11institutetext: Nicolaus Copernicus Astronomical Center, Polish Academy of
Sciences, Rabiańska 8, 87-100 Toruń, Poland, 11email<EMAIL_ADDRESS>22institutetext: Space Radio-Diagnostics Research Centre, University of Warmia
and Mazury, ul. Oczapowskiego 2, 10-719 Olsztyn, Poland 33institutetext:
Warsaw University Astronomical Observatory, Al. Ujazdowskie 4, 00-478
Warszawa, Poland

# Lithium in coalesced non-compact stars

Tomek Kamiński 11 Mirek Schmidt 11 Marcin Hajduk 22 Aleksandra Kiljan 33 Inna
Izviekova 11 Adam Frankowski 11

###### Abstract

Context. Galactic red novae are thought to be produced in stellar mergers
between non-compact stars, such as main-sequence stars and cool giants. They
are hoped to help in explaining physical processes involved in common envelope
evolution and stellar binary collisions.

Aims. We investigate the presence of lithium in three best-observed Galactic
red nova remnants. Explaining the origin of lithium may point to mixing
mechanism present before, during, or after the merger.

Methods. The lithium line at 6707.81 Å was compared to a feature of [Ca I at
6572.78 Å to derive relative abundances in circumstellar gas. Absolute
abundances were next calculated assuming the Solar calcium to lithium
abundance ratio.

Results. Lithium abundances were measured in the merger remnants of V838 Mon
with $A$(Li)=2.3, CK Vul with $A$(Li)=2.5, and V1309 Sco with $A$(Li)=1.8.

Conclusions. Lithium is overabundant in red novae suggesting that at least
some merger products activate the Cameron-Fowler mechanism whereby convective
mixing can reach the deep stellar interior. Whether deep convection and
associated diffusion alone or some other processes (e.g. spindown) can be
responsible for driving the Cameron-Fowler mechanism in the remnants requires
further studies. Early observations of lithium in V838 Mon hint that these
mechanisms activate early, perhaps already in the common envelope phase. These
observations should be taken into account in modelling these complex systems.

###### Key Words.:

Stars: abundances – circumstellar matter – Stars: mass-loss – Stars: peculiar
– Stars: individual: V838 Mon, CK Vul, V1309 Sco

##  1 Introduction

Red novae (or luminous red novae) are a group of eruptive stars that have been
postulated to outburst as stellar mergers of non-compact stars (Tylenda &
Soker, 2006; Tylenda et al., 2011a; Pastorello et al., 2019b). We currently
classify five Galactic objects in this group: CK Vul (a.k.a. Nova 1670), V4332
Sgr (Nova Sgr 1994), V838 Mon (Nova Mon 2002), V1309 Sco (Nova Sco 2008), and
EWS 2002-OGLE-BLG-360. While other eruptive stars, like supernovae and
classical novae, red novae are primarily powered by accretion and, in
consequence, do not reach high plasma temperatures. Their stellar remnants –
that is, the products of the coalescence – are cool giant or supergiant like
stars with effective temperatures of 2000–3000 K. Owing to this low
photospheric temperature and mass loss phenomena (at different stages), the
circumstellar media of the merger products are remarkably rich in cool gas,
often present in molecular form, and quickly become enshrouded in dust. The
circumstellar envelopes of red nova remnants are very complex as they combine
material lost prior, during, and after the merger. Many of the studied
structures are asymmetric, pointing into formation of disks, jets, wide-angle
outflows, and inhomogeneous winds. This complexity seems to hold the key to
our understanding of physical mechanisms driving stellar systems to mergers
and of the merger process itself.

Although still small in numbers, the Galactic red novae are rather an
inhomogeneous group of variables. While V838 Mon progenitor was a young ($<20$
Myr, Afşar & Bond, 2007) triple system (Kamiński et al., 2021b) with an 8 M☉
primary (Tylenda et al., 2005b), most of the other objects were systems of
solar mass stars which were evolved in the time of the outburst: V1309 Sco and
V4332 Sgr were contact systems of two subgiants (Stȩpień, 2011), while EWS
2002-BLG-360 and CK Vul were systems containing a red giant branch (RGB) star
(Tylenda et al., 2013; Kamiński et al., 2018b; MacLeod et al., 2022). The
timescales on which we have studied the remnants of Galactic red novae are
typically short, up to decades (V4332 Sgr erupted in 1994), but the outburst
of CK Vul was observed in 1670–72 (Hevelius, 1671; Shara et al., 1985), giving
us currently a timescale of over 350 years. The youngest Galactic red nova,
V1309 Sco, erupted in 2008. For a brief overview of the objects, see Kamiński
et al. (2018a).

Over a dozen extragalactic red nova analogs (a.k.a. luminous red novae or,
more broadly, intermediate luminosity optical/red transients) has been
identified in the Local Group (e.g., Bond & Siegel, 2006; Pastorello et al.,
2019a, b, 2021a, 2021b; Blagorodnova et al., 2017, 2020, 2021; Williams et
al., 2015, 2020; Kurtenkov et al., 2015; MacLeod et al., 2017; Smith et al.,
2016; Rau et al., 2007; Cai et al., 2022). However, information on their
progenitors and remnants is sparse. Due to lack of adequate spectral
observations, they are not the main subject of this study.

In recent years, red novae have been recognized as objects important for
understanding the physics of the common envelope evolution and stellar
mergers. The focus on red novae especially intensified when the V1309 Sco’s
progenitor, relatively well observed before the red nova outburst in 2008, was
recognized to had been an eclipsing contact binary inspiraling into a common
envelope phase (Tylenda et al., 2011a; Pejcha, 2014; Pejcha et al., 2016a).
Observations of red novae and their remnants are used to investigate the
mechanisms behind the common envelope evolution and merger events, including
the associated mass loss phenomena (e.g., Ivanova et al., 2013; Nandez et al.,
2014; Soker & Kaplan, 2021; Soker, 2021, 2020a, 2020b, 2019, 2016; Soker &
Kashi, 2012; MacLeod et al., 2022, 2018a, 2018b; MacLeod & Loeb, 2020; Pejcha,
2014; Pejcha et al., 2016a, b, 2017; Metzger & Pejcha, 2017; Iaconi et al.,
2018, 2019; Reichardt et al., 2019).

On the course of detailed studies of individual red nova remnants, we noticed
that at least three of them – the three best observed objects, that is V838
Mon, CK Vul and V1309 Sco – show clear signatures of lithium in their spectra.
Since lithium is generally destroyed in stars, the spectral features usually
imply lithium enhancement. Overabundance of lithium in stars, especially
evolved ones, is not entirely understood but generally points to effective
mixing mechanisms in stellar interiors (e.g., Magrini et al., 2021). Curious
whether lithium in red novae can identify similar mixing mechanisms and by
this shed light on the merger physics, we investigated the lithium features in
the Galactic red novae using methods described in Sect. 2. We next constrained
lithium content and discussed its origin in V838 Mon, CK Vul, and V1309 Sco in
Sects. 3, 4, and 5, respectively. We briefly comment on the presence of
lithium in V4332 Sgr in a subsection of Sect. 5. The fifth known Galactic red
nova, EWS 2002-BLG-360, is not discussed here as no relevant observational
data exist for this heavily dust-enshrouded object (Tylenda et al., 2013). We
draw our final conclusions in Sect. 6.

##  2 Methodology

We analyzed the presence of lithium in red nova remnants based on observations
of a single feature of 7Li I at 6707 Å which is formed in the circumstellar
gas, not in the photosphere. The feature is a tight doublet of lines separated
by 0.15 Å or 6.7 km s-1. This split is insignificant compared to the local
turbulence and kinematic broadening of the features in all analyzed objects.
We assume a weighted mean central wavelength of 6707.81 Å as the reference
wavelength of the doublet; we used the transitions’ oscillator strengths as
weights (Wiese & Fuhr, 2009) in the calculation. Stars rich in lithium
occasionally display also features near 6103.6 and 8126.5 Å which are non-
resonant transitions and are not expected to be strong in cool circumstellar
matter of our remnants. We do not find any photospheric Li I features at these
wavelengths, either. Additionally, the $\lambda$6103 transition of Li I
overlaps with a line of Ca I at 6102.72 Å and the information on the lithium
content cannot be easily extracted from this spectral range. The 8126.5 Å
transition is problematic because it is within a telluric band.

In order to derive relative and absolute (i.e. relative to H) abundances of Li
I we compared the $\lambda 6707$ feature to the resonance line of Ca I at
6572.78 Å. This Ca line is the only semi-forbidden transition among resonance
lines observed toward our sources, and thus it is the least affected by
saturation effects. Its profile is very similar to the Li I line profile in
all our targets. Calcium features have been used in the analysis of lithium
content in other objects (Wallerstein & Conti, 1969), including classical
novae (Izzo et al., 2022). However, in circumstellar and interstellar matter,
calcium is strongly depleted owing to condensation into dust, limiting its
usefulness as a reference element. For instance, in the outflow of the carbon
star CW Leo, calcium was found to be depleted by a factor of 300 with respect
to the solar abundances (Mauron & Huggins, 2010) and by over 3000 in the
diffuse interstellar medium (ISM; Whittet, 1992). Observed species with less
significant depletion, e.g. Na or K, even though are observed in the same
outflows of red nova remnants, have too saturated lines to be useful in our
analysis. If calcium is depleted into dust in the analyzed outflows, our
estimates on absolute abundances of lithium should be treated as upper limits.
Depletion of lithium into dust is, in contrast, very small. For instance,
Field (1974) found it is depleted by a factor of 2 in the ISM.

Assuming both absorption lines are optically thin or have very similar optical
depth, by measuring their equivalent widths, $W$, we can obtain the column
density of Li I relative to Ca I from

$\log\frac{N_{\rm LiI}}{N_{\rm CaI}}=\log\frac{W_{\rm Li}f_{\rm
CaI}\lambda_{\rm CaI}^{2}}{W_{\rm Ca}f_{\rm LiI}\lambda_{\rm LiI}^{2}}.$ (1)

Where $f$ and $\lambda$ are the oscillator strength and central wavelength of
the transition, respectively. We adopted $f_{\rm LiI}$=0.75 for the Li I line,
which is a sum of oscillator strengths of the doublet, and $f_{\rm
CaI}=5.1\cdot 10^{-5}$ for the Ca I line. Both species have similar ionization
potentials (5.4 and 6.1 eV for Li I and Ca I, respectively) and thus we assume
they have similar ionization fractions. We further assume that both species
are chiefly in the ground state. This is justified by very low temperatures,
of a few 100 K, in the analyzed media. Then, the derived column density ratio
can be interpreted as the relative abundance ratio. Adopting now the solar
abundance of Ca of $A$(Ca)=6.36 (Grevesse & Sauval, 1998) in our objects, we
can constrain the Li abundance relative to hydrogen as $A({\rm
Li})=\log(N_{\rm LiI}/N_{\rm CaI})+A({\rm Ca})$ (where $A({\rm X})=\log N_{\rm
X}/N_{\rm H}+12$).

##  3 Lithium in outflows of V838 Mon

The $\lambda$6707 line of Li I has been one of the strongest atomic features
in the post-outburst spectra of V838 Mon (see below), but it was also clearly
present during the outburst, which began at the beginning of January 2002 and
had a visual light curve with three peaks. Munari et al. (2002) found a broad
– thus circumstellar – feature of Li I as early as in late January, just
before the main (second) outburst of V838 Mon. In a quasi-photospheric
spectrum observed in March 2002, i.e. shortly after the third outburst, Kipper
et al. (2004) derived the lithium abundance of $A$(Li)=3.8. In an improved
analysis of the same observational material, Kipper & Škoda (2007) derived
$A$(Li)=3.4. Thus, the lithium abundance was certainly super-solar
($A_{\sun}$(Li)=1.1111In both papers, the authors compared V838 Mon’s lithium
abundance to the meteoritic Solar System abundance, not to the photospheric
values, and incorrectly concluded that V838 Mon had a nearly solar lithium
abundance., Grevesse & Sauval, 1998) and marginally super-meteoritic ($A_{\rm
met}$(Li)=3.31) in March 2002. We take a closer look at the circumstellar
lithium line monitored through high-resolution spectra obtained in much later
epochs, from 2005 to 2020.

The spectra used in our analysis represent five epochs and are listed in Table
1. As illustrated in Fig. 1, the Ca I profile follows closely the variations
in the Li I profile over the years. Both transitions, however, are located in
spectral ranges where it is difficult to assess the local pseudo-continua. The
presented profiles are normalized as a result of dividing the spectra by high-
order polynomials, which were fitted in arbitrarily selected spectral ranges.
The Li I line is located between the TiO $\gamma$ 1–0 $R_{1}$ and $R_{2}$
band-heads. Their shape influences the underlying “baseline”. Depending on the
actual shape of the TiO absorption bands, the Li I profile may possess or not
an emission component at velocities higher than 80 km s-1. This is illustrated
in Fig. 1 for the 2005 epoch, for which we present two alternative
normalizations of the Li I spectrum. This epoch was the most challenging case.

The Ca I line, on the other hand, is heavily polluted by overlapping
absorption bands of the TiO$\gamma$ (5–3, 6–4) and TiO$\gamma^{\prime}$ (0–1)
systems, which makes a correction for the baseline difficult, too. For both
species, we attempted a normalization through division by spectra of other
cool stars or by TiO simulations, but the temperature and velocity structure
of the TiO circumstellar gas of V838 Mon is too complex for this approach to
work satisfactorily well (but see an example in Fig. 6). We therefore
performed further analysis with spectra normalized by polynomials. In this
approach, we assume that the atomic line profile is formed by absorbing from
the TiO absorption spectrum, which may not be the case if TiO and neutral
atomic gas are well mixed along the line of sight. Full radiative transfer
including all opacity sources would be necessary to make a more realistic
analysis. Also, all profiles are contaminated by low-amplitude TiO features,
on the order of 0.1 of continuum, and Ca I line is additionally affected by
weak telluric features. All these shortcomings increase uncertainty in the
derived abundances, but do not change our overall conclusions.

Table 1: High-resolution observations of V838 Mon used in the analysis of the
Li I line.

Epoch | Obs. date (UT) | Telescope | Instrument | Resolution, $R$ | Reference
---|---|---|---|---|---
2005 | 13.10.2005 | Keck | HIRES | 34 000 | 1
2009 | 02.01, 13.02, 22.03.2009 | VLT | UVES | 60 000 | 2
2012 | 06.01.2012 | VLT | UVES | 12 000 | 3
2018 | 27.10.2018 | Subaru | HDS | 20 000 | 4
2020 | 23.12.2020 | SALT | HRS | 40 000 | 5

(1) Kamiński et al. (2009), (2) Tylenda et al. (2011b), (3) Kamiński et al.
(2018a), (4) Kamiński, in prep., (5) Kamiński et al. (2021b)

Figure 1: Profiles of the $\lambda$6572 line of Ca I and of the $\lambda$6707
line of Li I in V838 Mon. Normalization of the spectra is uncertain, in
particular for the Li I profile from 2005 which may display a profile with
(purple dashed) or without (purple solid) emission (see text). The alternative
profile normalization for the 2005 epoch shown with dashed line was divided by
1.3 for better clarity. The Ca I profile is contaminated by two weak telluric
features (light green), whose position and depth is different for all epochs,
but is close to the sample telluric spectrum presented as a green line. The
dashed vertical line indicates velocity of the SiO maser, which usually is
indicative of a systemic velocity. The left panel shows changes in each
profile over time, while the right panel compares profiles of both species for
two sample epochs. Spectral regions highlighted in cyan and yellow were
measured for relative line intensities (Table 2).

The profile evolution of both lines is very similar. The P-Cyg profile from
2005 changed into a double broad absorption-dominated feature in 2009. This
slight doubling had continued to 2012, but some time after that, the two
subcomponents had split entirely, so that since $\approx$2018 they have been
two well separated velocity components. Both components in both lines become
weaker with time, but the Ca I component at most negative velocities is barely
distinguishable from baseline variations in 2020. At a closer look, the
profiles of both species display, however, some differences. A Li I component
at extreme heliocentric velocities, i.e. $\leq-95$ km s-1, has no
corresponding absorption in the Ca I profile. Hereafter, we call this Li-rich
and Ca-poor component $V_{\rm maxL}$. A component, which formed the well
separated absorption in Ca I in later epochs, between about –100 and –55 km
s-1, is designated $V_{\rm maxC}$. The range of intermediate velocities,
between about –55 and +25 km s-1, where absorption nearly disappeared in later
epochs, is hereafter designated as $V_{\rm mid}$. All the components mentioned
so far can be assigned to outflows caused by the 2002 eruption that is known
to had been associated with mass loss at changing terminal velocities. In the
velocity range between about 25 and 75 km s-1, the latter being the systemic
velocity measured from SiO maser observations (Deguchi et al., 2005; Ortiz-
León et al., 2020), in late epochs we observe a well separated component which
we assign to the ongoing mass loss (wind) from the coalesced star (cf. Tylenda
et al., 2005a; Kamiński et al., 2021b). We designated it $V_{\rm wind}$. All
components are graphically presented in Fig. 1. There may be an extra narrow
component at velocities higher than the maser, indicative of a fallback onto
the star, but it is too weak to be analyzed here.

Table 2: Equivalent widths of the Li I and Ca I lines in V838 Mon.

|  $W$ (Å) for Li I
---|---
Component | 2005aa$a$Spectrum was normalized as showed with the dashed line in Fig. 1. | 2009 | 2012 | 2018 | 2020
totalbb$b$Intergration over the entire profile. | 1.8 | 3.4 | 3.6 | 1.1 | 1.1
$V_{\rm maxL}$ |  | 0.4 | 0.4 | 0.2 | 0.2
$V_{\rm maxC}$ |  | 0.8 | 0.8 | 0.3 | 0.2
$V_{\rm mid}$ |  | 1.4 | 1.5 | ¡0.1 | ¡0.1
$V_{\rm wind}$ |  | 0.8 | 0.8 | 0.6 | 0.6
|  $W$ (Å) for Ca I
Component | 2005 | 2009 | 2012 | 2018 | 2020
totalbb$b$Intergration over the entire profile. | 1.5 | 2.4 | 2.4 | 0.8 | 0.6
$V_{\rm maxL}$ |  | ¡0.1 | ¡0.1 | ¡0.1 | ¡0.1
$V_{\rm maxC}$ |  | 0.5 | 0.5 | ¡0.1 | ¡0.1
$V_{\rm mid}$ |  | 1.3 | 1.3 | 0.3 | 0.2
$V_{\rm wind}$ |  | 0.5 | 0.6 | 0.47 | 0.4
| log of the Li to Ca abundance
Component | 2005 | 2009 | 2012 | 2018 | 2020
totalbb$b$Intergration over the entire profile. | –4.1 | –4.0 | –4.0 | –4.0 | –3.9
$V_{\rm maxL}$ |  |  |  |  |
$V_{\rm maxC}$ |  | –4.0 | –4.0 |  |
$V_{\rm mid}$ |  | –4.2 | –4.1 |  |
$V_{\rm wind}$ |  | –4.0 | –4.1 | –4.0 | –4.0
222

We measured equivalent widths of the different components and derived the
relative and absolute abundances of Li I in V838 Mon. The results are
presented in Table 2. The ratios we obtain for different epochs and different
parts of the outflow are very consistent and on average give $\log(N_{\rm
LiI}/N_{\rm CaI})$=–4.0. This constrains an absolute abundance of lithium at
$A({\rm Li})$=2.3, an order of magnitude lower than that constrained from 2002
quasi-photospheric spectra. At the same time, in the component $V_{\rm maxL}$
at the highest outflow velocity, where no Ca I emission is observed, the
lithium abundance should be at least one order of magnitude higher, i.e. close
to or higher than the meteoritic value.

We have assumed that the analyzed lines are optically thin or at least have a
similar optical depth. By modelling of the Li I P Cyg features observed in
V838 Mon in 2005 using the same techniques as in Tylenda et al. (2005a) (their
Sect. 5.2), we found a mean optical depth of the $\lambda$6572 line of
$\tau\la 4$ — the line was only moderately thick. The lines very likely became
even optically thinner in later epochs due to expansion.

###  3.1 Origin of lithium in V838 Mon

The presence of lithium in the remnant of V838 Mon is especially perplexing
when one realizes that before 2002 the primary of the coalesced system was a
B3 main-sequences star of a mass of 8 M☉ and with a radiative envelope. There
should be practically no lithium in the primary. The observations show,
however, that lithium was present in the early phases of the outburst,
possibly at very high, nearly meteoritic, abundances. Our observations show
that lithium was indeed enhanced in the matter lost during the 2002 event, and
is present at a similar abundance in the wind of the coalesced star, which
indirectly indicates lithium is also present in the photosphere. We consider
generally two mechanisms that could be responsible for lithium in V838 Mon:
(i) internal lithium production in the common envelope system or in the
coalesced star and (ii) an external enrichment through the material of the
lower-mass star that merged with the B-type primary in 2002.

#### Internal lithium enrichment

The internal production scenario is justified by the similarity of the post-
outburst remnant of V838 Mon to stars on the asymptotic giant branch (AGB;
Soker & Tylenda, 2007) and to red supergiants. The coalesced star mimics well
a red giant or a red supergiant in the sense that it has a large luminosity of
104 L☉ and huge dimensions with a 460 R☉ radius (Kamiński et al., 2021b;
Chesneau et al., 2014), and has a convective envelope that puts it close to
the Hayashi limit on the Hertzsprung-Russel diagram. Lithium, or more
specifically 7Li, can be produced in some RGB and AGB stars through fast
mixing in the Cameron-Fowler mechanism (a.k.a. the “7Be-mechanism”) (Cameron,
1955; Cameron & Fowler, 1971; Sackmann & Boothroyd, 1992). In this process,
7Be is produced in the second proton-proton chain (pp II), via ${}^{3}{\rm
He}(\alpha,\gamma)^{7}{\rm Be}$, in deep internal layers of the star. In the
complete pp II cycle, beryllium is next transformed into 7Li, which in turn is
quickly destroyed in a reaction with a proton, ${}^{7}{\rm Be}({\rm
e}^{-},\nu)^{7}{\rm Li}(p,\alpha)^{4}{\rm He}$. However, fast mixing is
thought to be able to effectively transfer the freshly produced 7Be to outer
cooler layers where it decays into 7Li in conditions where lithium is not
entirely destroyed and can be farther dredged up to the surface. Convection
seems to be the major mixing mechanism explaining surface 7Li abundances in
some AGB stars (including hot bottom burning), while in RGB stars an extra
mixing mechanism is necessary, be it termohaline mixing, magneto-hydrodynamic
instabilities, or rapid rotation, to mention just a few of the proposed mixing
mechanisms (Charbonnel et al., 2010; Magrini et al., 2021). The currently
measured lithium abundance of 2.3 in V838 Mon is very close to average values
of 2.3 and 2.6 measured in RGB and AGB giants (Maciel & Costa, 2012). We
propose that the 2002 event turned the main-sequence star of 8 M ☉, i.e. a
star with a radiative envelope, into a giant/supergiant with a fully
convective envelope that can reach deep enough to dredge up 7Be and produce
7Li in the photosphere. This way, V838 Mon would be a more extreme case of a
scenario considered by Denissenkov & Weiss (2000) in which a planet or brown-
dwarf engulfment can trigger an active Cameron-Fowler mechanism; in the case
of V838 Mon a much heavier body was accreted, making this process even more
effective.

Given that high lithium abundances were already observed in V838 Mon during
the 2002 outburst, the hypothesized 7Li production mechanism must have been
activated very early on. It is tempting to assign this behavior to the common
envelope phase, which the system went through just before the coalescence.
Whether the presence of lithium enrichment is related to V838 Mon being a
merger in addition to being a common event is an open question which would
require dedicated modeling. We note, however, that convective mixing that does
not reach the H-burning regions would actually decrease the abundance of Li in
the photosphere (by dredging up Li-poor material from internal layers where
lithium is easily destroyed by protons).

Alternatively, rotation of the remnant may provide extra mixing that can
support the Cameron-Fowler mechanism. Some stellar merger models predict that
the collision product contains a lot of angular momentum and thus is rapidly
rotating (e.g. Schneider et al., 2019). There is, as of now, no direct
indication of a fast rotation of V838 Mon, with an upper limit of $\varv\sin
i<28$ km s-1), but if this is ever found to be the case, internal rotation
might provide extra mixing. Moreover, since such a star is also supposed to
spin down with time, extra mixing is even more likely, but perhaps operates on
longer time scales than the 20 yr life-span of the remnant of V838 Mon.

#### External lithium enrichment

In the V838 Mon merger scenario of Tylenda & Soker (2006), the main sequence
B-type star (or an A-type protostar) collides with a low-mass protostar of a
mass of 0.1–0.5 M☉. The latter mass constraint is very uncertain, since it was
calculated as a mass needed to be accreted on the 8 M☉ star to account for the
observed power of the outburst. The young age of the progenitor is implied by
its membership in a small open cluster (Afşar & Bond, 2007), which is still
surrounded by large quantities of molecular gas and dust (Kamiński et al.,
2011; Tylenda & Kamiński, 2012). The presence of early B-type main-sequence
stars implies an upper limit on the age of the cluster of $\leq$25 Myr, but
observations of the translucent molecular cloud near the cluster suggest an
even younger age, 3–10 Myr. The youngest low-mass protostars are known to be
lithium-rich ($A({\rm Li}))\approx 3$) but, as they are fully convective, they
lose lithium very quickly during their evolution to the main sequence after
reaching core temperatures above about 3 MK. It is estimated that the Li
depletion timescale is the quickest for stars of a mass of 0.6 M☉, which
deplete lithium in about 20 Myr (White & Hillenbrand, 2005). The time scales
are comparably quick for masses considered for the accreted star in V838 Mon
(Bildsten et al., 1997). For instance, for a 0.3 M☉ protostar, 100-fold Li
depletion is expected to take 28 Myr and a lithium reduction by a factor of 2
takes 18 Myr (Chabrier et al., 1996).

If the age of the V838 Mon progenitor is at the lower end of the suggested
range, i.e. $<$10 Myr, the low mass star could still have been rich in Li.
This lithium-rich material could have been seen during the outburst and is
seen now in the wind (and thus photosphere) of V838 Mon. Stellar-merger
simulations show that the matter of the lower-mass body forms the outer layers
of the merger product (e.g. Tylenda & Soker, 2006; Sills et al., 2001), which
would explain the persistent presence of Li in the remnant. Ongoing convection
and other mixing mechanisms that do not tap the hot 3He-rich layers will
decrease the Li abundance with time. Additionally, the low-mass protostar
could have been surrounded by a protoplanetary disk or other form of
circumstellar matter which is expected to be Li rich, as it would be enhanced
in lithium produced by spallation. We calculate that to produce the observed
lithium enhancement in an envelope of 0.1 M☉, the primary would have to
accrete 0.01 M☉ or 10 Jupiter masses of material with meteoritic abundance of
lithium. This is not impossible given typical masses of protoplanetary disks
of low-mass stars (e.g Andersen et al., 2019).

If the progenitor age is closer to 25 Myr or the actual lithium depletion is
more effective in low-mass stars than described by theoretical models, the
internal production of lithium remains the more likely mechanism responsible
for the observed lithium in V838 Mon. Some observations indicate that low-mass
protostars (as young as 8 Myr) may be Li depleted sooner than theoretically
expected (White & Hillenbrand, 2005; Yee & Jensen, 2010; Jeffries et al.,
2017). Because lithium enhancement is seen in other remnants of red novae, we
favor the scenario where lithium is commonly produced by coalesced stars
internally.

The proposed scenarios result in lithium abundances with different 7Li/6Li
isotopic ratios. While the external enhancement should result in the cosmic
7Li/6Li ratio of about 7, the internal production of lithium should
considerably increase the content of the 7Li isotope only. Unfortunately,
deriving the lithium isotopic ratio is very difficult even for sharp
interstellar lines, and it is hopelessly hard, if not impossible, for the wind
features in V838 Mon. The internal production of lithium predicts also
anomalous abundance of Be and B, but lines of those species are hardly
available for observations.

###  3.2 Li-bearing molecules

With the premise that lithium may be very highly enhanced in the circumstellar
medium of V838 Mon, we investigated whether there may be an enhanced presence
of Li-bearing molecules in the remnant. A detection of such molecules would be
instrumental in determining the dominant isotope of lithium. Overall,
observations of Li-bearing species in astronomical environments are very rare
(Muller et al., 2020). We considered two species, LiOH and LiCl, which are
thought to be among the main Li-bearing species that are easily accessible for
millimeter-wave spectroscopy through their pure rotational transitions. We
used the GGchem code (Woitke et al., 2018) to investigate molecular abundances
in gas at different temperatures and pressures, assuming the thermochemical
equilibrium. The latter assumption is questionable for the circumstellar gas
of the red nova remnant, but we used the calculation only as a first
guideline. We find that even if lithium abundance were extremely high, with
$A({\rm Li})$=3–5, the Li-bearing molecules would have low absolute abundances
and would produce lines that are not easily detectable with modern
submillimeter interferometers like ALMA. In particular, LiOH emission should
be 1–3 orders of magnitude weaker than AlOH lines detected in V838 Mon at a
modest S/N (Kamiński et al., 2021b).

##  4 Lithium in the CK Vul outflow

The photosphere of the stellar remnant of Nova 1670 is not observable, as it
is deeply embedded in circumstellar dust. The presumed merger product is
surrounded by dusty bipolar nebula, which contains cool molecular gas as well
as shock-excited atomic and H2 matter (Kamiński et al., 2020). Hajduk et al.
(2013) noticed two variable stars towards the southern lobe of that nebula.
The variability can be only explained if these are far background stars
affected by changing extinction originating from the rapidly expanding
southern lobe of CK Vul. In an optical spectrum of the brightest of the
variable stars, Hajduk et al. found a strong lithium feature near 6707 Å which
most likely originates in the outflow of CK Vul. Here we take a closer look at
this feature, aiming to constrain the abundance and origin of lithium in the
circumstellar medium of CK Vul.

We used a spectrum of Hajduk et al. (2013) obtained with the Gemini
Multiobject Spectrograph (GMOS) in 2010 to constrain first the parameters of
the background star. We assume that the spectrum is correctly calibrated in
fluxes, and we made no extra corrections for chromatic slit losses or effects
caused by the atmospheric chromatic aberration. Using a grid of synthetic
spectra of AMBRE (de Laverny et al., 2012), we found that the spectrum is best
approximated by a photosphere of an effective temperature of 4750$\pm 250$ K
(near spectral type $\sim$K3) and reddening (interstellar plus circumstellar)
of $A_{V}$=6.0 mag. Results of the best-fit search are shown in Fig. 7. We
used the standard interstellar extinction curve (Cardelli et al., 1989) with
the selective extinction parameter of $R_{V}$=3.1. Since no uncontaminated
photospheric features are present, the luminosity class of the photosphere
remains unconstrained. The effective temperature is rather well constrained.
Cooler stars would produce a conspicuous MgH band near 5200 Å, which is not
present, and hotter stars would be betrayed by a H$\alpha$ line which would be
much stronger than the observed one. The derived temperature is close to that
of a spectral type K4 ($T_{\rm eff}\\!\approx$4500 K) assigned to the star by
Hajduk et al. However, the $V$-band reddening of 6.0 mag we derived is
significantly higher than 4.4 mag postulated earlier.

Figure 2: Profiles of absorption lines towards the star located behind the
southern lobe of the CK Vul outflow. Left panel shows circumstellar absorption
of Li I and Ca I in the circumstellar gas (upper) and interstellar features
(shifted down by 0.3). Note different central velocities of the two groups.
The right panel shows features of the Na I doublet and of H$\alpha$ which
combine circumstellar, interstellar, and photospheric absorption. The filled
cyan area shows absorption common for both lines of the blending Na I doublet
and approximates the shape of the intrinsic Na I absorption.

Hajduk et al. measured the $R$ magnitude of the background star to be lower
than 19 mag in 2010. Adopting a spectral type K3 and assuming the star is on
the main-sequence, we would expect $(V-R)$ of 0.8 mag and then the dereddened
$V$ magnitude would be 18.6. Assuming now an absolute magnitude of 6.9 mag for
a K3 V star, we obtain a distance of 2.2 kpc. Gaia parallax measurements of
the background star (0.0660$\pm$0.1858 mas Lindegren et al., 2021) imply a
3$\sigma$ lower limit of 2.6 kpc for the distance, so a luminosity class V can
still be reconciled with the Gaia limits. However, higher luminosity classes,
like giants, are more likely because the distance to CK Vul is much larger
than 3 kpc (Kamiński et al., 2021a; Banerjee et al., 2020). Since only a small
fraction, 1–2%, of field giants have strong lithium features in their spectra
(e.g. Smiljanic et al., 2018), it is very unlikely we observe one here.

There are no clear photospheric features in the spectrum of the background
star, but, as mentioned, Na I and H$\alpha$ seem to have a photospheric
component partially blending with interstellar and outflow absorption
components. By comparing profiles of the Na I lines and H$\alpha$ to
circumstellar and interstellar lines, we tentatively assign the photosphere a
heliocentric velocity of $\approx$–250$\pm$50 km s-1. This is a much more
negative velocity than the dominant interstellar and circumstellar components
(cf. Fig. 2).

The spectrum contains pure interstellar features. Hajduk et al. mentioned that
the strongest diffuse interstellar band (DIB) is near 6284 Å, but we find it
too close to a telluric band. Instead, the strongest DIB we have identified is
the one near 5780.59 Å. There is also possibly a DIB at 6613.56 Å.

The most interesting for the discussion here are the circumstellar components.
The lithium line $\lambda$6707 is clear from any contamination – the closest
features known in this source from other observations are the $\lambda$6716
line of the [S II] doublet and $\lambda$6678 of He I, which, fortunately, do
not overlap with the Li line. Given the virtual absence of other photospheric
lines, the line of lithium is very unlikely to be contaminated by absorption
intrinsic to the field star, even if it happened to be a Li-rich giant. The Ca
I line, when observed at the modest resolution of the GMOS spectra, is
partially contaminated by an H$\alpha$ feature, which itself combines
photospheric absorption and absorption from the CK Vul lobe (at different
velocities; possibly also weak emission may be present). Despite this
complication and as shown in Fig. 2, the velocity profiles of the Li I and Ca
I lines overlap closely in velocity. This suggests that they are formed in the
same region with a central heliocentric radial velocity of about –110 km s-1
and the lines can be directly compared. Their radial velocity is certainly
different from that of the interstellar gas probed by DIBs, which we measure
is centered near 40 km s-1. They also overlap with the velocity range between
–180 and 0 km s-1 occupied by the undoubtedly circumstellar CO $J$=1–0
emission observed by ALMA along the same line of sight.

The equivalent widths of $W_{\rm LiI}$=1.3 and $W_{\rm CaI}$=0.6 (both
$\pm$0.1) Å yield $\log(N_{\rm LiI}/N_{\rm CaI})$=–3.8 ($\pm$0.1). In a
standard solar composition, this implies a lithium abundance of $A$(Li)=2.5.
This lithium abundance is by 1.4 dex higher than in the Sun, but sill well
below the meteoritic value.

###  4.1 Interpretation of Li in CK Vul

Formation of Li in CK Vul may be different from that in V838 Mon. There is a
large body of evidence that CK Vul’s circumstellar material has been heavily
processed by nucleosynthesis processes associated with the progenitor and the
merger event. Observations of optical atomic lines of the inner nebular
remnant indicate a highly processed elemental composition. Foremost, helium is
approximately twice more abundant relative to hydrogen than in the Solar
photosphere and the N/O ratio is nearly ten times higher than in the Sun
(Tylenda et al., 2019). Millimeter-wave observations of cool molecular gas of
the remnant of CK Vul also suggest a general overabundance of nitrogen
(Kamiński et al., 2015b). They additionally yield numerous isotopic ratios,
for instance 12C/13C$\approx$3.8, 14N/15N$\approx$20, and 16O/18O$\approx$36
(Kamiński et al., 2017) which together are very unusual. The chemical
characteristics can be coherently explained by processing in the CNO hydrogen-
burning cycles and by incomplete helium burning. Moreover, based on an
unprecedented detection of an AlF isotopologue containing 26Al, which implies
27Al/26Al$\approx$7, it was postulated that the progenitor system included an
RGB star (Kamiński et al., 2018b). (Currently, the most likely scenario is a
merger of an RGB star with a He white dwarf; Tylenda et al., in prep.). The
analysis of optical spectra of heavier elements (Ne, S, and Ar) show
additionally that the progenitor had a low metallicity with
[Fe/H]=–0.53$\pm$0.07 (Tylenda et al., 2019). While the general
nucleosynthesis channels responsible for the peculiar composition can be
identified, their relation to the stellar nucleosynthesis, nuclides processing
during the presumed merger, or how exactly the matter was dispersed are
currently not understood.

The level of overabundance of Li in CK Vul is close to these observed in other
Li-enriched objects, including RGB and AGB giants (J-type stars among them),
classical novae, or R CrB (RCB) stars. Relevant for the discussion here is
especially the similarity to RCB stars and J-type carbon stars, which are
thought to be products of mergers between either CO and He white dwarfs (RCB
stars), or He white dwarfs and red giants (J stars). The merger of CK Vul was
certainly different from that postulated for RCB stars, chiefly due to the
presence of a hydrogen envelope, but models of the latter still offer an
informative analogy. In nucleosynthesis models of merging white dwarfs (WDs),
7Li is produced from 3He through 7Be, i.e. ${}^{3}{\rm
He}(\alpha,\gamma)^{7}{\rm Be}({\rm e}^{-},\nu)^{7}{\rm Li}$, just as in the
Cameron-Fowler mechanism. White dwarf mergers are thought to effectively
produce 7Be and thus 7Li only if 3He is abundantly available from earlier
processing or directly from the matter of the He WD. It is also expected that
the Li-rich region is mainly located in the outer layers of the merger product
(Longland et al., 2012), which gives a chance of identification of the species
in stellar spectra.

Models of mergers leading to the formation of J (and early-R) type carbon
stars – that is, those that involve a red giant and thus are closer analogs of
CK Vul – are similar to RCB formation models. They produce lithium only
provided that large quantities of 3He are mixed down to the He-burning layer
and that freshly produced lithium has a way to escape to the outer layers of
the coalesced star (Zhang & Jeffery, 2013; Zhang et al., 2020). In those
models, lithium abundances steeply increase over the first $\sim$1000 yr after
the merger and can reach super-meteoritic values. The lithium abundance of 2.5
dex in CK Vul’s outflow observed some 350 yr after the merger can be
reconciled with these models, but more studies on the kinematic age of the Li-
bearing outflow are necessary to assure this, as there is chance that the
outflow is younger than 350 yr (Kamiński et al., 2021a). As mentioned,
observations of CK Vul indicate an abundance of He that amounts to over half
the mass of the optical nebula. There was therefore enough raw material to
produce Li in the ${}^{3}{\rm He}(\alpha,\nu)^{7}{\rm Li}$ process.

Dedicated merger models are necessary to verify the origin of Li in the
extended bipolar remnant of CK Vul.

##  5 Lithium in the emission spectrum of V1309 Sco

V1309 Sco is an important member of the red nova group, as it is the only
object with well documented photometric variability prior to the merger in
2008 (Tylenda et al., 2011a). Located in a dense field, it is a challenging
source for optical observations. After around 2010, its spectral evolution
could have been traced mainly through emission lines with no direct
observations of the stellar continuum, which only becomes dominant in the mid-
infrared (Kamiński et al., 2015a, 2018a).

Lithium features have not been reported in the spectra of V1309 Sco before
this work. We find the $\lambda$6707 feature in emission in an X-shooter
spectrum from May–July 2016. The spectrum was reported by Kamiński et al.
(2018a). Parts of the spectrum covering the lithium and calcium lines of our
interest are shown in Fig. 3. The Ca I line is very close to H$\alpha$ with
both lines having nearly equal strength, but the Ca I line is relatively clear
from contamination by H$\alpha$. The Li I line partially overlaps with a broad
blend composed of the TiO $\gamma$ 1–0 band and a line of [S II] (at a rest
wavelength of 6716.44 Å). As shown in Fig. 4, the lithium and calcium lines
have almost the same profiles, but the exact appearance relies on an arbitrary
definition of the local baselines. The measured line fluxes of $F_{\rm
CaI}$=33.2$\pm$0.7 and $F_{\rm LiI}$=4.8$\pm$0.4 10-16 erg s-1cm-2 yield a
calcium to lithium flux ratio of 6.9$\pm$0.6, but uncertainty in this value is
larger than indicated owing to uncertain baseline levels for both lines.

The spectrum of the V1309 Sco remnant is dominated by emission lines formed
through resonant scattering (Chandrasekhar), whereby absorption of a stellar
photon from the ground state is followed by emission to the ground. The
emission is thus coupled to the incident light of the star. This kind of
spectrum was well studied in V4332 Sgr, a spectroscopic twin of V1309 Sco
(Kamiński & Tylenda, 2013; Tylenda et al., 2015). Assuming that the lines of
Li I and Ca I are formed in the same volume of gas, the number density ratio
of both species can be calculated as:

$\frac{n_{\rm LiI}}{n_{\rm CaI}}=\frac{F_{\rm Li}f_{\rm CaI}\lambda_{\rm
CaI}^{2}}{F_{\rm CaI}f_{\rm LiI}\lambda_{\rm LiI}^{2}}\frac{F^{*}_{\rm
CaI}}{F^{*}_{\rm LiI}}\exp{-(\tau_{\rm CaI}-\tau_{\rm LiI})}.$

Here, $F$ and $F^{*}$ represent the integrated line fluxes and the incident
starlight fluxes at the wavelengths of the lines. Since the photospheric
spectrum of V1309 Sco was not observed directly (being attenuated by
circumstellar dust since before 2012, see Kamiński et al., 2015a), the ratio
of the intrinsic stellar fluxes is unknown. Nevertheless, due to many
analogies between V1309 Sco and V4332 Sgr, we can assume that the stellar
spectrum is close to a giant with an effective temperature of 3200 K, as
determined for V4332 Sgr in Kamiński et al. (2010). Synthetic spectra of
giants at this temperature (Gustafsson et al., 2008) show a flux ratio
$F^{*}_{\rm CaI}/F^{*}_{\rm LiI}\approx$1.5. However, the incident spectrum
that is scattered by the Ca I and Li I lines is already affected by
circumstellar absorption of TiO which greatly increases the ratio. For
instance, in the spectrum of V838 Mon observed in 2005, this ratio of about
4.5 (Kamiński et al., 2009, their Fig. 3). We therefore adopt a value of
3.0$\pm$1.5 for V1309 Sco.

Although the calcium line is a semi-forbidden transition, it can have a
significantly higher optical depth than the seven times weaker line of
lithium. Because we do not have a reliable constraint on the size of the
emission region, calculating line opacities solely from the observed fluxes is
not possible, and thus we ignore the saturation effects, that is, take
$\tau_{\rm CaI}$=$\tau_{\rm LiI}$=0.

We obtain the density ratio $n_{\rm LiI}/n_{\rm CaI}$=2.8($\pm
1.7$)$\cdot$10-5 and the absolute abundance $A$(Li)=1.8($\pm 0.3$). The
lithium abundance is a few times higher than in the Sun ($A_{\sun}$(Li)=1.1),
but is also a few times lower than in the two other red nova remnants analyzed
here.

Figure 3: Spectra of V1309 Sco from 2016 near the lithium and calcium
resonance lines. Neighboring features are identified. To better illustrate the
contribution of TiO and CrO emission, simple simulations of their emission are
shown in orange and green, respectively. The simulations were performed as in
Kamiński et al. (2015a). Spectra were dereddened, but were not corrected for
the spectroscopic baseline (which is not the true continuum of V1309 Sco).
Figure 4: Profiles of the Ca I and Li I resonance lines in X-shooter
observations of V1309 Sco from 2012 (upper) and 2016 (lower). The Ca I flux in
2016 was divided by 7. In 2016, at velocities higher than about 0 km s-1, the
Li I line partially overlaps with a blend of TiO and [S II].

Lines of both species are centered at –130 ($\pm$10) km s-1 and occupy a
velocity range typical for most other neutral resonance lines in the late 2016
spectrum of V1309 Sco, for example the K I doublet (cf. Kamiński et al.,
2015a, 2018a). This velocity is also close to our best constraints on the
systemic heliocentric velocity of the object, that is, –110 or –120 km s-1
(Kamiński et al., 2018a; Mason & Shore, 2022). As shown in Fig. 4, the calcium
line was also present in a spectrum obtained in 2012 and reported in Kamiński
et al. (2015a). However, the lithium line was not detected in 2012, most
likely due to a poorer signal-to-noise ratio—the non-detection is consistent
with the line ratio we measured for the 2016 features. Even earlier spectra,
as those from 2009 presented by Kamiński et al. (2015a), do not show any
feature near 6707 Å. We also analyzed UVES spectra from Mason et al. (2010)
taken during the outburst of V1309 Sco, that is, in September–October 2008,
but do not find any features at velocities near –130 km s-1 that could be
assigned to the Ca I and Li I resonance lines. However, multiple spectra in
this dataset show broad absorption features centered at $\approx$–105 km s-1
which can be assigned to Ca I and which may have a very faint corresponding
feature of lithium. We find these lines in spectra from 18, 20, 28 September
2008, 8 and 20 October 2008 – but not in spectra from early September 2008.
Sample spectra are shown in Fig. 5. The relatively low expansion velocity of
the calcium-bearing gas places it within the “low-velocity component” of Mason
& Shore (2022), which – according to these authors – originates in material
lost from the V1309 Sco system before the red nova outburst and which was
excited in 2008–2009 by the radiation of the outburst. If the weak features
that we have identified indeed belong to Li I, the enrichment in lithium may
even precede the 2008 eruption of V1309 Sco, possibly when the progenitor
system of subgiants (Stȩpień, 2011) was losing mass. Some solar-mass subgiants
are known to have super-solar lithium abundances (Randich et al., 1999), and
thus the progenitors could have been a direct source of the lithium reservoir.
However, the spectra from the outburst phase 2008–2009 are very complex, and
it is difficult to make unambiguous identifications of the observed features.
While it is certain that the remnant displays a lithium enrichment in
circumstellar gas in 2016, it is uncertain for dates earlier than 2016.

If we accept that lithium was only seen in late type spectra, its production
can be similar to that discussed for V838 Mon, that is, through internal
production due to enhanced mixing caused by the merger or premerger phases.

Figure 5: Sample spectra near the Li I and Ca I resonance lines as observed by
UVES in 2009 in V1309 Sco (Mason et al., 2010). The observing epochs are
indicated as UT times for each pair. Spectra were not normalized and for each
pair are in scale. The area highlighted in pink is where strongest Ca I
absorption is observed and where Li I may have a weaker corresponding feature.
The vertical dotted line marks the central velocity of this component. The
dashed line marks the peak position of the Li I and Ca I emission lines in the
2016 spectrum.

#### V4332 Sgr

V4332 Sgr is a close analog of V1309 Sco (Kamiński et al., 2015a, 2018a). Its
outburst was observed in 1994, that is, twelve years before that of V1309 Sco.
No lithium feature has ever been reported in the spectra of V4332 Sgr. Spectra
obtained during and soon after the 1994 eruption (Martini et al., 1999;
Kimeswenger, 2006; Banerjee et al., 2015; Tylenda et al., 2005a) do not show
any conspicuous features that could be assigned to Li I $\lambda$6707\.
However, most of these spectra have very modest spectral resolutions (¡10 000)
and signal-to-noise ratios, and thus we cannot definitively claim absence of
lithium. Because V4332 Sgr is an oxygen-rich object, its strong photospheric
and circumstellar bands of TiO $\gamma$ 1–0 hamper any meaningful tests on the
presence of Li I $\lambda$6707 in the poor quality spectra. Also, there was a
gap of eight years between the earliest spectra obtained in 1994 by Martini et
al. (1999) and the next one, i.e. from August 2002 (Kimeswenger, 2005). The
best-quality spectrum of the V4332 Sgr remnant was obtained with UVES in 2005
(Tylenda et al., 2015) and it does not show any Li features. If the Li I line
was a temporal feature in V4332 Sgr, it could have been missed.

##  6 Conclusions

Some overabundance of lithium is observed in all the three well-observed red-
nova remnants, V838 Mon, V1309 Sco, and CK Vul. The origin of the enhancements
is however not clear and, as discussed, may have a different origin in
different objects. However, it is always liked to the Cameron-Fowler mechanism
that requires deep mixing. Because in some objects there are hints of Li
overabundance during the outburst (V838 Mon) and possibly even before it
(V1309 Sco), some convective configuration may be responsible. Since common
envelope evolution appears to be a common feature in these objects, deep
convection in that phase is a promising scenario explaining lithium in red
novae. With this observational work, we would like to encourage theoretical
studies of lithium production through activation of the Cameron-Fowler
mechanism in common-envelope systems, merger events, and in merger remnants.

The curious presence of lithium in red novae and their remnants does not
exhaust links between lithium production and common-envelope evolution or
stellar mergers. Lithium enhancements are often found in non-binary giants
with a history of enhanced mass loss and with rapid rotation, including small
subsets of RGB stars (e.g. de La Reza et al., 1996) and AGB stars (e.g.
Reyniers & Van Winckel, 2001). A link of some of them to red novae was
recently discussed by Melis (2020). It also appears now that mergers are the
only sensible way to explain J-type and early R-type carbon stars rich in
lithium (Izzard et al., 2007; Zhang et al., 2020). Finally, the resonance line
of lithium was observed in extragalactic gap transients, including an
intermediate-luminosity red transient AT 2019abn (Williams et al., 2020) and
possibly a luminous red nova AT 2017jfs (Pastorello et al., 2019a, priv.
comm.), which may as well be powered by mergers. Lithium appears to be an
important element for understanding a wide range of systems suspected of
common envelope evolution or mergers.

This study identifies red novae and their remnants as copious lithium
producers. However, given their rarity, now estimated to two bright outbursts
per decade (Kochanek et al., 2014), they are unlikely to be key players in the
Milky Way lithium budget.

The derivation of circumstellar lithium abundances was rather difficult in red
nova remnants. Potential depletion of calcium into dust make our estimates of
$A$(Li) upper limits, and further studies of dust formation in these objects
may improve on the uncertainties. Radiative transfer modelling of the spectra
of these objects in different phases are underway to better address line
opacity effects.

###### Acknowledgements.

We are grateful to: Th. Rauch for kindly sending us spectra from the early
outburst of V838 Mon; P. Martini for providing us with spectra of V4332 Sgr;
A. Pastorello for sharing the spectra of AT 2017jfs; and Y. Cai for sharing
data on AT 2021biy. We thank R. Tylenda for comments on the manuscript and Y.
Pavlenko for attempts in modeling the high-excitation lines of lithium in V838
Mon. T.K. and I.I. acknowledge funding from grant SONATA BIS no
2018/30/E/ST9/00398 from the Polish National Science Center. Based on
observations collected at the European Organisation for Astronomical Research
in the Southern Hemisphere under ESO programmes 382.D-0152, 088.D-0112,
60.A-9445, 089.D-0041, 097.D-0092, and 281.D-5055. This research is based in
part on data collected at the Subaru Telescope, which is operated by the
National Astronomical Observatory of Japan. We are honored and grateful for
the opportunity of observing the Universe from Maunakea, which has the
cultural, historical, and natural significance in Hawaii. Some data presented
herein were obtained at the W. M. Keck Observatory, which is operated as a
scientific partnership among the California Institute of Technology, the
University of California and the National Aeronautics and Space
Administration. The Observatory was made possible by the generous financial
support of the W. M. Keck Foundation. The authors wish to recognize and
acknowledge the very significant cultural role and reverence that the summit
of Maunakea has always had within the indigenous Hawaiian community. We are
most fortunate to have the opportunity to conduct observations from this
mountain. Some observations reported in this paper were obtained with the
Southern African Large Telescope (SALT). Polish participation in SALT is
funded by grant No. MNiSW DIR/WK/2016/07. Based on observations obtained at
the international Gemini Observatory, a program of NSF’s NOIRLab, which is
managed by the Association of Universities for Research in Astronomy (AURA)
under a cooperative agreement with the National Science Foundation on behalf
of the Gemini Observatory partnership: the National Science Foundation (United
States), National Research Council (Canada), Agencia Nacional de Investigación
y Desarrollo (Chile), Ministerio de Ciencia, Tecnología e Innovación
(Argentina), Ministério da Ciência, Tecnologia, Inovações e Comunicações
(Brazil), and Korea Astronomy and Space Science Institute (Republic of Korea).

## References

  * Afşar & Bond (2007) Afşar, M. & Bond, H. E. 2007, AJ, 133, 387
  * Andersen et al. (2019) Andersen, B. C., Stephens, I. W., Dunham, M. M., et al. 2019, ApJ, 873, 54
  * Banerjee et al. (2020) Banerjee, D. P. K., Geballe, T. R., Evans, A., et al. 2020, ApJ, 904, L23
  * Banerjee et al. (2015) Banerjee, D. P. K., Nuth, Joseph A., I., Misselt, K. A., et al. 2015, ApJ, 814, 109
  * Bildsten et al. (1997) Bildsten, L., Brown, E. F., Matzner, C. D., & Ushomirsky, G. 1997, ApJ, 482, 442
  * Blagorodnova et al. (2020) Blagorodnova, N., Karambelkar, V., Adams, S. M., et al. 2020, MNRAS, 496, 5503
  * Blagorodnova et al. (2021) Blagorodnova, N., Klencki, J., Pejcha, O., et al. 2021, A&A, 653, A134
  * Blagorodnova et al. (2017) Blagorodnova, N., Kotak, R., Polshaw, J., et al. 2017, ApJ, 834, 107
  * Bond & Siegel (2006) Bond, H. E. & Siegel, M. H. 2006, AJ, 131, 984
  * Cai et al. (2022) Cai, Y. Z., Pastorello, A., Fraser, M., et al. 2022, arXiv e-prints, arXiv:2207.00734
  * Cameron (1955) Cameron, A. G. W. 1955, ApJ, 121, 144
  * Cameron & Fowler (1971) Cameron, A. G. W. & Fowler, W. A. 1971, ApJ, 164, 111
  * Cardelli et al. (1989) Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245
  * Chabrier et al. (1996) Chabrier, G., Baraffe, I., & Plez, B. 1996, ApJ, 459, L91
  * Charbonnel et al. (2010) Charbonnel, C., Tosi, M., Primas, F., & Chiappini, C., eds. 2010, Light Elements in the Universe (IAU S268), Vol. 268
  * Chesneau et al. (2014) Chesneau, O., Millour, F., De Marco, O., et al. 2014, A&A, 569, L3
  * de La Reza et al. (1996) de La Reza, R., Drake, N. A., & da Silva, L. 1996, ApJ, 456, L115
  * de Laverny et al. (2012) de Laverny, P., Recio-Blanco, A., Worley, C. C., & Plez, B. 2012, A&A, 544, A126
  * Deguchi et al. (2005) Deguchi, S., Matsunaga, N., & Fukushi, H. 2005, PASJ, 57, L25
  * Denissenkov & Weiss (2000) Denissenkov, P. A. & Weiss, A. 2000, A&A, 358, L49
  * Field (1974) Field, G. B. 1974, ApJ, 187, 453
  * Grevesse & Sauval (1998) Grevesse, N. & Sauval, A. J. 1998, Space Sci. Rev., 85, 161
  * Gustafsson et al. (2008) Gustafsson, B., Edvardsson, B., Eriksson, K., et al. 2008, A&A, 486, 951
  * Hajduk et al. (2013) Hajduk, M., van Hoof, P. A. M., & Zijlstra, A. A. 2013, MNRAS, 432, 167
  * Hevelius (1671) Hevelius, J. 1671, Philosophical Transactions of the Royal Society of London Series I, 6, 2197
  * Iaconi et al. (2018) Iaconi, R., De Marco, O., Passy, J.-C., & Staff, J. 2018, MNRAS, 477, 2349
  * Iaconi et al. (2019) Iaconi, R., Maeda, K., De Marco, O., Nozawa, T., & Reichardt, T. 2019, MNRAS, 489, 3334
  * Ivanova et al. (2013) Ivanova, N., Justham, S., Chen, X., et al. 2013, A&A Rev., 21, 59
  * Izzard et al. (2007) Izzard, R. G., Jeffery, C. S., & Lattanzio, J. 2007, A&A, 470, 661
  * Izzo et al. (2022) Izzo, L., Molaro, P., Cescutti, G., et al. 2022, MNRAS, 510, 5302
  * Jeffries et al. (2017) Jeffries, R. D., Jackson, R. J., Franciosini, E., et al. 2017, MNRAS, 464, 1456
  * Kamiński et al. (2015a) Kamiński, T., Mason, E., Tylenda, R., & Schmidt, M. R. 2015a, A&A, 580, A34
  * Kamiński et al. (2015b) Kamiński, T., Menten, K. M., Tylenda, R., et al. 2015b, Nature, 520, 322
  * Kamiński et al. (2017) Kamiński, T., Menten, K. M., Tylenda, R., et al. 2017, A&A, 607, A78
  * Kamiński et al. (2020) Kamiński, T., Menten, K. M., Tylenda, R., et al. 2020, A&A, 644, A59
  * Kamiński et al. (2010) Kamiński, T., Schmidt, M., & Tylenda, R. 2010, A&A, 522, A75
  * Kamiński et al. (2009) Kamiński, T., Schmidt, M., Tylenda, R., Konacki, M., & Gromadzki, M. 2009, ApJS, 182, 33
  * Kamiński et al. (2021a) Kamiński, T., Steffen, W., Bujarrabal, V., et al. 2021a, A&A, 646, A1
  * Kamiński et al. (2018a) Kamiński, T., Steffen, W., Tylenda, R., et al. 2018a, A&A, 617, A129
  * Kamiński & Tylenda (2013) Kamiński, T. & Tylenda, R. 2013, A&A, 558, A82
  * Kamiński et al. (2011) Kamiński, T., Tylenda, R., & Deguchi, S. 2011, A&A, 529, A48
  * Kamiński et al. (2021b) Kamiński, T., Tylenda, R., Kiljan, A., et al. 2021b, A&A, 655, A32
  * Kamiński et al. (2018b) Kamiński, T., Tylenda, R., Menten, K. M., et al. 2018b, Nature Astronomy, 2, 778
  * Kimeswenger (2005) Kimeswenger, S. 2005, Astronomische Nachrichten, 326, 657
  * Kimeswenger (2006) Kimeswenger, S. 2006, Astronomische Nachrichten, 327, 44
  * Kipper et al. (2004) Kipper, T., Klochkova, V. G., Annuk, K., et al. 2004, A&A, 416, 1107
  * Kipper & Škoda (2007) Kipper, T. & Škoda, P. 2007, in Astronomical Society of the Pacific Conference Series, Vol. 363, The Nature of V838 Mon and its Light Echo, ed. R. L. M. Corradi & U. Munari, 231
  * Kochanek et al. (2014) Kochanek, C. S., Adams, S. M., & Belczynski, K. 2014, MNRAS, 443, 1319
  * Kurtenkov et al. (2015) Kurtenkov, A. A., Pessev, P., Tomov, T., et al. 2015, A&A, 578, L10
  * Lindegren et al. (2021) Lindegren, L., Klioner, S. A., Hernández, J., et al. 2021, A&A, 649, A2
  * Longland et al. (2012) Longland, R., Lorén-Aguilar, P., José, J., García-Berro, E., & Althaus, L. G. 2012, A&A, 542, A117
  * Maciel & Costa (2012) Maciel, W. J. & Costa, R. D. D. 2012, Memorie della Societa Astronomica Italiana Supplementi, 22, 103
  * MacLeod et al. (2022) MacLeod, M., De, K., & Loeb, A. 2022, arXiv e-prints, arXiv:2205.07929
  * MacLeod & Loeb (2020) MacLeod, M. & Loeb, A. 2020, ApJ, 895, 29
  * MacLeod et al. (2017) MacLeod, M., Macias, P., Ramirez-Ruiz, E., et al. 2017, ApJ, 835, 282
  * MacLeod et al. (2018a) MacLeod, M., Ostriker, E. C., & Stone, J. M. 2018a, ApJ, 868, 136
  * MacLeod et al. (2018b) MacLeod, M., Ostriker, E. C., & Stone, J. M. 2018b, ApJ, 863, 5
  * Magrini et al. (2021) Magrini, L., Smiljanic, R., Lagarde, N., et al. 2021, The Messenger, 185, 18
  * Martini et al. (1999) Martini, P., Wagner, R. M., Tomaney, A., et al. 1999, AJ, 118, 1034
  * Mason et al. (2010) Mason, E., Diaz, M., Williams, R. E., Preston, G., & Bensby, T. 2010, A&A, 516, A108
  * Mason & Shore (2022) Mason, E. & Shore, S. N. 2022, arXiv e-prints, arXiv:2206.03191
  * Mauron & Huggins (2010) Mauron, N. & Huggins, P. J. 2010, A&A, 513, A31
  * Melis (2020) Melis, C. 2020, Research Notes of the AAS, 4, 238
  * Metzger & Pejcha (2017) Metzger, B. D. & Pejcha, O. 2017, MNRAS, 471, 3200
  * Muller et al. (2020) Muller, S., Roueff, E., Black, J. H., et al. 2020, A&A, 637, A7
  * Munari et al. (2002) Munari, U., Henden, A., Kiyota, S., et al. 2002, A&A, 389, L51
  * Nandez et al. (2014) Nandez, J. L. A., Ivanova, N., & Lombardi, J. C., J. 2014, ApJ, 786, 39
  * Ortiz-León et al. (2020) Ortiz-León, G. N., Menten, K. M., Kamiński, T., et al. 2020, A&A, 638, A17
  * Pastorello et al. (2019a) Pastorello, A., Chen, T. W., Cai, Y. Z., et al. 2019a, A&A, 625, L8
  * Pastorello et al. (2021a) Pastorello, A., Fraser, M., Valerin, G., et al. 2021a, A&A, 646, A119
  * Pastorello et al. (2019b) Pastorello, A., Mason, E., Taubenberger, S., et al. 2019b, A&A, 630, A75
  * Pastorello et al. (2021b) Pastorello, A., Valerin, G., Fraser, M., et al. 2021b, A&A, 647, A93
  * Pejcha (2014) Pejcha, O. 2014, ApJ, 788, 22
  * Pejcha et al. (2016a) Pejcha, O., Metzger, B. D., & Tomida, K. 2016a, MNRAS, 461, 2527
  * Pejcha et al. (2016b) Pejcha, O., Metzger, B. D., & Tomida, K. 2016b, MNRAS, 455, 4351
  * Pejcha et al. (2017) Pejcha, O., Metzger, B. D., Tyles, J. G., & Tomida, K. 2017, ApJ, 850, 59
  * Randich et al. (1999) Randich, S., Gratton, R., Pallavicini, R., Pasquini, L., & Carretta, E. 1999, A&A, 348, 487
  * Rau et al. (2007) Rau, A., Kulkarni, S. R., Ofek, E. O., & Yan, L. 2007, ApJ, 659, 1536
  * Reichardt et al. (2019) Reichardt, T. A., De Marco, O., Iaconi, R., Tout, C. A., & Price, D. J. 2019, MNRAS, 484, 631
  * Reyniers & Van Winckel (2001) Reyniers, M. & Van Winckel, H. 2001, A&A, 365, 465
  * Sackmann & Boothroyd (1992) Sackmann, I. J. & Boothroyd, A. I. 1992, ApJ, 392, L71
  * Schneider et al. (2019) Schneider, F. R. N., Ohlmann, S. T., Podsiadlowski, P., et al. 2019, Nature, 574, 211
  * Shara et al. (1985) Shara, M. M., Moffat, A. F. J., & Webbink, R. F. 1985, ApJ, 294, 271
  * Sills et al. (2001) Sills, A., Faber, J. A., Lombardi, James C., J., Rasio, F. A., & Warren, A. R. 2001, ApJ, 548, 323
  * Smiljanic et al. (2018) Smiljanic, R., Franciosini, E., Bragaglia, A., et al. 2018, A&A, 617, A4
  * Smith et al. (2016) Smith, N., Andrews, J. E., Van Dyk, S. D., et al. 2016, MNRAS, 458, 950
  * Soker (2016) Soker, N. 2016, New A, 47, 16
  * Soker (2019) Soker, N. 2019, MNRAS, 483, 5020
  * Soker (2020a) Soker, N. 2020a, ApJ, 893, 20
  * Soker (2020b) Soker, N. 2020b, Galaxies, 8, 26
  * Soker (2021) Soker, N. 2021, Research in Astronomy and Astrophysics, 21, 112
  * Soker & Kaplan (2021) Soker, N. & Kaplan, N. 2021, Research in Astronomy and Astrophysics, 21, 090
  * Soker & Kashi (2012) Soker, N. & Kashi, A. 2012, ApJ, 746, 100
  * Soker & Tylenda (2007) Soker, N. & Tylenda, R. 2007, MNRAS, 375, 909
  * Stȩpień (2011) Stȩpień, K. 2011, A&A, 531, A18
  * Tylenda et al. (2005a) Tylenda, R., Crause, L. A., Górny, S. K., & Schmidt, M. R. 2005a, A&A, 439, 651
  * Tylenda et al. (2015) Tylenda, R., Górny, S. K., Kamiński, T., & Schmidt, M. 2015, A&A, 578, A75
  * Tylenda et al. (2011a) Tylenda, R., Hajduk, M., Kamiński, T., et al. 2011a, A&A, 528, A114
  * Tylenda & Kamiński (2012) Tylenda, R. & Kamiński, T. 2012, A&A, 548, A23
  * Tylenda et al. (2019) Tylenda, R., Kamiński, T., & Mehner, A. 2019, A&A, 628, A124
  * Tylenda et al. (2011b) Tylenda, R., Kamiński, T., Schmidt, M., Kurtev, R., & Tomov, T. 2011b, A&A, 532, A138
  * Tylenda et al. (2013) Tylenda, R., Kamiński, T., Udalski, A., et al. 2013, A&A, 555, A16
  * Tylenda & Soker (2006) Tylenda, R. & Soker, N. 2006, A&A, 451, 223
  * Tylenda et al. (2005b) Tylenda, R., Soker, N., & Szczerba, R. 2005b, A&A, 441, 1099
  * Wallerstein & Conti (1969) Wallerstein, G. & Conti, P. S. 1969, ARA&A, 7, 99
  * White & Hillenbrand (2005) White, R. J. & Hillenbrand, L. A. 2005, ApJ, 621, L65
  * Whittet (1992) Whittet, D. C. B. 1992, Dust in the galactic environment (CRC Press)
  * Wiese & Fuhr (2009) Wiese, W. & Fuhr, J. 2009, Journal of physical and chemical reference data, 38, 565
  * Williams et al. (2015) Williams, S. C., Darnley, M. J., Bode, M. F., & Steele, I. A. 2015, ApJ, 805, L18
  * Williams et al. (2020) Williams, S. C., Jones, D., Pessev, P., et al. 2020, A&A, 637, A20
  * Woitke et al. (2018) Woitke, P., Helling, C., Hunter, G. H., et al. 2018, A&A, 614, A1
  * Yee & Jensen (2010) Yee, J. C. & Jensen, E. L. N. 2010, ApJ, 711, 303
  * Zhang & Jeffery (2013) Zhang, X. & Jeffery, C. S. 2013, MNRAS, 430, 2113
  * Zhang et al. (2020) Zhang, X., Jeffery, C. S., Li, Y., & Bi, S. 2020, ApJ, 889, 33

##  Appendix A Spectral fits

Figure 6: A sample correction of the spectrum near lithium lines by a
simulation of TiO bands. The observed spectrum (gray) was obtained with SALT
in 2020. The simulation (magenta) included TiO opacity of a cool (800 K) slab
of circumstellar gas. Aside from Li I absorption components, most of the
spectral features are TiO rotational lines in the $P_{2}$ ($J_{\rm
low}$=24–27) and $Q_{2}$ ($J_{\rm low}$=36–40) branches. Weak absorption from
a couple of Fe I lines may be present as well. Figure 7: Best fit synthetic
spectrum (orange) attempting to reproduce the observed spectrum (blue) of the
field star observed behind the outflow of CK Vul. The model spectrum was
reddened by $A_{V}$=6 mag. Main features are identified.