate GALAH symbiotic stars
measured with the Asiago 1.82m telescope + Echelle spectrograph compared with
the corresponding values of rv_guess listed in GALAH DR3. The mean radial
velocity listed in Gaia DR2 is also given, with its formal error and the
number of epoch transits over which it has been computed. Figure 9: Example of
variability observed on H$\alpha$ comparing the original GALAH spectra (top
row) with Asiago 1.82m Echelle spectra obtained at later epochs (the left-
bottom spectrum is from Varese 84cm telescope). Only GALAH spectra are
corrected for telluric absorption lines.

A quick summary of the photometric data extracted from spectra is presented in
Fig. 8. The $B$ and $V$ data lie well on a 1:1 relation with APASS mean data,
the scatter being at least in part related to the variability all objects
display (as per column $\Delta V$ of Table 5). The difference in the $B$$-$$V$
colour (indicative of the accuracy of the slope of fluxed spectra) between
APASS and spectra appears well distributed around 0.0 with a FWHM of about 0.1
mag. The most interesting panel of Fig. 8 is the last one, showing the
distribution of $U$$-$$B$ measured on spectra compared with the distribution
of $U$$-$$B$ for a hundred M giants of the Solar neighborhood (suffering from
negligible reddening) as measured by Fluks et al. (1994). It is clear how, on
average, the $U$$-$$B$ colour of candidate acc-SySt is bluer that those of
(supposedly) single, normal M giants, supporting the notion that for a
sizeable fraction of them we have actually observed an excess in $U$ as coming
from the companion and the accretion disk. The acc-SySt scoring an "spc" in
Table 5 are those with an $U$$-$$B$ bluer than +0.4 in Fig. 8, and "spc:" for
those between +0.4 and +0.8.

###  5.2 Radial velocity variability

Armed with evidence from Fig. 2 that the low amplitude, irregular variability
affecting the M giant in SU Lyn does not reverberate into changes in radial
velocity, we observed some of the 53-sample north of $-$25∘ in declination
with the Echelle spectrograph mounted on the Asiago 1.82m telescope. The
results are summarized in Table 6. As a check, some normal stars from the
initial 15k-sample and a few radial pulsators from the 223-sample were also
observed, with results presented in Table 12. In both tables, the quoted error
for Asiago RVs is the internal error, i.e. that derived by comparing the RV
obtained from different Echelle orders by cross-correlation with spectra of
IAU radial velocity standards of M spectral type. We selected to limit the
measurement to 5 of the 32 orders covered by the Asiago Echelle spectrograph,
favoring those at redder wavelengths (where S/N is higher), and avoiding the
orders with TiO band-heads or telluric atmospheric absorptions that would
affect the cross-correlation results. A comparison of Asiago RVs with the
GALAH rv_guess and Gaia DR2 values in Table 12 for the GALAH M giants sub-
sample returns a negligible offset and an rms $\sim$1.5 km s-1. The latter
could in part be accounted by variability of epoch radial velocities intrinsic
to the sources.

From the data in Table 6, a clear RV variability has been observed for five
objects, and borderline for an additional three. They are marked with "var" in
column "Asiago RV" of Table 5. We favor an interpretation in terms of orbital
motion for this RV variability.

###  5.3 Emission line variability

The acc-SySt of this paper show a clear variability of the intensity and
profile of H$\alpha$ emission line on spectra collected with the Asiago 1.22m
and 1.82m telescopes. An example is presented in Fig. 9. GaSS 1-31 showed a
clear emission in H$\alpha$ on the GALAH spectrum recorded on Apr 9, 2017, a
weaker one on the Asiago spectrum for Oct 2, 2020, and no emission on Jul 5,
2020 (see also Section 5.5 below) and on Oct 17, 2020 (this last spectrum has
been obtained at a lower resolving power - 12,000 because of 2$\times$2 CCD
binning - with an Echelle spectrograph mounted on the Varese 84cm telescope).
The H$\alpha$ of GaSS 1-12 has been always observed in emission on all three
visits with the Asiago 1.82m, with a varying degree of intensity and velocity
of the emission relative to the M giant. Most interestingly, the last Asiago
spectrum for Oct 3, 2020 shows the appearance of a second narrow absorption,
blue shifted with respect to the primary one, an event reminiscent of what
happened to SU Lyn in 2017 (cf. Fig. 2).

###  5.4 Optical flickering

The accretion in binaries is characterized by chaotic processes (e.g. density
fluctuations in the accretion flow from the donor to the accreting star) that
develop on time scales (seconds, minutes or a few hours) much shorter than
those of other sources of variability (orbital motion, pulsation, rotation,
outbursts, etc.). The possibility to detect flickering on optical photometric
data depends primarily on its contrast with the steady sources in the system:
the fainter the two components of the binary, the easier for flickering to
become observable (e.g. Hellier, 2001). Therefore, if flickering is easy to
observe in cataclysmic variable stars that radiate about 1 L⊙ at optical
wavelengths, it is an entirely different matter for symbiotic binaries in
which the optical luminosity of the cool giant alone amounts to 102 L⊙.

The regions in the disk responsible for flickering are hot (Warner, 1995;
Zamanov et al., 2020), and their emission is consequently rising toward the
blue; that of the giant rises instead toward the red. Therefore, to reduce the
background glare of the giant it is convenient to go as blue as possible
(Skopal, 2005). Detection of flickering at satellite ultraviolet wavelengths
is straightforward as demonstrated by the Swift observations of SU Lyn by
Mk16. On ground-based observations, the bluest photometric band is $U$ (in its
many variants: Landolt’s $U$, $u$ from Stromgren, or SLOAN $u^{\prime}$). $U$
band was relatively easy to observe in the old times of photo-electric
photometers mounted on all-reflecting telescopes with bare-aluminum surfaces,
thanks to the high instrumental throughput (no coatings or refractive
components) and the excellent sensitivity of photo-multipliers like the
standard RCA 31034A and Hamamatsu R934-02. Many searches for flickering in
symbiotic stars were performed at that time (e.g. Slovak & Africano, 1978;
Mikolajewski et al., 1990; Tomov et al., 1996. The advent of CCDs, with their
much lower sensitivity at $U$-band wavelengths, badly affected the study of
flickering, with rarer attempts carried out (e.g. Dobrzycka, Kenyon, & Milone,
1996; Sokoloski, Bildsten, & Ho, 2001; Gromadzki et al., 2006; Zamanov et al.,
2011). To further complicate the matter, the sensitivity of CCDs, that extends
up and beyond 1 $\mu$m, allows signal to be recorded from the red-leak
affecting most $U$ filters, especially those of the modern multi-layer
dielectric type. Such a red-leak may be so severe in the case of cool giants,
that the fraction of recorded photons coming through the red-leak easily
outnumbers those going through the proper transmission profile of the $U$-band
(Munari & Moretti, 2012).

To search for flickering among our candidate acc-SySt we used a 50cm
telescope, with a 40arcmin well corrected field of view and quality
photometric filters, operated robotically for ANS Collaboration in Chile at
San Pedro de Atacama. For the above described limitations about $U$ filters,
we selected to search for flickering in $B$-band, with interspersed
observations in $V$-band serving to construct the $B$$-$$V$ colour base for
the transformation from the instantaneous local photometric system to the
Landolt’s standard one. Local photometric sequences were extracted from the
APASS multi-epoch, all-sky $B$$V$$g$$r$$i$ photometric survey (Henden et al.,
2016) in its DR8 version.

Given the much brighter emission of the M giant in $B$ compared to $U$, the
expected amplitude of observable flickering decreases to (at most) a few
hundredths of a magnitude. To be detected, such a tiny amount of variation
requires highly accurate observations and careful data reduction, under clear
and stable skies as usually enjoyable in Atacama. We adopted integration times
in $B$ of 30sec for the brightest targets and 5min for the faintest, with
60sec for the bulk of objects, thus ensuing that all objects have been
observed at high S/N ($>$100 on each single point) while remaining short of
the saturation thresholds and non-linear regime. Every 10 exposures in $B$,
one was obtained in $V$ to compute the coefficients of the transformation
colour equations. The procedure was repeated identically for 70min for each
program star. The photometry was performed for most objects in aperture mode,
their high galactic latitudes implying sparsely populated fields and therefore
no need to revert to PSF-fitting, which was instead necessary for three
targets. On average, the 50 field stars closest on the image to the symbiotic
star, of a similar magnitude and well isolated from neighboring stars were
also measured on all recorded images in exactly the same way as the symbiotic
star. The photometry of these 50 field stars was then inspected looking for
those with a $B$$-$$V$ colour as close as possible to the symbiotic star in
the center. Typically 5 such stars were found, with a few cases scoring just 2
and others up to 8. These field stars with magnitude and colour closely
matching those of the symbiotic star, serve as samplers of the observational
noise above which the flickering has to be detected: the closely similar
colour nulls the effects of differential atmospheric transmission, and the
closely similar brightness nulls the non-linear effects in the differential
statistical noise.

The program acc-SySt subjected to the search for flickering are listed in
Table 13, together with the HJD of the central image for the 70min time
series, and the average of $B$ and $V$ measured magnitudes during such time
series. The error of such $B$ and $V$ data is totally negligible in its
Poissonian component, while the amount due to transformation from the local
system to the APASS one (closely adherent to Landolt’s) is systematic and
constant for a given object and amounts on average to 0.008 mag (which
obviously cancels out when comparing flickering data within a given time-
series).

An example of the collected data is presented in Fig. 10, with one object each
for the 30sec, 1min and 5min sample times. At the top and bottom panels we
show objects with a clear detection of flickering well in excess of the noise
affecting the field stars as usually quantified by their respective rms
deviation from the mean (Dobrzycka, Kenyon, & Milone, 1996), while in the
center we present a non-flickering symbiotic that remains stable at a few
millimag-level through the 70min monitoring. GaSS 1-25 at the top of Fig. 10
exemplifies the frequent case of incessant flickering (e.g. RS Oph as observed
by Dobrzycka, Kenyon, & Milone, 1996), while the behavior displayed by GaSS
1-13 at the bottom of Fig. 10, i.e. isolated flarings over an otherwise
relative flat background, is a close match to that observed by Zamanov et al.
(2017) in EF Aql.

Figure 10: Examples of the $B$-band time series collected on the program
stars, with 30sec, 1min and 5min sampling time (from top to bottom), in search
for flickering from the accretion disk. The regular gaps correspond to the
acquisition of $V$-band frames for the calibration of colour equations to
transform data to the standard Landolt system. The measured $B$$-$$V$ colour
of the symbiotic star (in red and filled circles) and of nearby field stars of
similar brightness and colour (in black and crosses), is marked to the right,
together with the dispersion around the median of the plotted $B$-band data.
Symbiotic stars on the top and bottom panels clearly show flickering well in
excess of the noise affecting field stars, while that on the center panel is
constant at the same noise level of the field stars.

The objects for which flickering was detected well above observational noise
are marked with an ’F’ in the ’Flk’ column of Table 5. The detectability of
flickering and its amplitude are usually dependent on the actual epoch of
observation (Stoyanov et al., 2018). Orbital phases with the hot spot in plain
view are favoured (WD at quadrature moving toward superior conjunction), while
others can be detrimental, e.g. during eclipses or passage at apoastron in a
highly eccentric orbit (with consequent reduction in mass transfer rate).
Therefore, stars in Table 5 that failed to show flickering at the single epoch
of our observation could do otherwise at a later revisit. This is also the
case for SU Lyn. When it was observed in late 2015 by Mk16 its was in a high
accretion state and it flickered wildly, while it was in a low accretion state
and did not flicker when Dimitrov et al. (2018) performed their observations
in January 2018. Similarly, while large amplitude flickering has always been
the rule for MWC 560 (eg. Lucy et al., 2020), it has disappeared altogether
during the last couple of years (Zamanov et al., 2019).

We postpone a detailed analysis of the data collected during our search for
flickering for the program stars to a follow-up paper. Derivation of
temperature and radius of the region(s) responsible for the observed
flickering (e.g. Zamanov et al., 2020) is straightforward but largely out of
the goals of this paper, as it is the search about possible quasi-
periodicities or signatures of the WD rotation (e.g. Sokoloski & Bildsten,
1999; Zhekov & Tomov, 2019). Such analysis will strongly benefit from further
observations we have planned for the SySt reported in this paper and are
carrying out also for other SySt assigned to later papers in this series.

###  5.5 Satellite X–ray/UV observations

The presence and nature of the accreting source in acc-SySt systems, either a
WD or a NS, is best defined by satellite observations performed in the
ultraviolet and X–rays, where the peak of the emissivity for accretion-induced
processes is located (e.g. Masetti et al., 2007a and references therein, and
also Kuranov & Postnov, 2015 for a recent review). The X–ray/UV properties of
the acc-SySt prototype SU Lyn have been investigated by Mk16, Lopes de
Oliveira et al. (2018), and Kumar et al. (2021) based on observations with
Swift, NuSTAR and ASTROSAT satellites. They found SU Lyn to show (i) a large
(by a factor of a few hundred in flux) UV excess with respect to a ‘normal’
red giant; (ii) long-term (months to years) hard X–ray variations by an order
of magnitude between low- and high-level states in the 15–35 keV hard X–ray
emission; and (iii) fast UV variability during the high X–ray states
(flickering). In the latter ones, the average X–ray flux was $\sim$10-11 erg
cm-2 s-1 in the 0.3–50 keV band, with the spectrum modelled with a thermal
plasma of temperature $kT\sim 20$ keV, plus a fluorescent iron emission at 6.4
keV, absorbed by a hydrogen column NH as large as $\sim$3$\times$1022 erg
cm-2. This absorption has however little effect on the UV excess observed from
this source. At a distance of 650 pc, this corresponds to unabsorbed X-ray
luminosities of up to $\sim$1033 erg s-1 in the high state, and a factor of 10
larger in the UV. The passage from low- to high-level states in the UV and
X-ray emission is attributed to large excursions in the transfer and accretion
rates.

Table 7: Results of the Swift XRT and UVOT observations of some of the acc-SySt presented in this paper. The fluxes are in units of 10-13 erg cm-2 s-1. The bottom three objects are radial pulsators observed for comparison. The data for 20 35 04.52 $-$05 47 35.3 include also additional 485sec exposure obtained on 8 Sep. 2019. name / RA,DEC | Obs. | Start | XRT obs. | 0.3–10 keV | 0.3–10 keV | UVOT obs. | UVM2
---|---|---|---|---|---|---|---
| date | time (UT) | duration (s) | count rate | flux | duration (s) | mag.
GaSS 1-11 | 13 Jun. 2020 | 12:05 | 1885 |  $<$3.0e-3 |  $<$2.6e-13 | 1864 | 17.33$\pm$0.06
GaSS 1-12 | 16 Sep. 2019 | 00:02 | 1516 |  $<$3.2e-3 |  $<$2.8e-13 | 1701 | 16.21$\pm$0.04
GaSS 1-31 | 14 Jun. 2020 | 01:01 | 839 |  $<$3.7e-3 |  $<$3.3e-13 | 830 |  $>$19.92
05 11 30.88 $-$61 29 03.6 | 10 Jun. 2020 | 06:31 | 1973 |  $<$1.7e-3 |  $<$1.5e-13 | 1948 |  $>$20.65
20 06 56.32 $-$28 35 32.3 | 11 Oct. 2019 | 00:49 | 1536 |  $<$2.5e-3 |  $<$2.2e-13 | 1527 |  $>$20.71
20 35 04.52 $-$05 47 35.3 | 04 Sep. 2019 | 12:23 | 1504 |  $<$3.6e-3 |  $<$3.2e-13 | 1485 |  $>$20.82

In anticipation of a devoted observing campaign, we have performed some quick
exploratory observations of three of the acc-SySt discovered in this paper
with the Swift satellite (Gehrels et al., 2004), and for comparison of three
of the identified radial pulsators considered as a control sample. The
observations have been carried out in Target-of-Opportunity mode. This type of
observations is generally limited to roughly 2000 seconds per object, so this
can be considered as a quick and rather shallow survey of our target sample.
The pointings were performed in two time slots: the first one in September-
October 2019, and the second one in June 2020. The results are summarized in
Table 7.

Our Swift observations were acquired 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 covering the
X–ray band between 0.3 and 10 keV band, whereas UVOT data were collected using
the UV filter $UVM2$ (reference wavelength: 2246 Å; see Poole et al., 2008 and
Breeveld et al., 2011 for details). On-source pointings were simultaneously
performed with the two instruments and lasted between $\sim$800 and $\sim$2000
s.

All observations were reduced within the ftools environment (Blackburn, 1995).
The 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) within the same XRT acquisition. In all
cases, no emission above a signal-to-noise threshold $S/N=3$ was detected. The
XRT count rate upper limits in the 0.3–10 keV range were then measured within
the xspec package.

X–ray flux limits were determined using the webpimms online
tool333https://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3pimms/w3pimms.pl by
assuming a spectral model similar to that of Mk16 for SU Lyn, i.e. a thermal
plasma emission with temperature $kT$ = 17 keV absorbed by a column density NH
= 2.9$\times$1022 cm-2, which implies a count rate-to-flux conversion factor
of $\sim$8.8$\times$10-11 erg cm-2 s-1 cts-1. We note that, with this model,
the unabsorbed fluxes (in Table 8) in this same band are about 50% larger than
the absorbed ones.

Count rates on Level 2 (i.e. calibrated and containing astrometric
information) UVOT images at the position of the objects of our sample were
measured through aperture photometry using 5′′ apertures, whereas the
corresponding background was evaluated for each image using a combination of
several circular regions in source-free nearby areas. Magnitudes and upper
limits 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.

Of the three acc-SySt observed with Swift, two were found to be strong UV
emitters, confirming their symbiotic nature. The third, GaSS 1-31 was not
detected. There may be a clear reason for that, though. The Swift pointing was
carried out on 14 June 2020. Three weeks later, on 5 July 2020, we observed
the same star with the Asiago 1.82m and the Echelle spectrograph. The
H$\alpha$ profile for that date is presented in Fig. 9, and doesn’t show the
faintest trace of an emission. In addition, the equivalent width of the
H$\alpha$ absorption is same as for field stars of the same spectral type (and
greater than for the other dates in Fig. 9), indicating no partial filling
from even very faint emission. In short, it looks like the emission from the
accretion disk was ’switched-off’ at the time of the Swift pointing, with
consequent non-detection. An eclipse behind the M-giant seems improbable in
view of the limited variability of the epoch GALAH and Asiago radial
velocities listed in Table 6 (and their similarity with the average velocity
provided by Gaia DR2), precluding a high orbital inclination for GaSS 1-31. It
rather seems a drastic reduction in the mass accretion rate was taking place
at the time of Swift observations. This could have been caused by a reduction
in the wind blown-off by the M giant or passage at apoastron in a highly
eccentric orbit. The spectrum taken a hundred days later (that for 3 Oct 2020
in Fig. 9) shows that accretion has weakly resumed, with both a reduction in
the equivalent width of H$\alpha$ absorption (partially filled-in) as well as
emission wings extending above the local continuum.

None of the three radial pulsators we tried with Swift were detected, in spite
of having been selected among those showing the strongest H$\alpha$ and
H$\beta$ emission on GALAH spectra (far stronger than typically observed in
acc-SySt). This supports the clear-cut role of satellite UV observations in
segregating acc-SySt (Sahai et al., 2015), in particoular from radial
pulsators of similar spectral appearance on optical spectra.

A few final words are in order to explain the non-detection of GaSS 1-11 and
1-12 in X-rays in our shallow observations. The upper limit to their flux is
about 0.01$\times$ the flux Mk16 recorded from SU Lyn during their Swift
observations. Those observations of SU Lyn were obtained during an
exceptionally bright state of the accretion disk (cf. Fig. 2 and the
discussion at the end of Section 3), while the flux at other times (when the
H$\alpha$ shown by SU Lyn is much weaker and more similar in appearance to
that revealed by GALAH spectra for our sample of 33 acc-SySt) was just 1/10 of
that (cf. Lopes de Oliveira et al., 2018). There are two other factors playing
against GaSS 1-11 and 1-12. The most important is the distance. SU Lyn is just
0.65 kpc away, while GaSS 1-11 and 1-12 are respectively at 6.6 and 4.6 kpc
distance according to Gaia $e$DR3 parallaxes, corresponding to dilution
factors of 100 and 50. In addition, SU Lyn is seen face-on, while the unknown
orbital inclination of GaSS 1-11 and 1-12 introduces a reduction in flux
proportional to $\cos i$ and a further reduction for limb-darkening reasons.
Considering all factors together, it is no wonder that GaSS 1-11 and 1-12 fell
below the X–ray detection threshold of snapshot ToO with Swift.

###  5.6 Accretion luminosities

Table 8: Unabsorbed accretion luminosity from the Swift data in Table 7 (see Section 5.6 for details). symbiotic | LUV | LUV | Macc
---|---|---|---
star | (erg s-1) | (L⊙) | (M⊙ yr-1)
GaSS 1-11 | 1.43e34 | 3.7 | 1.7e-9
GaSS 1-12 | 3.30e34 | 8.5 | 4.0e-9
GaSS 1-31 |  $<$4.29e32 |  $<$0.11 |  $<$5.1e-11
Table 9: Flux recorded in H$\alpha$ and corrected for reddening (F(H$\alpha$)), isotropic luminosity radiated in H$\alpha$ at Gaia $e$DR3 parallax (L(H$\alpha$)), and accretion luminosity (Lacc) derived as described in Section 5.6 for the program star with available Asiago spectra to flux-calibrate GALAH spectra. symbiotic | F(H$\alpha$) | L(H$\alpha$) | Lacc
---|---|---|---
star | (erg cm-2 s-1) | (erg s-1) | (L⊙)
GaSS 1-12 | 1.391E-13 | 3.52e+32 | 14
GaSS 1-14 | 3.307E-14 | 2.93e+32 | 12
GaSS 1-15 | 3.205E-14 | 3.23e+31 | 1
GaSS 1-23 | 7.257E-14 | 2.00e+32 | 8
GaSS 1-25 | 1.179E-13 | 2.61e+32 | 10
GaSS 1-28 | 2.785E-14 | 1.88e+32 | 7
GaSS 1-31 | 2.367E-13 | 5.24e+32 | 22

We have previously stated that the aim of the present paper is to search for
new SySt accreting at low rates similar to SU Lyn. Mass accretion rate
($M_{\rm acc}$) and accretion luminosity ($L_{\rm acc}$) are related by the
usual

$L_{\rm acc}=G\frac{M_{\ast}M_{\rm acc}}{R_{\ast}}$ (4)

where $M_{\ast}$ and $R_{\ast}$ are the mass and radius of the accreting
object, respectively. The accretion on a WD, as it is the case for SU Lyn,
radiates most of its output in the UV and much less into X-rays, while the
reverse applies to neutron stars and their much deeper potential well.

The UVM2 magnitudes of the new SySt in Table 7 correspond to flux densities
(in units of 10-15 erg cm-2 s-1 Å-1) of 0.54, 1.52 and $<$0.05, for GaSS 1-11,
1-12 and 1-31, respectively (from the countrate-to-flux calibration given in
Poole et al., 2008 and Breeveld et al., 2011). Then, considering the distances
and $E(B-V)$ colour excesses reported in Table 5, correcting for the Galactic
reddening according to the prescription in section 3.2 of Kataoka et al.
(2008) and the UV correction coefficients in their table 5, and assuming for
the three objects a flat spectral distribution in the UV as Mk16 did for SU
Lyn, we obtain the unabsorbed UV luminosities listed in Table 8. In the
hypothesis of accretion onto a WD of mass 1 M⊙, as Mk16 has inferred for SU
Lyn, in Table 8 we report also the corresponding accretion rates.

For the other new SySt that were not observed by Swift, an estimate of $L_{\rm
acc}$ can be derived from the flux recorded in the H$\alpha$ emission. If an
emission line forms in an accretion disk, the flux radiated by the line and
the flux emitted by the disk as a whole are related. The correlation can be
tight, with temporal variations in $L_{\rm acc}$ that may be accurately
tracked and replicated by $L$(H$\alpha$) (cf. Munari et al., 2019, their fig.
14).

GALAH spectra are unfortunately not flux-calibrated. To circumvent the
problem, we look for assistance from the fluxed Asiago spectra. After de-
reddening the Asiago spectra for the $E_{B-V}$ values listed in Table 5, we
used them to flux-calibrate the GALAH spectra of the corresponding objects.
Being primarily interested in deriving just an order-of-magnitude, and noting
the new SySt exhibit only a limited photometric variability (cf. Section 4.6
above), we ignore the fact that Asiago and Galah spectra are not simultaneous.
On the template-subtracted version of the flux-calibrated GALAH spectra, we
measured the integrated flux of H$\alpha$ emission line $F$(H$\alpha$)
(therefore already corrected for reddening), and list it in Table 9, where we
consider only the new SySt with Gaia $e$DR3 distances in Table 5. If $d$ is
the distance, the luminosity $L$(H$\alpha$) radiated isotropically by the disk
in H$\alpha$ is obviously

$L{\rm(H\alpha)}=4\pi d^{2}F{\rm(H\alpha)}$ (5)

As customary practice, we may write a power-law proportionality between the
luminosity radiated in the line and by the disk as a whole:

$\log(L_{\rm acc})=\alpha_{\rm H\alpha}+\beta_{\rm
H\alpha}\times\log(L{\rm(H\alpha)})$ (6)

where $\alpha_{\rm H\alpha}$ and $\beta_{\rm H\alpha}$ are coefficients
specific to H$\alpha$. Adopting their value from Mendigutía et al. (2011), we
have derived the $L_{\rm acc}$ listed in the last column of Table 9. The
accretion luminosities derived directly from Swift ultraviolet data or
estimated from the flux radiated in H$\alpha$ are similar, and confined to the
1$-$20 L⊙ range, the same roughly spanned by SU Lyn while varying between
high- and low-states. Interestingly, there is an estimate of $L_{\rm acc}$ for
GaSS 1-12 both from the Swift UV observations and also from H$\alpha$ emission
line flux, and the two differ by less than a factor of 2. While this can be
read as an indirect confirmation of the reliability of $L_{\rm acc}$ estimated
from H$\alpha$, it also points to the constant variability of $L_{\rm acc}$
observed in acc-SySt. In fact, the $F$(H$\alpha$) of GaSS 1-12, measured on
the fluxed version of the Echelle spectra presented in Fig. 9, varies by more
than the above 2$\times$ factor.

###  5.7 Orbital-like periodicities

Figure 11: Examples of the modulation of the $B$$V$$R$$I$ lightcurves induced
by orbital motion in non-eclipsing symbiotic stars (see Section 5.7). Note the
increasing impact of flickering going toward bluer bands for the accreting-
only T CrB, and its absence for the burning-type LT Del. The lightcurves of
both objects are obtained from data extending over at least 10 consecutive
orbital cycles. Adapted from Munari et al. (2019).

The long-term lightcurves of SySt are usually reconstructed from multi-band
photometric campaigns (e.g. Skopal et al., 2007, 2012; Sekeráš et al., 2019),
from digging historical plate archives (e.g. Munari & Jurdana-Šepić, 2002), or
searching the public databases of all-sky patrol surveys (e.g. Gromadzki et
al., 2009; Gromadzki, Mikołajewska, & Soszyński, 2013; Munari et al., 2021).
Aside from the obvious effect of eclipses for systems seen at high
inclination, the orbital motion may affect also the lightcurve of SySt seen at
intermediate inclinations. The basic type of orbital-induced modulation of the
lightcurve of non-eclipsing SySt is shown in Fig. 11 (adapted from Munari,
2019), aiming to highlight the interplay between the irradiation effect
(better visible in the bluest photometric bands and for burn-SySt) and the
ellipsoidal distortion (which is more prominent in the reddest bands and in
acc-SySt).

We have already mentioned that the M3III giant of T CrB fills its Roche lobe,
causing an ellipsoidal deformation of the surface, that along an orbital
period presents to observers twice the larger projected area (the two maxima
in the lightcurve) and twice the smaller one (the two minima). The reason for
the change in brightness is primarily geometric, so its shape is similar at
all wavelengths as the left panel of Fig. 11 shows. In that figure, the
presence of flickering superimposed to the regular orbital motion is quite
obvious, in the form of a scatter of points of an increasing amplitude moving
from redder to bluer bands (cf. Section 5.4 above). A search for periodicity
on such a lightcurve would return half the value of the true period: 114 days
instead of 228. It is the spectroscopic orbit that would solve the uncertainty
in favor of 228 days.

Figure 12: Phased light-curves (see ephemeris at the top of each panel) for
the twelve acc-SySt marked in Table 5 as showing a modulation reminiscent of
orbital motion.

The cool giant in LT Del, the other symbiotic star in Fig. 11, equally fills
its Roche lobe, but contrary to T CrB, its WD is burning on its surface. The
hot and bright companion irradiates the facing side of the cool giant, thus
heating its photosphere and partially ionizing the atmosphere above it, with
the consequence to make the system brighter when the irradiated side is in
full view (Kenyon, 1986). The effect is larger going to bluer wavelengths.
Thus, while the behavior of T CrB and LT Del is the same in $I$-band where the
ellipsoidal distortion of the cool giant dominates, differences begin to
emerge moving toward shorter wavelengths: the irradiation effect takes the
lead in shaping the $B$-band lightcurve of LT Del. If the support of radial
velocities is required for T CrB, the orbital period of LT Del can be fixed
without uncertainty on the base of multi-band photometry alone. Unfortunately,
most all-sky patrol surveys observe in just one band, and a disentangling like
that for LT Del is not possible with their data alone, leaving on the desk the
ambiguity between P and 2$\times$P for the true orbital period.

We have searched the ASAS-SN photometric data for the 33 candidate acc-SySt
looking for a type of lightcurve modulation that could be orbitally induced
along the lines just outlined for T CrB and LT Del in Fig. 11. We found a
suitable shape for 12 of them, and they are presented in Fig. 12, with the
corresponding periods listed in Table 5. Their length is rather typical for
symbiotic stars, the majority of known cases going from 7 months to 2.5 years
(Mikołajewska, 2003).

In two cases, GaSS 1-27 and GaSS 1-28, a slight difference in the depth of
minima seems enough to resolve the ambiguity between P and 2$\times$P. Given
the limited time span of ASAS-SN coverage, such a difference could also result
from other causes for photometric variability, and needs confirmation over
more orbital cycles. Even more so for a possible third case for unequal
minima, that of GaSS 1-33, and its sparsely covered lightcurve (cf Fig. 12). A
borderline case is that of GaSS 1-8, with just a hint of a secondary minimum
around phase 0.5 which could be spurious. The modulation displayed by GaSS
1-28, while somewhat larger than usual, is a fine match in both shape and
amplitude to that of Hen 3-1591, whose modulation Gromadzki et al. (2007)
attribute to orbital modulation and report that it is confirmed by variations
in radial velocity.

The remaining objects in Fig. 12 show equal-depth minima, so the actual
orbital period could be twice the value listed in Table 5 and on the
ephemerides in their respective panels of Fig. 12. To help resolve the
controversy, we have started a $B$$V$$R$$I$ monitoring of them all, but the
fruits of this effort will be collected only in a few years time, when at
least a few orbital cycles will be covered.

###  5.8 Lithium

Wallerstein et al. (2008), in their chemical analysis of high resolution
spectra of M giants in symbiotic stars, found that the recurrent novae and
accreting-only systems T CrB and RS Oph show an over-abundance of 7Li in their
spectra. None of the other symbiotic stars in their large sample showed 7Li
enhancement and none was known to have undergone nova outbursts similar to T
CrB and RS Oph. The over-abundance of 7Li in T CrB has been recently revised
upward to A(Li) = 2.4$\pm$0.1 by Woodward et al. (2020). The apparent link
between 7Li and nova outbursts has been further reinforced when the symbiotic
star V407 Cyg erupted in 2010. 7Li over-abundance was discovered earlier in
V407 Cyg by Tatarnikova et al. (2003) at a time when the object was simply
known as a Mira variable with a possible companion, and therefore attributed
to indigenous production via the Cameron & Fowler (1971) mechanism activated
in the Mira variable. We know now that also V407 Cyg undergoes nova outbursts,
and this sheds an alternative explanation for its 7Li overabundance.

The production of 7Li during nova eruptions has been predicted theoretically
(e.g. José & Hernanz, 1998; Rukeya et al., 2017; Starrfield et al., 2020) and
confirmed observationally (e.g. Tajitsu et al., 2015; Izzo et al., 2015;
Molaro et al., 2016). Up to $\sim$1$\times$10-5 of the whole mass of the
ejecta can be in the form of 7Li, actually 7Be that will decay into 7Li. The
large amounts produced during outbursts (especially those occurring on white
dwarfs of the Carbon-Oxygen type), suggests that novae could be the main
producer of 7Li in the Milky Way (Molaro et al., 2020).

Figure 13: The GALAH spectrum of GaSS 1-26 (black) around the position of LiI
6707.8 Å compared with that (red) of the template for the same $f$-ratio bin.
The overabundance of LiI is quite obvious.

It is tempting to attribute the over-abundance of 7Li observed in T CrB, RS
Oph and V407 Cyg to pollution of the cool giant by the ejecta of the nova
outbursts experienced by the WD companion. Their eruptions occur at a high
frequency (eight have been observed in historic times for RS Oph) such that
they could replenish the surface of the cool giant with 7Li faster than
convection can dilute 7Li into the interior of the star. Alternatively, 7Li
could instead come from the interior of the cool giant, and conditions leading
to an efficient transport of 7Li to the surface could also lead more probably
to nova outbursts as a consequence of a higher mass loss rate and/or an
enrichment in CNO nuclei.

In analogy with T CrB, RS Oph and V407 Cyg, we have searched for 7Li
overabundance in our 33 candidate new SySt by comparing with the template
spectrum for their $f$-ratio, and found a clear case of enhancement in GaSS
1-26, as shown in Fig. 13. The object is completely anonymous (no matching
entry in SIMBAD, VSX and similar catalogs) and its optical variability is
inconspicuous (cf Fig. 22 in the Appendix).

We have subjected GaSS 1-26 to a search for flickering and got a robust
detection (cf Table 5), which supports its binary nature and the presence of
an accreting companion in the form of a WD.

Armed with that, we turned to the Harvard plate stack in search for an
unnoticed nova outburst in the distant past. DASCH project provides access to
calibrated scans of the huge collection of sky patrol photographic plates
obtained over a century with Harvard astrographs located in both hemispheres
(Grindlay et al., 2009). The area of the sky where GaSS 1-26 is located has
not yet been included in the available DASCH data releases (1 to 6). We have
been however granted access to the data for the area of sky around GaSS 1-26
prior to their publication in DASCH DR7, currently planned for July 2021\.
This led to some initial excitement when two plates for June 6, 1908 and
September 17, 1916 were logged as reporting GaSS 1-26 much brighter than at
other epochs and at a blue magnitude compatible with a nova outburst. This
pair of observations belongs to a time interval with very few Harvard plates
covering the GaSS 1-26 position and being deep enough to have recorded it in
quiescence (we measured it at $B$=15.4 during our search for flickering, cf.
Table 13, thus quite faint for the photographic means of the time). To assess
if any nova outburst had really affected GaSS 1-26 in 1908 and/or 1916, we
asked for and were provided with a scan of both plates (Edward Los, private
communication). Inspection of these scans reveals the presence of artefacts at
the astrometric position for GaSS 1-26, in the form of ink from a pen marking
for the 1908 plate and a scratch on the emulsion for the 1916 one. Therefore,
no nova outburst affected GaSS 1-26 in either 1908 and 1916. Even if the
search on Harvard plates for unnoticed past nova outbursts has not been
fruitful, it should be extended to other photographic plate archives, although
not many cover the southern hemisphere at faint magnitudes and over protracted
periods of times, and even less have preserved the plates in good conditions
and provide access to them in person or via digitized scans.

Figure 14: Location of the new SySt on the near-IR isochrone plane, compared
with Padova PARSEC isochrones (top) and observed ones (bottom) for clusters
spanning a range in metallicity (see Section 5.9 for details). Figure 15:
Distribution in absolute 2MASS $K_{S}$ magnitude of Belczyński et al. (2000)
symbiotic stars with an M spectral type, and those from the present paper.
Only objects with a Gaia $e$DR3 parallax satisfying the accuracy condition
$\sigma(\pi)$/$\pi$$<$0.2 are plotted.

In favor of a possible pollution scenario for the 7Li in GaSS 1-26 is the fact
that indigenous 7Li production via the Cameron & Fowler (1971) mechanism is
expected to occur during hot-bottom burning conditions (HBB), when the outer
part of the burning H-shell is included in the envelope convection. A key
aspect of HBB is the existence of a lower mass limit for it to occur, which
depends on metallicity (Herwig, 2005). The minimum initial mass for HBB is
Mini$\geq$5 M⊙ at solar metallicity (Forestini & Charbonnel, 1997), which can
decrease down to 3 M⊙ for very metal-poor stars (Siess, Livio, & Lattanzio,
2002). The latter is obviously not the case given the otherwise normal
spectrum displayed by GaSS 1-26. The faintness and negligible reddening place
GaSS 1-26 at a great distance from us; its Gaia $e$DR3 parallax is affected by
a large error, but it seems safe to argue that a lower limit to the distance
is 10 kpc. At a galactic latitude $-$10${}^{\circ}_{.}$6, it means a height
over the galactic plane of $z$$\geq$1.9 kpc. It would be rather atypical to
find a young and massive Mini$\geq$5 M⊙ star at such a large height above the
Galactic plane.

###  5.9 Evolutionary status

The nature of the new SySt will be investigated in detail elsewhere, including
their chemical partition and Galactic orbits. However, some preliminary
comments seem in order here.

To asses the evolutionary status of the new SySt, we have taken their
$J$$H$$K_{S}$ near-IR photometry from 2MASS survey (Cutri et al., 2003) and
after correcting for reddening and distance given in Table 5, we have plotted
their positions on the absolute $M(K_{S})$/de-reddened $(J-K_{S})_{\rm o}$
plane in Fig. 14.

On the top panel of Fig. 14, the comparison is carried out against the Padova
PARSEC isochrones (Bressan et al., 2012; Marigo et al., 2017; Pastorelli et
al., 2020). It provides evidence that the new SySt are all located close to
the tip of the Giant Branch, and that they are relatively old and metal poor
objects. This conforms with their significant height $z$ above the Galactic
plane as given in Table 11. For that age and metallicity, their mass estimated
from the isochrones ranges from 1.03 to 1.22 M⊙. Similar conclusions are drawn
by comparison with the observed isochrones compiled by Karaali, Bilir, & Gökçe
(2013) for a sample of clusters spanning a large range of observed
metallicities, as illustrated in the bottom panel of Fig. 14.

How do the red giants of the new SySt compare to well established SySt ? This
is addressed in Fig. 15, where their distributions in $M(K_{S})$ are compared.
For the well established SySt we adopted the catalog by Belczyński et al.
(2000), and retained only the 44 objects with a giant of M spectral type and
with a Gaia $e$DR3 parallax accurate enough to satisfy the inversion condition
$\sigma(\pi)$/$\pi$$<$0.2. To compute their $M(K_{S})$ we adopted for them the
interstellar extinction estimated from the Bayestar 2019 3D model of the
Galaxy (Green et al., 2019).

The distribution in Fig. 15 of the new and the well established SySt is the
same. The fact that GALAH is a magnitude-limited survey ($12<V_{JK}<14$) means
that intrinsically brighter objects need to be at larger distances to fall
between the limits in magnitude. The lack of GALAH new SySt at the brighter
$M(K_{S})$ values in Fig. 15 simply reflects the fact that current Gaia $e$DR3
parallaxes are not accurate enough to satisfy the condition
$\sigma(\pi)$/$\pi$$<$0.2 for the most distant, and therefore intrinsically
brightest objects. It may be guessed that future Gaia data releases will push
forward the horizon of parallax accuracy and allow to include the remaining
new SySt and let them populate the histogram of Fig. 15.

##  6 Conclusions and future prospects

We present the discovery and characterization of 33 bona-fide candidates of
the accreting-only type of symbiotic stars. So far, this class of binary
systems has evaded a thorough scientific scrutiny, otherwise received by its
counterpart, the nuclear-burning class of symbiotic stars. Both describe the
same underlying configuration, a binary system composed of a red giant and a
degenerate companion, however, they are distinguished by the physical
processes under way inside the system and the consequent observational
implications, favoring the study of burn-SySt. Because acc-SySt are believed
to be much more abundant in the Galaxy than burn-SySt, while the ratio of
known (investigated) objects is just the reverse, it is important to uncover
this hidden population of interacting binary stars in order to assess the
viability of symbiotic stars as progenitors of type Ia supernovae, their
impact on the pollution of the ISM, and to better understand the cycle of
exchange between the accreting-only and the nuclear-burning phases.

To this end, we have investigated a large sample of 600 255 stars from the
GALAH spectroscopic all-sky survey, which, given its unbiased observational
strategy, promises to enable the envisaged statistical analysis, while
supported by a suite of follow-up observations, provides a rich source of
information on individual symbiotic stars. The first 33 candidates presented
in this paper were selected using relatively conservative criteria: parallax
and colour cut, strength of molecular (TiO) absorption bands, prominence of
H$\alpha$ emission line, rejection based on the presence of various
observational and reduction artefacts. Furthermore, the photometric light
curves were checked to filter out radial pulsators, and great care was taken
to reinforce our diagnostics for the symbiotic nature of program stars,
building on the knowledge provided by the prototype acc-SySt SU Lyn and by
conducting follow-up ground and space based observations examining the near-UV
excess, RV and emission line variability, optical flickering, and X–ray/UV
luminosities.

One of the 33 objects, Gass 1-26, shows a clear enhancement in 7Li, which
makes it an excellent candidate for follow-up observations aiming to
illuminate the possible connection between pollution by novae outbursts and
the galactic lithium abundance. Future studies will extend the analysis
presented in this work in diverse directions, including:

  * •

weaker emission features in spectra will be considered, as we currently only
accept H$\alpha$ profiles which, after subtraction of the template spectrum,
reach higher than 0.5 above the adjacent continuum

  * •

other spectral diagnostic features will be established to identify acc-SySt
despite a possible absence of emission lines, for example a peculiar chemical
signature in the spectrum of a giant star that was polluted by eruptions from
the compact companion and thus by nuclearly processed material

  * •

earlier spectral types of the primary stars will be investigated, also to
account for the more metal-poor systems which do not achieve the same large
radius and low surface temperatures as M giants

  * •

automatic and manual inspection of legacy observations is being carried out,
such as digging through diverse photometric archives and historical plates
stacks in order to find e.g. past active phases or nova outbursts

  * •

existing and future high-energy instruments/surveys (e.g. SWIFT, $e$-ROSITA)
can be exploited for follow-up observations or cross-matching of catalogued
detections

  * •

combining the chemical abundances and galactic orbits with the aid of the next
Gaia releases will provide the means to confirm/reject the belief of symbiotic
stars being an old population, primarily belonging to the bulge, thick disk,
and halo. This view is built from the infrared characteristics of their cool
giants, but it is heavily biased by how (and where in the Galaxy) the
classical symbiotic stars were discovered by objective prism surveys of half a
century ago. Determining the parent stellar population of symbiotic stars has
a profound implication considering that Ia are the only type of supernovae
known to erupt in elliptical galaxies, whose stellar populations resemble the
bulge of our Galaxy.

The proposed future work will greatly expand the number of characterized acc-
SySt and thus provide the community with large enough samples for a
statistical study of the population of symbiotic stars as a whole. In the
meantime, we encourage the readers to further investigate the 33 candidates
presented here and help confirm their symbiotic nature with further follow-up
observations.

##  7 Data availability

The data underlying this article will be shared on reasonable request to the
corresponding author.

## Acknowledgements

This research has made use of the ASI Science Data Center Multimission
Archive; it also used the NASA Astrophysics Data System Abstract Service,
which is operated by the Jet Propulsion Laboratory, California Institute of
Technology, under contract with the National Aeronautics and Space
Administration. This research has also made extensive use of the SIMBAD and
VIZIER databases operated at CDS, Strasbourg, France. We thank Dr. Jamie
Kennea and the Swift team for the quick approval and the rapid acquisition of
the observations we requested. We thank R. Jurdana for the initial contact,
and J. Grindlay and E. Los for grating access to DASCH data and scans prior to
publication. NM acknowledges financial support through ASI-INAF agreement
2017-14-H.0, and UM through MAINSTREAM - PRIN INAF 2017 "The origin of
lithium: a key element in astronomy" (P.I. Paolo Molaro). We would also like
to thank S. Dallaporta, G. Cherini and F. Castellani (ANS Collaboration) for
contributing to the photometric observations of SU Lyn. This work was
supported by the Swedish strategic research programme eSSENCE. G.T. 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. G.T. acknowledges financial support of the Slovenian Research Agency
(research core funding No. P1-0188 and project N1-0040). This research made
use of Astropy (http://www.astropy.org) a community-developed core Python
package for Astronomy (Astropy Collaboration et al., 2013, 2018).

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