# The Pristine Inner Galaxy Survey (PIGS) VII: a discovery of the first inner
Galaxy CEMP-r/s star††thanks: Based on observations made with the Very Large
Telescope (VLT)

L. Mashonkina,1 A. Arentsen,2 D. S. Aguado,3,4 A. Smogorzhevskii,1,5 M.
Hampel,6 A.I Karakas,6,7 F. Sestito,8 N. F. Martin,9,10 K. A. Venn,8 J. I.
González Hernández3,4
1 Institute of Astronomy of the Russian Academy of Sciences, Pyatnitskaya st.
48, 119017, Moscow, Russia
2 Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge
CB3 0HA, UK
3Instituto de Astrofísica de Canarias, Vía Láctea, 38205 La Laguna, Tenerife,
Spain
4Universidad de La Laguna, Departamento de Astrofísica, 38206 La Laguna,
Tenerife, Spain
5 M. V. Lomonosov Moscow State University, Kolmogorova st. 1, 119991, Moscow,
Russia
6 School of Physics & Astronomy, Monash University, Clayton, VIC 3800,
Australia
7 ARC Centre of Excellence for All Sky Astrophysics in 3 Dimensions (ASTRO
3D), Melbourne, VIC, Australia
8 Department of Physics and Astronomy, University of Victoria, PO Box 3055,
STN CSC, Victoria BC V8W 3P6, Canada
9 Observatoire Astronomique de Strasbourg, Université de Strasbourg, CNRS, UMR
7550, F-67000 Strasbourg, France
10 Max-Planck-Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg,
Germany E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

Well-studied very metal-poor (VMP, [Fe/H] $<-2$) stars in the inner Galaxy are
few in number, and they are of special interest because they are expected to
be among the oldest stars in the Milky Way. We present high-resolution
spectroscopic follow-up of the carbon-enhanced metal-poor (CEMP) star
Pristine_184237.56-260624.5 (hereafter Pr184237) identified in the Pristine
Inner Galaxy Survey. This star has an apocentre of $\sim 2.6$ kpc. Its
atmospheric parameters ($T_{\rm eff}$ = 5100 K, $\log{g}$ = 2.0, [Fe/H] =
$-2.60$) were derived based on the non-local thermodynamic equilibrium (NLTE)
line formation. We determined abundances for 32 elements, including 15 heavy
elements beyond the iron group. The NLTE abundances were calculated for 13
elements from Na to Pb. Pr184237 is strongly enhanced in C, N, O, and both s-
and r-process elements from Ba to Pb; it reveals a low carbon isotope ratio of
12C/13C = 7. The element abundance pattern in the Na-Zn range is typical of
halo stars. With [Ba/Eu] = 0.32, Pr184237 is the first star of the CEMP-r/s
subclass identified in the inner Galaxy. Variations in radial velocity suggest
binarity. We tested whether a pollution by the s- or i-process material
produced in the more massive and evolved companion can form the observed
abundance pattern and find that an i-process in the asymptotic giant branch
star with a progenitor mass of 1.0-2.0$M_{\odot}$ can be the solution.

###### keywords:

stars: abundances – stars: atmospheres – galaxies: abundances.

††pubyear: 2023††pagerange: The Pristine Inner Galaxy Survey (PIGS) VII: a
discovery of the first inner Galaxy CEMP-r/s star††thanks: Based on
observations made with the Very Large Telescope (VLT)–References

##  1 Introduction

Studies of very metal-poor (VMP, [Fe/H]111In the classical notation, where
[X/Y] = $\log(N_{\rm X}/N_{\rm Y})_{star}-\log(N_{\rm X}/N_{\rm Y})_{\odot}$
for each pair of elements X and Y. $<-2.0$) stellar populations in the Milky
Way are important to understand the early Universe. Their detailed chemical
abundances teach us about the properties of the First Stars and early star
formation, and their chemo-dynamics shed light on the early formation history
of our Galaxy (Frebel & Norris, 2015). The central regions of our Galaxy
($\lesssim 5$ kpc) are predicted to host the oldest metal-poor stars
(Tumlinson, 2010), which are important probes of the earliest metal-free stars
in the Universe. In recent years, significant efforts have been made to build
larger samples of VMP stars in the inner Galaxy (Howes et al., 2015, 2016;
Lucey et al., 2019; Arentsen et al., 2020b) – a challenging endeavour because
the overwhelming majority of bulge stars have high metallicity, and the high
extinction and the relatively large distance from the Sun make clean
selections of metal-poor stars difficult. These surveys typically employ
narrow-band photometry to identify metal-poor stars, which has been very
effective.

Recent studies show that VMP stars in the inner Galaxy appear to have similar
abundances to stars in previous halo samples, although there are some subtle
differences in, e.g., the scatter and correlations of various abundances
(Howes et al., 2016; Koch et al., 2016; Lucey et al., 2019; Sestito et al.,
2023). One other striking difference is an apparent lack of carbon-enhanced
metal-poor (CEMP) stars (Howes et al., 2016; Arentsen et al., 2021), with the
common definition for CEMP stars having [C/Fe] $>+0.7$ (Aoki et al., 2007). In
the Galactic halo, many of the most metal-poor stars are found to be rich in
carbon, representing 20% and 43% of stars with [Fe/H] $\leq-2$ and $\leq-3$,
respectively (Placco et al., 2014). For the bulge, Arentsen et al. (2021)
found a CEMP fraction of only 16% for [Fe/H] $<-2.5$ in the low-resolution
spectroscopic Pristine Inner Galaxy Survey (PIGS). They discuss that it might
partly be related to selection effects against carbon-rich stars, but argue
that it cannot be explained away entirely. The lower CEMP fraction could be
due to a lower number of binary stars in the inner Galaxy, and/or due to
faster and more intense early star formation – both scenarios would be
indications that the metal-poor stellar populations in the inner Galaxy and
the more distant halo formed in different environments.

Depending on abundances of the elements produced presumably in the slow (s) or
rapid (r) neutron-capture processes, which can be derived from high-resolution
spectroscopy, CEMP stars are separated into CEMP-r ([Eu/Fe] $>1$), CEMP-s
([Ba/Fe] $>1$ and [Ba/Eu] $>0.5$), CEMP-r/s (enhanced in both barium and
europium with $0<$ [Ba/Eu] $<0.5$), and CEMP-no ([Ba/Fe] $<0$) stars, as
suggested by Beers & Christlieb (2005). Slightly different numbers are
recommended by Abate et al. (2016) and Frebel (2018) to characterise different
groups. Being purely phenomenological, this separation turns out to be
indicative of the nature of stars in different groups. Excess of carbon and
s-process elements in a CEMP-s star is thought to be the result of mass
transfer from an asymptotic giant branch (AGB) star to a companion in a binary
system. This hypothesis is supported by radial velocity variations observed
for many CEMP-s stars (Hansen et al., 2016). The CEMP-no stars, which are
typically not in binary systems, likely formed out of the interstellar matter
enriched by the Population III stars, which either were rapidly rotating
(Chiappini et al., 2006) or exploded as faint supernovae (Nomoto et al., 2013,
and references therein). CEMP-r stars are similar to CEMP-no stars, except
that they are enhanced in r-process elements, which they were likely born with
as well. The nature of CEMP-r/s stars is under debate, and various scenarios
of enrichment in heavy elements are considered, from mass transfer in a binary
system formed out of the interstellar matter rich in the r-process elements to
the operation of the intermediate neutron-capture process (i-process, see
Goswami et al., 2021, for a review).

Detailed studies of CEMP stars can provide a unique information about the
early chemical enrichment of the inner Galaxy, but they are only available in
the literature for a very few stars. Four CEMP-no stars, all with [Fe/H] $<-3$
and low Ba abundances, were discovered in the EMBLA (Extremely Metal-poor
BuLge stars with AAOmega) sample (Howes et al., 2015, 2016): one was reported
by Howes et al. (2015), and three additional stars were found by Arentsen et
al. (2021) when applying the evolutionary carbon corrections from Placco et
al. (2014). The three CEMP stars with close together metallicities of [Fe/H]
$\simeq-2.5$ were found to be strongly enhanced in Ba, with [Ba/Fe] $>+1$. One
of them was classified by Koch et al. (2016) as CEMP-s based on low upper
limits for abundances of the r-process elements Eu and Dy. None of the
r-process elements were measured by Sestito et al. (2023) in the remaining two
CEMP stars, due to a limited wavelength range of their Gemini GRACES spectra.

In this work, we report the detailed analysis of a new CEMP star in the inner
Galaxy, Pristine_184237.56-260624.5 (for short, Pr184237), with [Fe/H] =
$-2.6$ and [C/Fe] = $+1.8$, which is enhanced in both s- and r-process
elements. It was observed as part of a high-resolution spectroscopic follow-up
campaign of the most metal-poor stars in PIGS, undertaken with
UVES222Ultraviolet and Visual Echelle Spectrograph at the VLT. It deserved a
careful, dedicated analysis due to the high carbon and neutron-capture element
abundances, which we report on here separately from the main sample. We
performed a detailed line-by-line analysis of the spectrum of Pr184237 in a
wide wavelength range and determined abundances for 32 chemical elements,
including 15 neutron-capture elements. Abundances for 13 chemical elements are
based on the non-local thermodynamic equilibrium (NLTE) line formation, which
is important because the physical conditions in atmospheres of VMP giants are
favorable for the departures from local thermodynamic equilibrium (LTE).

The paper is organised as follows. We introduce PIGS and the high-resolution
observations of Pr184237 in Sect. 2. Section 3 describes the NLTE methods and
codes used. Section 4 discusses the determination of the star’s atmospheric
parameters. Elemental abundances are derived in Sect. 5, and Sect. 6 discusses
the possible origins of the heavy element abundance pattern of Pr184237. Our
conclusions are given in Sect. 7.

##  2 Observations of Pr184237

The Pristine Inner Galaxy Survey (PIGS), an extension of the main Pristine
survey (see Starkenburg et al., 2017, for an overview), was initiated to build
an unprecedentedly large sample of metal-poor and very metal-poor stars for
chemo-dynamical studies of the inner regions of the Galaxy (Arentsen et al.,
2020a, b, 2021). Using metallicity-sensitive $CaHK$ photometry from the
Canada-France-Hawaii-Telescope (CFHT), PIGS selects metal-poor candidates for
spectroscopic follow-up. The next step is an intermediate-resolution
spectroscopy with the AAOmega spectrograph (Saunders et al., 2004) on the
Anglo-Australian Telescope combined with determinations of the stellar
atmosphere parameters using the ULySS (Koleva et al., 2009) and FERRE (Allende
Prieto et al., 2006) tools. A total of 1900 VMP stars were identified in the
PIGS-AAT sample. Aiming to understand the nucleosynthesis processes in the
ancient inner Galaxy, $\sim$20 of the most metal-poor stars were selected from
the AAT sample for high-resolution spectroscopy with the UVES/VLT, one of
which was Pr184237.

Pr184237 was found to have [Fe/H] = $-2.8$ and [C/Fe] = $+2.0$ from the low-
resolution spectroscopy and was recognised as a CEMP star in Arentsen et al.
(2021). The AAT barycentric radial velocity is $-12.5\pm 2.0$ km s-1. It is
located at (RA, Dec) = (18:42:37.56,$-$26:06:24.5) and
$(l,b)=(8.514105^{\circ},-9.783936^{\circ})$ and identified with Gaia DR3
source_id 4073253907337129472 of G magnitude = 15.52 (Gaia Collaboration et
al., 2022). It has E(B-V) = 0.38 (Green et al., 2019).

###  2.1 UVES observations

The UVES observations of Pr184237 were originally scheduled for period 105 in
queue mode. However, due to COVID restrictions in Paranal (Chile),
observations were taken in period 107. The employed setup was
$\\#Dic1\,(390+580)$ with a $1\aas@@fstack{\prime\prime}2$ slit, $1\times 1$
binning and low readout speed. This configuration led to a resolving power of
R$\sim 45,000$ in the blue part of the spectrum ($330-452$ nm) and R$\sim
41,500$ for the red ($480-680$ nm). Three observations were taken between the
3rd and the 4th of June 2021 with an exposure time of 3000 s each and an
average signal-to-noise ratio (SNR) of 7 and 32 at 380 and 510 nm,
respectively.

The data were reduced by the ESO pipeline and retrieved by a query on the
phase 3 online interface333https://www.eso.org/sci/observing/phase3.html. The
barycentric correction was applied with IRAF and the radial velocity was
calculated with a cross correlation function to a UVES template with the fxcor
package, after which the three spectra were combined in the rest frame with a
“sigclip” algorithm. The three radial velocities are reported in Table 1,
together with the AAT radial velocity.

A first full-spectrum fitting analysis of the UVES spectrum was performed with
the FERRE444http://github.com/callendeprieto/ferre code following the same
procedure explained in Aguado et al. (2021a, b). In brief, the code is able to
fit the data and interpolate between the nodes of a grid of stellar models
computed with ASSET (Koesterke et al., 2008) and ATLAS (Sbordone, 2005) codes
of spectral synthesis. These models already contain carbon-enrichment, an
important feature when analyzing metal-poor stars (see e.g. Aguado et al.,
2017, for more information about the carbon models). The four parameters FERRE
derived for Pr184237 are effective temperature ($T_{\rm eff}=5448\pm 123$ K),
surface gravity ($\log{g}=2.01\pm 0.26$), overall metallicity
($\rm{[Fe/H]}=-2.46\pm 0.10$), and carbon enrichment ($\rm{[C/Fe]}=2.19\pm
0.12$). A preliminary automatic abundance analysis with FERRE indicated that
Pr184237 is enhanced in both s-process ([Ba/Fe] = 1.04, [La/Fe] = 0.56) and
r-process ([Eu/Fe] = 0.6) elements, suggesting that this is the first CEMP-r/s
star in the inner Galaxy. This encouraged us to do a dedicated detailed
analysis (see Sect. 4-5). The FERRE analysis for the full PIGS-UVES sample
will be presented in a forthcoming paper (Aguado et al., in prep.).

Table 1:  Radial velocities of Pr184237.  Instrument | MJD | RV
---|---|---
| [days] | [km s-1]
AAT | 58337.44 | -12.5 $\pm$ 2.0
UVES | 59368.37 | -39.3 $\pm$ 1.0
UVES | 59369.33 | -39.1 $\pm$ 1.0
UVES | 59369.36 | -39.8 $\pm$ 1.0

###  2.2 Variability and binarity

The AAT baricentric radial velocity ($-12.5\pm 2.0$ km s-1) is significantly
different from the three UVES radial velocities taken almost three years later
($-39.4$ km s-1 with a standard deviation of 0.3 km s-1). The precision of the
AAT radial velocities was estimated by Arentsen et al. (2020b) to be around 2
km s-1. A more recent comparison with Gaia DR3 radial velocities for the full
AAT sample (where available, only for brighter stars and not for Pr184237)
shows that the uncertainty estimate was correct and possibly even slightly
over-estimated (Arentsen et al. in prep). Therefore the difference between the
AAT and UVES radial velocities can be taken as a clear indication that this
star is in a binary (or multiple) system.

We also found that Pr184237 was flagged as a photometrically variable star in
Gaia DR3, with the classification "ELL" (Rimoldini et al., 2022, best class
score = 0.58) – an ellipsoidal variable, which is a close binary system where
the stars are deforming and projection effects during the orbit cause
photometric variability. There is a Gaia light curve available for this star,
with variations on the order of $\sim 0.03$ mag on a short timescale (less
than a few days). Fitting the epoch photometry with lightkurve (Lightkurve
Collaboration et al., 2018) we found a best period of 1.8 days, but it was not
very significant. Such a short period (if this star really was an ellipsoidal
variable in a binary system) is not consistent with our UVES observations. The
UVES radial velocities taken over two days do not show any significant radial
velocity variation, and there is also no sign of severe rotational line
broadening. One could maybe imagine a scenario where this star is in a triple
system, orbiting in a longer period around an inner short-period binary. In
this scenario the short-term variability could come from the inner binary. If
the close binary consists of two white dwarfs, they should be too faint to
contribute to the photometric variability since a white dwarf is $\sim 10$ mag
fainter than an RGB star like Pr184237 (related to the rapidly accreting white
dwarf scenario discussed in Section 6.2).

However, inspecting images of Pr184237, we found that there is a faint star
($\sim 4$ mag fainter, G = 19.9) within $\sim 2^{\prime\prime}$ of the main
target. It has different proper motions so is not associated with Pr184237,
but it might affect the photometry. We checked the Gaia scanning law and found
that, for the epochs where Pr184237 appears brighter, Gaia scanned the sky
exactly in the direction in which the two stars are aligned, and, when
Pr184237 appears fainter, it was scanned perpendicular to that axis. The
difference in flux between the two stars is similar to the amplitude of the
photometric variability. Combined with the lack of evidence for a short
orbital period from our UVES observations, we therefore concluded that the
Gaia photometric variability is most likely spurious and not related to
Pr184237 itself.

###  2.3 Orbit in the Milky Way

Stars that are currently in the inner Galaxy do not necessarily stay in the
inner Galaxy – they could be outer halo stars just moving through. To test
this for Pr184237, we derived its distance and orbital properties. The
distance and orbits were derived adopting the same methodology used by Sestito
et al. (2023) for the PIGS follow-up with the GRACES high-resolution
spectrograph. Briefly, the Bayesian distance derivation uses the Gaia DR3
parallax and its uncertainty plus a prior on the stellar density distribution
containing a disc and a halo (see Sestito et al. 2019 for more details about
the method). For Pr184237, it resulted in a median distance of $7.97\pm 2.57$
kpc (corresponding to a distance from the Galactic centre
$d_{\mathrm{GC}}\approx 1.8$ kpc). The orbital integration is performed with
GALPY (Bovy, 2015) with the inclusion of a rotating bar in the gravitational
potential. Then the orbital parameters are calculated from the distance
probability distribution function. More details can be found in Sestito et al.
(2023). We adopt a radial velocity of $-26.0$ km s-1, the average between the
AAT and the UVES epochs, with an inflated uncertainty of 10 km s-1 to
represent that we do not actually know the systemic velocity. Adopting the
above distance and radial velocity with their uncertainties and combining it
with the Gaia DR3 proper motions, we derive the orbital properties of
Pr184237, and find it has a median apocentre of $2.56\pm 0.93$ kpc with a
pericentre less than 1 kpc – it appears to be a real inner Galaxy star and not
just a halo star passing through. Given that it is so tightly confined, the
star is most likely to be of a “proto-Galactic” origin (Rix et al., 2022;
Belokurov & Kravtsov, 2022) rather than accreted from a dwarf galaxy, or if it
was accreted, this must have happened very early on.

##  3 Analysis method

We determined the NLTE abundances for Na, Mg, Al, Si, Ca, Sc, Ti, Fe, Zn, Sr,
Ba, Eu, and Pb. Hereafter, these elements are referred to as NLTE species. For
individual lines of Ti ii, Fe i, and Zn i, the NLTE abundances were computed
by adding the NLTE abundance corrections, $\Delta_{\rm
NLTE}=\log\varepsilon_{\rm\rm NLTE}-\log\varepsilon_{\rm\rm LTE}$, to the LTE
abundances. The $\Delta_{\rm NLTE}$ grids and interpolation tools are provided
by Mashonkina et al. (2016b, Ti ii, Fe i) and Sitnova et al. (2022, Zn i). For
the remaining NLTE species, we performed the NLTE calculations using
individual atmospheric parameters and elemental abundances of Pr184237 and the
NLTE methods developed in our previous studies for Na i (Alexeeva et al.,
2014), Mg i (Mashonkina, 2013), Al i (Mashonkina et al., 2016a), Si i-Si ii
(Mashonkina, 2020), Ca i-Ca ii (Mashonkina et al., 2017b), Sc ii (Mashonkina &
Romanovskaya, 2022), Sr ii (Mashonkina et al., 2022), Ba ii (Mashonkina &
Belyaev, 2019), Eu ii (Mashonkina & Gehren, 2000), and Pb i (Mashonkina et
al., 2012).

The coupled radiative transfer and statistical equilibrium (SE) equations were
solved with the code detail (Butler, 1984; Giddings, 1981), where the
background opacity was updated as described by Mashonkina et al. (2011). The
computed departure coefficients, $b_{i}=n_{i}^{\rm NLTE}/n_{i}^{\rm LTE}$,
were then used by the code synthV_NLTE (Tsymbal et al., 2019) to calculate the
NLTE line profiles for a given NLTE species. Here, $n_{i}^{\rm NLTE}$ and
$n_{i}^{\rm LTE}$ are the SE and thermal (Saha-Boltzmann) number densities,
respectively.

The LTE and NLTE abundances were derived from line profile fitting using the
code synthV_NLTE (Tsymbal et al., 2019). The best fit to the observed spectrum
is obtained automatically using the IDL binmag code by Kochukhov (2018). The
line list required for calculations of the synthetic spectra, together with
atomic data, were taken from the most recent version of Vienna Atomic Line
Database (VALD, Ryabchikova et al., 2015) that includes isotopic splitting
(IS) and hyper-fine splitting (HFS) components where they are available
(Pakhomov et al., 2019). The lines of the NLTE species were computed by
implementing the departure coefficients from detail, while LTE was assumed for
the lines of the other species.

We used the homogeneous spherical model atmospheres with standard abundances
(Gustafsson et al., 2008) from the MARCS website555https://marcs.astro.uu.se.
They were interpolated at the given effective temperature ($T_{\rm eff}$),
surface gravity ($\log{g}$), and metallicity ([Fe/H]), applying the
interpolation routine written by Thomas Masseron and available on the same
website.

##  4 Atmospheric parameters

For Pr184237, the preliminary automatic FERRE analysis (Aguado et al., in
prep.) yielded $T_{\rm eff}$ = 5450$\pm 123$ K, $\log{g}$ = 2.01$\pm 0.26$,
and [Fe/H] = $-2.46\pm 0.10$. In view of the extreme importance of derived
elemental abundances for understanding the phenomenon of a CEMP-r/s star and
for learning mechanisms of the early chemical enrichment of the inner Galaxy,
we double checked and improved the atmospheric parameters of this star.

The effective temperature is derived by fitting the theoretical profiles
computed by Amarsi et al. (2018) with three-dimensional (3D) hydrodynamic
model atmospheres and based on the NLTE line formation to the observed
wavelength regions free of blending lines in the wings of the Balmer Hα and Hβ
lines (Fig. 1). For both lines, the theoretical profiles corresponding to
$T_{\rm eff}$ = 5450 K are apparently deeper compared with the observed ones.
Using a $\chi^{2}$ minimization, we find that the best fit to Hα is achieved
with $T_{\rm eff}$ = 5100 K ($\chi^{2}$ = 0.127). For Hβ, equal preference can
be given to $T_{\rm eff}$ = 5100 K and 4900 K due to similar $\chi^{2}\simeq$
0.4. We give the greater weight to $T_{\rm eff}$ from Hα because of less
important blending effects and higher S/N, and adopt $T_{\rm eff}$ = 5100 K as
our final effective temperature. From a comparison of the results from two
lines, we estimate the uncertainty in $T_{\rm eff}$ as 100 K.

Figure 1:  Hα (top panel) and Hβ (bottom panel) line profiles in Pr184237
(bold dots) compared with the theoretical 3D-NLTE profiles from Amarsi et al.
(2018) for $T_{\rm eff}$ = 5100 K (red continuous curve), $T_{\rm eff}$ = 5450
K (blue dashed curve), and $T_{\rm eff}$ = 4900 K (green long-dashed curve,
for Hβ only). Everywhere, $\log{g}$ = 2.0 and [Fe/H] = $-2.5$. The yellow
shaded regions show the spectral windows used for the fit.

The surface gravity is constrained from a requirement of equal NLTE abundances
determined from lines of the two ionization stages, Fe i and Fe ii. We use the
Fe i and Fe ii lines and their atomic parameters, carefully selected and
checked in our previous studies of VMP stars (Mashonkina et al., 2017a). In
brief, for Fe i, we use presumably experimental $gf$-values, as listed by
VALD, while, for Fe ii, we apply $gf$-values from Raassen & Uylings (1998)
that were corrected by $+0.11$ dex, following the recommendation of Grevesse &
Sauval (1999). The exceptions are Fe ii 4923 and 5018 Å, for which $gf$-values
were obtained by averaging the data from the four sources – Bridges (1973);
Moity (1983), Raassen & Uylings (1998), and Meléndez & Barbuy (2009). The van
der Waals damping constants were taken from VALD.

Table 2:  Line atomic data, LTE, and NLTE abundances, $\log\varepsilon$, for
individual lines in Pr184237.

Species | $\lambda$ | $E_{\rm exc}$ | log $gf$ | LTE | NLTE
---|---|---|---|---|---
| [Å] | [eV] |  |  |
12CH | 4210.94 | 0.46 | -1.34 | 7.491 |
12CH | 4210.99 | 0.46 | -1.32 | 7.491 |
13CH | 4211.44 | 0.46 | -1.34 | 7.491 |
13CH | 4211.49 | 0.46 | -1.32 | 7.491 |
O i | 6300.30 | 0.00 | -9.94 | 7.45 |
Na i | 5889.95 | 0.00 | 0.12 |  $\leq$3.94 |  $\leq$3.37
Mg i | 5528.41 | 4.33 | -0.50 | 5.36 | 5.40

Notes. 1 For the 12C/13C = 7 isotope ratio. This table is available in its
entirety in a machine-readable form in the online journal. A portion is shown
here for guidance regarding its form and content.

Figure 2:  NLTE abundances from individual lines of Fe i (open circles) and Fe
ii (filled circles) in Pr184237 as a function of equivalent width EW (top
panel) and excitation energy of the lower level (bottom panel). The dotted
line indicates the mean Fe i-based abundance and the shaded grey area its
statistical error.

We checked the surface gravity $\log{g}$ = 2.0, which is very close to the
FERRE value. The LTE and NLTE abundances are determined from lines of Fe i and
Fe ii using the model atmosphere with $T_{\rm eff}$/$\log{g}$/[Fe/H] =
5100/2.0/$-2.5$ and a microturbulent velocity of $\xi_{t}$ = 2 km s-1, which
was calculated with the empirical formula deduced by Mashonkina et al. (2017a)
for metal-poor Galactic halo giants. It provides an accuracy of 0.2 km s-1.
The results for individual lines are presented in Table 2. The mean Fe i\- and
Fe ii-based NLTE abundances are found to be consistent within 0.02 dex:
$\log\varepsilon_{\rm FeI}$ = 4.86$\pm$0.14 (49 lines) and
$\log\varepsilon_{\rm FeII}$ = 4.84$\pm$0.12 (9 lines), while a difference in
the LTE abundances amounts to $-0.12$ dex. Hereafter, we employ the abundance
scale where $\log\varepsilon_{\rm H}$ = 12. The abundance error is calculated
as the dispersion in the single line measurements around the mean,
$\sigma=\sqrt{\Sigma(\overline{x}-x_{i})^{2}/(N_{l}-1)}$. Here, $N_{l}$ is the
number of lines used. Taking into account that two independent methods provide
consistent results, we estimate the uncertainty in $\log{g}$ as 0.1. Shifts of
+0.1/$-0.1$ in $\log{g}$ result in $\log\varepsilon_{\rm
FeI}-\log\varepsilon_{\rm FeII}=-0.02/+0.06$ dex. Figure 2 shows no trend of
the NLTE abundances with equivalent width, although there are four outliers
that reveal a lower abundance compared with the mean, by more than 0.25 dex.
The final Fe abundance is computed as the mean of $\varepsilon_{\rm FeI}$ and
$\varepsilon_{\rm FeII}$.

Table 3:  Determined stellar parameters of Pr184237. Parameter | Value | Uncertainty
---|---|---
$T_{\rm eff}$ | 5100 K |  $\pm 100$ K
$\log{g}$ | 2.0 | $\pm 0.1$
$\mathrm{[Fe/H]}$ | $-2.60$ | $\pm 0.14$
$\xi_{t}$ | 2.0 km s-1 |  $\pm 0.2$ km s-1

Despite large reddening for the star observed towards the bulge, we have
attempted to derive the effective temperature of our target from the infrared
flux method (IRFM) of González Hernández & Bonifacio (2009), and that attempt
was successful. We used the available photometry in the infrared bands from
the Two Micron All-Sky Survey (2MASS, Skrutskie et al., 2006). The 2MASS JHKs
magnitudes are accurate with J = 13.665$\pm 0.024$, H = 13.114$\pm 0.027$, and
Ks = 12.984$\pm 0.030$. We estimated the Johnson V magnitude (V = 15.942) $\pm
0.012$ converted from the Sloan PanSTARRS magnitudes g = 16.370 $\pm 0.003$
and r = 15.550 $\pm 0.001$ using the equation (6) and coefficients from Table
6 in Tonry et al. (2012). Magnitudes were de-reddened using E(B-V) = 0.38
(Green et al., 2019). The same reddening value of E(B-V)${}_{\rm SFD}=0.373$
was computed from the dust maps of Schlegel et al. (1998, SFD). Finally,
E(B-V)SFD was multiplied by a factor of 0.86 following Schlafly & Finkbeiner
(2011). With E(B-V) = 0.321, we obtain $T_{\rm eff}$(IRFM) = $4999\pm 102$,
$5116\pm 67$, and $5080\pm 86$ K for the 2MASS bands J, H, and Ks,
respectively. The IRFM-based mean temperature of Pr184237, $T_{\rm eff}$(IRFM)
= $5073\pm 83$ K, is consistent with that from the Hα wings. These
calculations were performed with $\log{g}$ = 2 and [Fe/H] = $-2.5$.

We also inferred the effective temperature and surface gravity as in Sestito
et al. (2023). Briefly, they develop a Monte Carlo method that consists in
deriving the effective temperature using the Gaia photometry-temperature
relation from Mucciarelli et al. (2021) and the surface gravity from the
Stephan-Boltzmann equation. This method has been tested to provide compatible
stellar parameters with spectroscopic inferences (e.g. Lardo et al., 2021;
Sestito et al., 2023). We find a photometric $T_{\rm eff}$ = 5190 $\pm$ 100 K
and $\log{g}$ = 2.25 $\pm$ 0.16. The spectroscopic and photometric
temperatures agree within 0.63$\sigma$ and the distance-based surface gravity
for such a distant object and the spectroscopic one are compatible within
1.32$\sigma$.

Thus, we adopt $T_{\rm eff}$ = 5100$\pm 100$ K, $\log{g}$ = 2.0$\pm 0.1$,
[Fe/H] = $-2.60\pm 0.14$, $\xi_{t}$ = 2$\pm 0.2$ km s-1 as the final
atmospheric parameters of Pr184237 (Table 3). In order to compute the
metallicity, we use the solar Fe abundance $\log\varepsilon_{\rm\odot}$ = 7.45
from Lodders (2021).

##  5 Chemical abundances

The strongest lithium line, Li i 6707 Å, cannot be extracted from noise in the
spectrum of Pr184237, suggesting a low Li abundance, as expected for a red
giant (Iben, 1967). In addition to iron, we are able to determine abundances
for 31 chemical elements. Spectral lines for abundance analysis, together with
their atomic data, are selected from the list provided by Mashonkina et al.
(2010) for an r-process enhanced star, HE 2327-5642, with similar atmospheric
parameters. For lines of Sc ii, Mn i, Co i, Sr ii, Ba ii, La ii, Eu ii, and Yb
ii, the HFS and/or IS structure were properly taken into account.

The LTE and NLTE abundances obtained from individual lines are presented in
Table 2. Table 4 lists the elemental average abundances together with their
stochastic errors. Stochastic errors are caused by random uncertainties in the
continuum placement, line profile fitting, and $gf$-values. They are
represented by the dispersion in the measurements $\sigma$ when $N_{l}>1$.
When only a single line or a molecular band is measured, the abundance error
resembles the fit uncertainty provided by MPFIT (Markwardt, 2009) and computed
as $PERROR\cdot\sqrt{\chi^{2}/N_{DOF}}$, where $N_{DOF}$ is the number of
fitting degrees of freedom and $PERROR$ is a formal parameter uncertainty,
computed from the covariance matrix. In order to compute the abundance ratios
relative to iron, [X/Fe], we used the present Solar system abundances
recommended by Lodders (2021) and [Fe/H] = $-2.60$. For Na, Mg, Al, Si, Ca,
Sc, Ti, and Zn, the [X/Fe] ratios are based on the NLTE abundances. For Cr,
Mn, Co, and Ni, Table 4 indicates the LTE abundance ratio [X/Fe i], which is
free, in part, from the departures from LTE. Since the NLTE abundances were
determined for 4 of 15 heavy elements, the [X/Fe] ratios for all the heavy
elements in Table 4 and the heavy element abundance pattern in our further
analysis (Sections 5.4 and 6) are based on the LTE abundances, for
consistency. In Sect 6.3, we discuss the influence of using the NLTE
abundances on our analysis of the nucleosynthesis scenarios.

Table 4:  Elemental abundances of Pr184237.

Species | $\log\varepsilon_{\rm\odot}$ | $\log\varepsilon$ | $N_{l}$ | [X/Fe]
---|---|---|---|---
|  | LTE | NLTE |  |
C(C2) | 8.47 | 7.64(0.03) |  | 1 | 1.77
N(CN) | 7.85 | 6.54(0.10) |  | 1 | 1.29
O i | 8.73 | 7.45(0.03) |  | 1 | 1.32
Na i | 6.27 |  $<$ 3.94 |  $<$ 3.37 | 2 | $<-0.30^{\rm N}$
Mg i | 7.52 | 5.36(0.02) | 5.40(0.02) | 1 | 0.48N
Al i | 6.42 | 3.31(0.15) | 3.76(0.15) | 1 | –0.06N
Si i | 7.51 | 5.30(0.10) | 5.27(0.10) | 1 | 0.36N
Ca i | 6.27 | 4.04(0.11) | 4.14(0.11) | 7 | 0.47N
Sc ii | 3.04 | 0.90(0.05) | 0.90(0.05) | 1 | 0.46N
Ti ii | 4.90 | 2.72(0.12) | 2.72(0.12) | 9 | 0.42N
V ii | 3.95 | 1.28(0.05) |  | 2 | –0.08
Cr i | 5.63 | 2.81(0.06) |  | 3 | –0.09I
Mn i | 5.47 | 2.41(0.10) |  | 1 | –0.33I
Fe i | 7.45 | 4.72(0.13) | 4.86(0.14) | 49 | 0.01N
Fe ii | 7.45 | 4.84(0.12) | 4.84(0.12) | 9 | –0.01N
Co i | 4.86 | 2.44(0.05) |  | 3 | 0.31I
Ni i | 6.20 | 3.40(0.18) |  | 5 | –0.07I
Zn i | 4.61 | 2.05(0.10) | 2.22(0.10) | 1 | 0.21N
Sr ii | 2.88 | 0.27(0.05) | 0.22(0.05) | 2 | –0.02
Y ii | 2.15 | –0.24(0.20) |  | 8 | 0.21
Zr ii | 2.55 | 0.67(0.22) |  | 4 | 0.72
Ba ii | 2.17 | 0.74(0.08) | 0.40(0.10) | 3 | 1.17
La ii | 1.17 | –0.33(0.13) |  | 7 | 1.10
Ce ii | 1.58 | 0.09(0.19) |  | 5 | 1.11
Nd ii | 1.45 | –0.08(0.14) |  | 10 | 1.07
Sm ii | 0.94 | –0.55(0.13) |  | 2 | 1.11
Eu ii | 0.51 | –1.24(0.06) | –1.09(0.06) | 4 | 0.85
Gd ii | 1.05 | –0.77(0.08) |  | 2 | 0.78
Dy ii | 1.12 | –0.47(0.09) |  | 5 | 1.01
Er ii | 0.92 | –0.67(0.07) |  | 3 | 1.01
Tm ii | 0.11 | –1.00(0.05) |  | 1 | 1.49
Yb ii | 0.91 | –0.66(0.32) |  | 1 | 1.03
Pb i | 2.03 | 1.55(0.16) | 2.21(0.16) | 2 | 2.12

Notes. The numbers in parentheses are the abundance errors. N Based on the
NLTE abundance. I LTE abundance ratio [X/Fe i].

We examined the systematic uncertainties in the derived elemental abundances
of Pr184237, which are linked to the adopted stellar parameters. These were
estimated by varying $T_{\rm eff}$ by 100 K, $\log{g}$ by 0.1, and $\xi_{t}$
by $-0.2$ km s-1 in the model atmosphere. Table 5 summarizes the various
sources of uncertainties. The total impact of varying $T_{\rm eff}$,
$\log{g}$, and $\xi_{t}$ is computed as the quadratic sum of Cols. 2-4 and
denoted $\Delta(T,g,\xi_{t})$. Column 6 lists the stochastic errors
$\sigma_{\rm obs}$, as given in Table 4. The total uncertainty $\sigma_{\rm
tot}$ in the absolute abundance of each element is computed by the quadratic
sum of the stochastic and systematic (Col. 5) errors.

Table 5:  Error budget for elemental abundances in Pr184237. El. | $\Delta T$ | $\Delta\log{g}$ | $\Delta\xi_{t}$ | $\Delta$ | $\sigma_{\rm obs}$ | $\sigma_{\rm tot}$
---|---|---|---|---|---|---
| 100 K | 0.1 | –0.2 km s-1 | ($T,g,\xi_{t}$) |  |
(1) | (2) | (3) | (4) | (5) | (6) | (7)
C2 | 0.14 | -0.02 |  $<$0.01 | 0.14 | 0.03 | 0.14
CN | 0.17 | -0.02 |  $<$0.01 | 0.17 | 0.10 | 0.20
O i | 0.07 | 0.03 |  $<$0.01 | 0.08 | 0.03 | 0.08
Mg i | 0.06 |  $<$0.01 | 0.03 | 0.07 | 0.02 | 0.07
Al i | 0.11 |  $<$0.01 | 0.10 | 0.15 | 0.15 | 0.21
Si i | 0.11 |  $<$0.01 | 0.05 | 0.12 | 0.10 | 0.16
Ca i | 0.07 |  $<$0.01 | 0.03 | 0.08 | 0.11 | 0.13
Sc ii | 0.05 | 0.04 | 0.02 | 0.07 | 0.05 | 0.08
Ti ii | 0.05 | 0.04 | 0.02 | 0.07 | 0.12 | 0.14
V ii | 0.05 | 0.04 | 0.01 | 0.06 | 0.05 | 0.08
Cr i | 0.11 |  $<$0.01 | 0.02 | 0.11 | 0.06 | 0.13
Mn i | 0.11 |  $<$0.01 | 0.05 | 0.12 | 0.10 | 0.16
Fe i | 0.11 |  $<$0.01 | 0.05 | 0.12 | 0.13 | 0.18
Fe ii | 0.02 | 0.04 | 0.04 | 0.06 | 0.12 | 0.13
Co i | 0.11 |  $<$0.01 |  $<$0.01 | 0.11 | 0.05 | 0.12
Ni i | 0.11 |  $<$0.01 | 0.04 | 0.12 | 0.18 | 0.21
Zn i | 0.07 | 0.01 |  $<$0.01 | 0.07 | 0.10 | 0.12
Sr ii | 0.09 | 0.01 | 0.11 | 0.14 | 0.05 | 0.15
Y ii | 0.06 | 0.03 | 0.03 | 0.07 | 0.20 | 0.21
Zr ii | 0.07 | 0.03 | 0.03 | 0.08 | 0.22 | 0.23
Ba ii | 0.09 | 0.03 | 0.14 | 0.17 | 0.08 | 0.19
La ii | 0.09 | 0.03 |  $<$0.01 | 0.09 | 0.13 | 0.16
Ce ii | 0.09 | 0.03 |  $<$0.01 | 0.09 | 0.19 | 0.21
Nd ii | 0.09 | 0.03 | 0.01 | 0.10 | 0.14 | 0.17
Sm ii | 0.09 | 0.03 |  $<$0.01 | 0.09 | 0.13 | 0.16
Eu ii | 0.09 | 0.03 |  $<$0.01 | 0.09 | 0.06 | 0.11
Gd ii | 0.09 | 0.03 |  $<$0.01 | 0.09 | 0.08 | 0.12
Dy ii | 0.09 | 0.03 | 0.01 | 0.10 | 0.09 | 0.13
Er ii | 0.09 | 0.03 |  $<$0.01 | 0.09 | 0.07 | 0.12
Tm ii | 0.09 | 0.03 |  $<$0.01 | 0.09 | 0.05 | 0.11
Yb ii | 0.06 | 0.03 | 0.09 | 0.11 | 0.32 | 0.34
Pb i | 0.11 |  $<$0.01 | 0.02 | 0.11 | 0.16 | 0.20

The results for different groups of elements are described below.

###  5.1 Light elements C, N, O and isotope ratio 12C/13C

Figure 3:  Molecular 12CH and 13CH (denoted as (12)CH and (13)CH) lines in
spectrum of Pr184237 (open circles). The best fit (continuous curve) to the
observations was achieved for the 12C/13C = 7 isotopic ratio and
$\log\varepsilon_{\rm C}$ = 7.49. The blue dashed and magenda three-dot-dashed
curves correspond to 12C/13C = 95 (the solar ratio) and 4, respectively. In
two latter cases, the total carbon abundances are $\log\varepsilon_{\rm C}$ =
7.44 and 7.54.

The FERRE analysis discovered a strong enhancement of Pr184237 in carbon. In
this study, we use a pronounced C2 Swan band (5150 – 5170 Å) and determine the
C abundance with an uncertainty of 0.03 dex. We confirm the FERRE conclusion
although obtain the lower [C/Fe] = 1.77 due to adopting the lower effective
temperature.

New results are measurements for the N and O abundances. With fixed C
abundance, the N abundance was determined from the CN band in the 4214-4216 Å
region. For oxygen, we use the [O i] 6300 Å forbidden line, which is clearly
detected in the spectrum of Pr184237. The star is found to be strongly
enhanced in N and O, with [N/Fe] = 1.29 and [O/Fe] = 1.32.

Owing to the high carbon abundance, we derive the 12C/13C isotope abundance
ratio from pairs of 12CH 4210.9 Å and 13CH 4211.4 Å molecular lines (Fig. 3).
The calculations are performed with the solar mixture of C isotopes, 12C/13C =
95 (Lodders, 2021), and with 12C/13C = 9, 7, and 4. The best fit to the
observed spectrum was achieved for 12C/13C = 7. A variation in 12C/13C only
influences a little the theoretical profile of 12CH 4210.9 Å, such that the C
abundance obtained from fitting this line increases from $\log\varepsilon_{\rm
C}$ = 7.44 to 7.54, when 12C/13C decreases from 95 (12C fraction of 98.965%)
to 4 (12C fraction of 80%).

A supersolar fraction of the 13C isotope can be a signature of AGB
nucleosynthesis (e.g., Karakas, 2010; Lugaro et al., 2012; Ventura et al.,
2021). Here, the donor AGB star would likely have been of intermediate-mass
($M\gtrsim 2M_{\odot}$), such that it experienced proton capture
nucleosynthesis at the convective envelope, leading to a low 12C/13C ratio and
a large overabundance of N. Such stars are important contributors to 13C and N
in the early Galaxy (Kobayashi et al., 2020). Low-metallicity AGB stars can
also produce large overabundances of O, which are dredged to the surface
alongside C and heavy elements (Karakas, 2010; Lugaro et al., 2012). With
$T_{\rm eff}$ = 5100 K and $\log{g}$ = 2.0, Pr184237 is far from the AGB
evolutionary phase, and the self-pollution hypothesis should be rejected.
Variations in the radial velocity suggest binarity of the system, where
Pr184237 is a less massive secondary component. The primary evolved to the AGB
phase and via the mass transfer provided Pr184237 with the nuclear-processed
material. At present, the primary is, probably, a white dwarf.

###  5.2 Elements Na to Zn

Lines of the four $\alpha$-process elements, namely, Mg i, Si i, Ca i, and Ti
ii are measured in the spectrum of Pr184237. The star reveals an
$\alpha$-enhancement with tight measurements of [X/Fe] = 0.48, 0.36, 0.47, and
0.42, respectively. These numbers correspond to the NLTE calculations.

Pronounced NLTE effects are found for the resonance lines of Na i and Al i,
such that NLTE leads to the lower abundance of Na compared with the LTE case,
by 0.57 dex, and, in contrast, to a higher abundance of Al, by 0.45 dex. Red
wings of the Na i 5889 and 5895 Å lines in spectrum of Pr184237 are affected
by interstellar absorption lines, and we can only evaluate an upper limit for
the element abundance: [Na/Fe] $\leq-0.3$ in the NLTE calculations. It is
important to note that Na is certainly depleted relative to Fe and does not
follow the enhanced carbon. Aluminum does not reveal a notable deviation from
the scaled solar abundance, with [Al/Fe] = $-0.06$ in the NLTE calculations.

Abundances of Sc and V were determined from lines of their majority species,
Sc ii and V ii, which are expected to be weakly affected by the departures
from LTE. Indeed, the NLTE calculations resulted in $\Delta_{\rm NLTE}<0.01$
dex for the Sc ii 4246 Å line observed in Pr184237. Scandium appears to be
enhanced relative to Fe to the same extent, as are the $\alpha$-process
elements. Vanadium follows iron, with [V/Fe] = $-0.08$.

Abundances of Cr, Mn, Co, and Ni were derived from lines of their minority
species, which are subject to the ultraviolet overionization, resulting in
weakened lines and positive NLTE abundance corrections. For Cr i and Mn i, the
NLTE corrections are available in the database
NLTE_MPIA666http://nlte.mpia.de/ (M. Kovalev et al., 2018). They were computed
based on the methods of Bergemann & Cescutti (2010) and Bergemann et al.
(2019), respectively. For the atmospheric parameters of Pr184237, we extracted
$\Delta_{\rm NLTE}$ = 0.56 dex, 0.58 dex, 0.48 dex, and 0.26 dex for Cr i
4254, 4274, 5409 Å and Mn i 4033 Å, respectively, and obtained [Cr/Fe](NLTE) =
0.33 and [Mn/Fe](NLTE) = $-0.20$. Thus, manganese in Pr184237 is deficient
relative to Fe, independent of whether we rely on NLTE or LTE. Overabundance
of Cr relative to Fe is unlikely. When taking the Cr i\- and Fe i-based LTE
abundances, we obtained a close-to-solar ratio of [Cr i/Fe i] = $-0.09$.

For the Co i lines, the NLTE corrections are not available in Bergemann et al.
(2010) and, currently, not accessible at the above cited website. However,
they are certainly positive. Therefore, cobalt is enhanced in Pr184237, at the
level of [Co/Fe] $>$ 0.2.

When taking the LTE abundances from lines of Ni i and Fe i, we obtained [Ni
i/Fe i] = $-0.07$, suggesting that the Ni abundance in Pr184237 is close to
the scaled solar one.

Zinc is slightly enhanced, with [Zn/Fe](NLTE) = 0.21.

###  5.3 Neutron-capture elements Sr to Pb

Figure 4:  Best fits (continuous curve) to the Sr ii 4215 Å (top panel), Ba ii
5853 Å (middle panel), and Eu ii 4129 Å (bottom panel) lines in Pr184237 (open
circles). The derived abundances are indicated in Table 2. The dotted curves
in the top and bottom panels show the synthetic spectra with no Sr and Eu in
the atmosphere. The blue long-dashed and green short-dashed curves show the
effect of a 0.2 dex variation in the abundance on the synthetic spectrum.

We determine abundance for 15 elements beyond Sr (hereafter, heavy elements).
Seven of them, namely, Sr, Y, Zr, Ba, La, Ce, and Pb, are commonly referred to
as s-process elements because of the dominant contribution of the s-process to
their solar abundances, of greater than 60% according to Bisterzo et al.
(2014). Five measured elements, namely, Eu, Gd, Dy, Er, and Tm, are presumably
of r-process origin, with more than 80% contribution of the r-process to their
solar abundances. The remaining elements Nd, Sm, and Yb are intermediate
cases, with approximately equal contributions of the r- and s-process to their
solar abundances.

For Sr to Er, their abundances are derived from two to ten lines. The best
fits to the representative lines are shown in Fig. 4. Ytterbium is observed in
the only line, Yb ii 3694 Å, and the only unblended line was found for Tm ii.

Figure 5:  Best fits (continuous curve) to the Pb i 4057 Å (top panel) and
3683 Å (bottom panel) lines in Pr184237 (open circles). The derived abundances
are indicated in Table 2. The blue dashed curves show the synthetic spectra
computed with [Pb/Fe] = 0. The red dash-dotted curve in the bottom panel
corresponds to [Pb/Fe] = 1.0. The green dotted curve in the top panel shows
the synthetic spectrum computed without carbon in the atmosphere.

A challenge was to determine the Pb abundance from the Pb i 4057 Å line, which
is heavily blended by the molecular CH line. As shown in Fig. 5 (top panel),
there are three unblended or weakly blended lines of 12CH in the 4058.0-4059.7
Å range. They all are well fitted with $\log\varepsilon_{\rm C}$ = 7.58, which
agrees within 0.06 dex with the C abundance derived from the C2 Swan band. We
note that the 13CH 4058.448 Å line is reproduced reasonably well using 12C/13C
= 7, as derived in this study. It is evident that neither the position, nor
the total absorption of the 4057 Å blend can be reproduced in the case of
[Pb/Fe] = 0 ($\log\varepsilon_{\rm Pb}=-0.57$). In the LTE calculations, the
best fit to the 4057 Å blend is achieved with $\log\varepsilon_{\rm Pb}$ =
1.55. The obtained Pb abundance is supported by another line of Pb i, at 3683
Å. This line is located in a rather crowded region (Fig. 5 bottom panel),
which is nevertheless well fitted using the derived atmospheric parameters
(important for H i 3682.8 Å!) and abundances of C and Fe. For Pb i 3683 Å, we
show the synthetic profiles computed with [Pb/Fe] = 0, 1, and 2.12. It is
evident that the first two profiles are too shallow compared with the observed
one. Only Pb i 4057 Å was used to represent in Table 5 the systematic
uncertainties in the Pb abundance. When varying $T_{\rm eff}$ or $\log{g}$ or
$\xi_{t}$, we first evaluated the C abundance from the CH 4058.2, 4059.2,
4059.5 Å lines, and then the Pb abundance was obtained from fitting the 4057.7
Å blend.

We obtain that elements of the second s-process peak (Ba, La, Ce) and elements
of the intermediate group (Nd, Sm, Yb) are strongly enhanced in Pr184237, with
[X/Fe] $>1$. An even higher excess was found for Pb, with [Pb/Fe] = 2.12 and
2.78 in the LTE and NLTE calculations, respectively.

Pr184237 is the first star in the inner Galaxy for which the five r-process
elements are measured, and they all reveal an enhancement, with the average
[r/Fe] = 1.03$\pm$0.28.

Figure 6:  The element abundance pattern of Pr184237. An upper limit for the
Na abundance is indicated by downward-facing triangle. The dashed line
indicates the star’s iron abundance [Fe/H] = $-2.60$. Abundances of Cr, Mn,
Co, and Ni are increased by 0.14 dex, see text for more details.

The obtained element abundance pattern of Pr184237 is presented in Fig. 6. For
illustration purpose, the LTE abundances of Cr, Mn, Co, and Ni were increased
by 0.14 dex, which is the difference between the NLTE and LTE abundances for
Fe i. By doing so, we assumed similar NLTE effects for neutral species of
these five elements and, in the first approximation, accounted for the NLTE
effects for Cr i, Mn i, Co i, and Ni i. For all elements beyond Sr, the LTE
abundances are used for consistency.

###  5.4 Comparisons with the VMP stellar samples

For abundance comparisons, we use the VMP stellar samples observed towards the
bulge from high-resolution spectroscopic studies of Howes et al. (2016), Koch
et al. (2016), and Sestito et al. (2023).

Both Koch et al. (2016) and Sestito et al. (2023) conclude that the majority
of their carbon-normal stars exhibit abundance patterns typical for the
Galactic halo stars, although the [Na/Mg] ratios in Sestito et al. (2023)
reveal a substantial star-to-star scatter. In contrast, Howes et al. (2016)
find that the chemistry of their sample stars deviates from that for halo
stars of the same metallicity. However, we see in their Figures 11-13 that
most of the elemental ratios lie well on the halo abundance trends. The star-
to-star scatter is, indeed, observed for [Mg/Fe]. In the Na–Zn range, our CEMP
star Pr184237 does not look exceptional among the non-CEMP stars.

Compared with the three carbon and s-process enhanced stars from Koch et al.
(2016) and Sestito et al. (2023), Pr184237 has the lowest [Ba/Fe] = 1.17,
however, this is the first star discovered in the inner Galaxy that is
strongly enhanced in both s- and r-process elements. Not only Eu, but also Gd,
Dy, Er, and Tm were measured in this star. With [Ba/Eu] = $0.32$, Pr184237
likely belongs to the CEMP-r/s subclass.

Pr184237 is the only CEMP star in the inner Galaxy, for which the isotope
abundance ratio 12C/13C and the Pb abundance are derived. Earlier, the Pb
abundance, [Pb/Fe] = 1.50, was measured for a more metal-rich and non-CEMP
star with [Fe/H] = $-1.5$ and [C/Fe] = 0.4 (Koch et al., 2019).

Figure 7:  CEMP-s (blue filled diamonds), CEMP-r/s (filled circles), and the
stars with unknown Eu abundance (open green diamonds) in the abundance
diagrams [C/Fe] versus [Fe/H] (left panel), [Pb/Ba] versus [Ba/H] (middle
panel) and [Sr/Ba] versus [Ba/H]. The data are taken from the SAGA database
(Suda et al., 2008). Pr184237 is shown by the red star symbol.

In order to compare Pr184237 with the Galactic halo CEMP-s and CEMP-r/s stars,
we used the Stellar Abundances for Galactic Archaeology (SAGA) database (Suda
et al., 2008). In total, 58 stars were selected with [Fe/H] $<-1.5$, [C/Fe]
$>0.7$, and the Pb abundance measured. We separated them into two subclasses,
using the criteria by Beers & Christlieb (2005), namely, [Ba/Fe] $>1$ and
[Ba/Eu] $>0.5$ for CEMP-s and [Ba/Eu] = 0 to 0.5, [Ba/Fe] $>0$ and [Eu/Fe]
$>0$ for CEMP-r/s. The obtained subsamples of 24 CEMP-s and 18 CEMP-r/s stars,
as well as the remaining CEMP stars with unknown Eu abundance are displayed
together with Pr184237 in Fig. 7. The three groups of stars fully overlap on
the abundance diagrams [Pb/Ba] and [Sr/Ba] versus [Ba/H]. The position of
Pr184237 in these diagrams does not contradict either the CEMP-s or the
CEMP-r/s status. Only the lower [Ba/Eu] ratio compared to the criterion for
the CEMP-s subclass leads us to refer to Pr184237 to as a CEMP-r/s star.

##  6 Origin of the heavy element abundance pattern

Figure 8:  Heavy element abundance pattern of Pr184237 (diamonds) compared to
abundances of an r-II star CS 31082-001 from Siqueira Mello et al. (2013).
Abundances of CS 31082-001 were scaled to match the Eu abundance of Pr184237.

One commonly believes that the first heavy element nuclei in our Galaxy were
produced in the r-process, although astrophysical sites for the r-process are
still debated (see Cowan et al., 2021, for a thorough review). In the Galactic
halo, a contribution of the main s-process to heavy element production becomes
notable in stars starting from [Fe/H] $\simeq-2.6$, according to Simmerer et
al. (2004), or even from higher [Fe/H] $\simeq-1.4$ (Roederer et al., 2010).

Spectroscopic analyses of MP stars in the inner Galaxy provide signatures of
early r-process nucleosynthesis. Johnson et al. (2013) and Howes et al. (2015,
2016) report on europium enhancements at the level of [Eu/Fe] = 0.5 to 1.0.
Some of the stars, for example, 2MASS 18174532-3353235 (${\rm[Fe/H]=-1.67}$,
${\rm[Eu/Fe]}$ = 0.99, ${\rm[Ba/Eu]=-0.5}$, Johnson et al., 2013) and
J181505.16-385514.9 (${\rm[Fe/H]=-3.29}$, ${\rm[Eu/Fe]}$ = 0.96,
${\rm[Ba/Eu]=-0.92}$, Howes et al., 2015), can be referred to as strongly
r-process enhanced (r-II) stars, following the classification of Beers &
Christlieb (2005). As shown by Johnson et al. (2012); Van der Swaelmen et al.
(2016), and Forsberg et al. (2019); Forsberg et al. (2023), the r-process was
the dominant neutron-capture production process in the bulge until metallicity
reached [Fe/H] $\sim-1$.

Therefore, we have every reason to believe that Pr184237 has formed from the
matter where the heavy elements were of the r-process origin. Later the star’s
atmosphere was enriched with the nuclear-processed material transferred from
the evolved companion. The enrichment scenario should explain enhanced
abundances of the heavy elements and also enhancements in C, N, O and the low
12C/13C isotope ratio.

###  6.1 Pollution by the s-process products

In this scenario, we assume that heavy elements in Pr184237 originate from two
sources. The first one is the interstellar matter out of which Pr184237 has
formed, and it has the r-process element abundance pattern, with [r/Fe] =
1.03, as shown in Sect. 5.3. The second one is the s-process enriched material
transferred onto a surface of Pr184237 from the more massive primary component
when it evolved to the AGB phase. The transferred material was diluted with
the atmospheric matter.

Figure 9:  Heavy element abundance pattern of Pr184237 (diamonds) compared to
different s-process models from Lugaro et al. (2012). See the text for more
details.

Initial abundances of the heavy elements in Pr184237, $\varepsilon_{\rm r}$,
were calculated using the empirical r-process relative yields, which were
obtained from abundance analysis of a benchmark r-II star CS 31082-001
(Siqueira Mello et al., 2013). The r-II stars reveal very similar heavy
element abundance patterns (Sneden et al., 2008), suggesting a universal
r-process. In Fig. 8, the heavy element abundance pattern of Pr184237 is
displayed together with the scaled abundances of CS 31082-001. The scaling
factor was calculated from matching the Eu abundances of Pr184237 and CS
31082-001. Abundances of Eu, Gd, Dy, and Er in Pr184237 lie well on the
empirical r-process curve, suggesting an origin of these elements in the
r-process. The Tm abundance, which is also expected to be on the r-process
curve, appears to be higher. Three of the four inspected lines of Tm ii, at
3700, 3701, and 3848 Å, are heavily affected by noise in the spectrum of
Pr184237, and the abundance from Tm ii 3795 Å seems to be overestimated.
Neglecting Tm due to the uncertainty in its abundance, we computed the average
[r/Fe] = 0.91$\pm$0.12.

For the s-process abundances, $\varepsilon_{\rm s}$, we used models of the
s-process in the low-metallicity ([Fe/H] = $-2.3$) AGB stars computed by
Lugaro et al. (2012) in the broad range of stellar masses (0.9–6 $M_{\odot}$)
by varying the size of the 13C pocket, which is an uncertain parameter in AGB
nucleosynthesis.

The quality of each enrichment scenario was judged in terms of the $\chi^{2}$
statistics. The statistical estimator is determined by

$\chi^{2}=\sum(\log\varepsilon_{\rm obs}-\log\varepsilon_{\rm
mod})^{2}/\sigma_{\rm obs}^{2},$ (1)

where $\varepsilon_{\rm mod}=d\cdot\varepsilon_{\rm
r}+(1-d)\cdot\varepsilon_{\rm s}$, $d$ is a dilution factor, which is a free
parameter and takes a value between 0 and 1, $\sigma_{\rm obs}$ is the
abundance error, and the summation is carried out for the 15 elements from Sr
to Pb observed in Pr184237.

Lugaro et al. (2012) provide two grids of the s-process abundances computed
with two independent stellar evolution codes, stars and Stromlo. For the first
grid, the best fit to the observations, with $\chi^{2}$ = 210 and $d$ = 0.982,
was obtained using the model with an initial mass of $M_{init}$ = 2$M_{\odot}$
and the mass of the partial mixing zone $M_{mix}$ = 0.004$M_{\odot}$. Note
that the 13C pocket forms in the partially mixed zone; we refer to Lugaro et
al. (2012) for further details. In the Stromlo grid, the best model
($\chi^{2}$ = 200, $d$ = 0.980) has $M_{init}$ = 1.5$M_{\odot}$ and $M_{mix}$
= 0.002$M_{\odot}$. Both models are shown in Fig. 9 compared with the
observations.

Lugaro et al. (2012) computed also the models with the assumption that an
initial composition of the stars producing s-nuclei was different from a
scaled-solar composition. Namely, the matter was r-process enriched with
[r/Fe] = +0.4, +1.0, and +2.0. We neglect the models with [r/Fe] = +2.0,
because the binary components should have common initial abundances, but
Pr184237 does not reveal such a strong r-process enhancement. The other models
are checked by adopting $\varepsilon_{\rm mod}=\varepsilon_{\rm s}$. The best
fit model ($\chi^{2}$ = 331) shown in Fig. 9 has $M_{init}$ = 3$M_{\odot}$,
$M_{mix}$ = 0.0005$M_{\odot}$, and [r/Fe] = +0.4.

Judging by a purely numerical criterion, the Stromlo model of $M_{init}$ =
1.5$M_{\odot}$, with the least $\chi^{2}$ = 200, can be considered as the
"best model". However, it overpredicts abundances of Sr and Y. This is in line
with Lugaro et al. (2012), who conclude that the light-s/heavy-s element
abundance ratios in their models are too high compared with the observations
of CEMP-r/s stars. In contrast, the Pb abundance of the "best model" is lower
compared with that for Pr184237. The observed Pb abundance is reproduced by
the higher mass models, namely, of 2$M_{\odot}$ in the stars grid and of
3$M_{\odot}$ in the grid of the r-process enhanced models. Carbon enhancement
of Pr184237 ([C/Fe] = 1.77) is reasonably well reproduced by each of three
selected models ([C/Fe] = 1.51 to 1.94), when applying the dilution factor
obtained in fitting the heavy element abundance pattern and assuming the
initial C abundance to be equal to the scaled solar one. However, the
predicted 12C/13C isotope ratio is supersolar, in contrast to our finding for
Pr184237.

###  6.2 Pollution by the i-process products

It is possible to reproduce the observed abundances of CEMP-r/s stars by using
nucleosynthesis models for the intermediate (i) neutron capture process
(Dardelet et al., 2014; Hampel et al., 2016, 2019; Goswami et al., 2021). This
nucleosynthesis process introduced first by Cowan & Rose (1977) runs at
neutron densities of $N_{\rm n}$ = 1011-1015 cm-3, which are intermediate to
that of the main s-process, which typically occurs at $N_{\rm n}\lesssim
10^{8}$ cm-3, and considerably lower than r-process neutron densities. While
He-shell burning in intermediate-mass AGB stars can result in $N_{\rm n}\sim
10^{12}$ cm-3 (Fishlock et al., 2014), owing to the operation of the
22Mg($\alpha$,n)25Mg reaction, the neutron exposure (Busso et al., 1999) is
likely not high enough to initiate an i-process. The operation of the
i-process can produce enhancements in elements typically produced by both the
s- and r-process (e.g., Ba and Eu), although the resulting abundance
distribution is unique and not a superposition of the two (Hampel et al.,
2019). Favourable conditions for the i-process can be found in multiple
stellar sites, for example, proton-ingestion episodes (PIEs) in AGB stars,
super-AGB stars, and rapidly accreting white dwarfs (see, for example Hampel
et al., 2016, 2019; Denissenkov et al., 2017, 2019; Choplin et al., 2021,
2022).

In order to understand an origin of the heavy elements in Pr184237, we explore
the i-process nucleosynthesis in AGB stars, using the models of Hampel et al.
(2019) and Choplin et al. (2022), and in rapidly accreting white dwarfs
(RAWD), using the models of Denissenkov et al. (2019). In the AGB star
scenario, the investigated star is in a binary system and has accreted the
i-process material from the AGB companion (which is now a white dwarf). In the
scenario of enrichment of Pr184237 with the i-process products from the RAWD,
the investigated star would have been a low-mass third companion orbiting the
close, previously interacting, binary system with the RAWD. Denissenkov et al.
(2019) proposed that the i-process enriched star would have accreted a wind
material ejected by the RAWD. There are two possible scenarios: either the
RAWD exploded as a type Ia supernova (SN Ia), and the i-process enriched star
was ejected from the triple system, or it did not explode and the i-process
enriched star is still part of the triple system. Variations in the radial
velocity of Pr184237 suggest that our star could still orbit the surviving
component in this scenario.

For each of the i-process models, we assume some dilution of the produced
i-process nuclei with the heavy elements of the atmospheric matter, such that
$\varepsilon_{\rm mod}$ in formula (1) is determined by

$\log\varepsilon_{\rm mod}=\log(d\cdot\varepsilon_{\rm
init}+(1-d)\cdot\varepsilon_{\rm i-proc}),$ (2)

where $\varepsilon_{\rm i-proc}$ is the i-process abundance of a given
element, $\varepsilon_{\rm init}$ is the initial atmospheric abundance, and
$d$ is a dilution factor. The model abundances are fitted to the abundances of
15 heavy elements observed in Pr184237.

Figure 10:  Heavy element abundance pattern of Pr184237 (diamonds) compared to
the best fit i-process models calculated following Hampel et al. (2019).

In Fig. 10 we show the results from Hampel et al. (2019) of an i-process
calculated with constant neutron densities of different magnitude. The
calculations assume conditions typical of the intershell region of a low-mass,
low-metallicity AGB star. In the fitting procedure, the scaled solar
abundances were adopted as $\varepsilon_{\rm init}$. The best three fits
correspond to $N_{\rm n}=10^{11}$ cm-3 ($\chi^{2}\simeq$ 107), 1012 cm-3
($\chi^{2}\simeq$ 115), and 1014 cm-3 ($\chi^{2}\simeq$ 117), although
qualitatively the fits look similar. Differences show up in the region between
Sr and Ba, which are elements not observed for Pr184237.

Figure 11:  Heavy element abundance pattern of Pr184237 (diamonds) compared to
the models of an i-process nucleosynthesis in the AGB star (Choplin et al.,
2022, continuous and dotted curves for two models) and the RAWD (Denissenkov
et al., 2019, dashed curve). See text for more details.

Choplin et al. (2022) calculated 12 stellar evolution models with various
initial masses and metallicities. Proton ingestion events happen in six of
them, namely, in 1 and 2$M_{\odot}$ AGB models with [Fe/H] = $-3$, $-2.5$, and
$-2.3$. The companion to Pr184237 had enough time to evolve to the AGB phase.
Choplin et al. (2022, Table 1) predicted total lifetimes of 6.3 and 0.8 Gyr
for the $M$ = 1 and 2 $M_{\odot}$ models, respectively. When fitting the
i-process yields to the observed heavy element abundance pattern, we assume
that they are diluted with $\varepsilon_{\rm init}$ = $\varepsilon_{\rm r}$.
Here, $\varepsilon_{\rm r}$ is determined as in the previous subsection. We
obtain that the model with $M_{init}$ = 1$M_{\odot}$, [Fe/H] = $-3$, and
$\alpha$-enhancement (M1.0z3.0a) has the smallest $\chi^{2}$ = 109 ($d$ =
0.98) and reproduces well the heavy element abundance pattern of Pr184237,
including the light neutron-capture elements Sr and Y (Fig. 11). There are a
few subtle differences though for the r-process elements Gd, Dy, Er and also
for Zr and Yb. This model predicts two pulses, and the PIE arises during the
second one, such that the AGB surface is enriched in i-process products right
after the PIE. With the dilution factor $d$ = 0.98 and the scaled solar
abundance as the initial C abundance, the M1.0z3.0a model predicts the present
C enhancement as [C/Fe] = 1.42, which is slightly less than the observed
value, and the isotope ratio of 12C/13C = 4.3, which is close to the observed
value.

Figure 11 also shows the model for 2$M_{\odot}$ and [Fe/H] = $-2.5$
(M2.0z2.5). Keeping in mind that the Tm abundance of Pr184237 seems to be
overestimated (see Sect. 6.1), we find the best fits to the observed abundance
pattern that does not include Tm. The lowest $\chi^{2}$ = 54.7 ($d$ = 0.95)
was obtained for the M2.0z2.5 model. However, the M1.0z3.0a model has a
similarly small $\chi^{2}$ = 58.4 ($d$ = 0.985). Compared to M1.0z3.0a, the
M2.0z2.5 model reproduces well observed abundances of Eu to Yb (Tm is not
included) and not only Sr and Y, but also Zr. However, M2.0z2.5 is worse for
Ba, La, Ce and, with 12C/13C = 24.5, also for the carbon isotope ratio. Right
after the PIE pulse the model predicts low 12C/13C = 3.5, but an enrichment of
the AGB surface in 13C reduces during ten subsequent thermal pulses due to a
dilution of the i-process material.

For an i-process in RAWD, we use the G (iRAWD_2.6) model of Denissenkov et al.
(2019) with exactly the same metallicity as [Fe/H] = $-2.6$ for Pr184237. They
calculated an evolution of a close binary system with $M_{1}$ = 2.5$M_{\odot}$
and $M_{2}$ = 2.0$M_{\odot}$ and the i-process nucleosynthesis when the
primary reached the stage of rapidly accreting white dwarf. The best fit
($\chi^{2}$ = 119) displayed in Fig. 11 was obtained with a dilution factor of
$d$ = 0.9997. The RAWD model predicts slightly greater [C/Fe] = 2.02 value
than the observed C enhancement, but extremely high 12C/13C isotope ratio.

###  6.3 Uncertainties in the choice of enrichment scenario

In all of the considered scenarios, the initial atmospheric chemical
composition of Pr184237 is polluted by nucleosynthesis products transferred
from the more massive, evolved companion. Can we decide on the type of
nucleosynthesis and initial mass of the companion from a comparison of our
observed data with the theoretical models?

Compared with an i-process, we have more arguments against the s-process
enrichment. The best fit s-process models have greater $\chi^{2}$, by a factor
of more than two. This is mostly due to the overproduction of the light
neutron-capture elements Sr and Y. When ignoring Sr to Zr, we obtained
$\chi^{2}$ = 119 for the Stromlo model of 1.5$M_{\odot}$ shown in Fig. 9. The
s-process models produce a C enhancement close to the observed one, but with
much lower enhancements in N and O. For example, assuming the initial
abundances [N/Fe] = 0 and [O/Fe] = 0.5, we obtain [N/Fe] = 0.14 and [O/Fe] =
0.68 after the pollution. To remind the reader, Pr184237 reveals [N/Fe] = 1.29
and [O/Fe] = 1.32. The s-process models predict very high 12C/13C isotope
ratios, however, this and the low [N/Fe] predicted from the model can be due
to uncertainties in the theoretical models of low-metallicity AGB stars.

Of two sites for an i-process, AGB and RAWD, no preference can be given to any
of them with respect to $\chi^{2}$ and enhancements in C, N, and O. For
Pr184237, we derive supersolar and similar [C/N] and [C/O] ratios at the level
of 0.45. Supersolar [C/O] $\simeq$ 0.3 ratios are predicted by the M1.0z3.0a
and M2.0z2.5 models of Choplin et al. (2022) and the iRAWD_2.6 model of
Denissenkov et al. (2019). The initial O abundance was assumed to be [O/Fe] =
0.5. There are discrepancies in [C/N] between the models and observations,
however, we realise that our estimate of the N abundance in Pr184237 is less
reliable compared with that for C and O.

From a chemistry point of view, an advantage of the AGB i-process models of
Choplin et al. (2022) is a subsolar 12C/13C ratio at the surface of the
polluted star which tends to be in line with observations of Pr184237.

We realise that our conclusion about a preferrable enrichment scenario can be
influenced by the uncertainties in the theoretical models and in the element
abundances determined for Pr184237.

####  6.3.1 Uncertainties in modelling the i-process

Theoretical modelling of low and intermediate-mass stars on the AGB comes with
significant uncertainties when considering mixing between the He-shell and
envelope, and from mass-loss, some of which is likely accreted onto the
companion by stellar wind accretion. For low-metallicity AGB stars the
situation is even more complicated owing to proton-ingestion episodes, which
can drive an i-process, along with the previous uncertainties (we refer to
Karakas & Lattanzio, 2014, for a detailed discussion). Choplin et al. (2021)
perform detailed AGB calculations of proton-ingestion and found that the event
can terminate the AGB phase, a phenomenon not found by Campbell & Lattanzio
(2008). However Choplin et al. have coupled the surface opacity to the
envelope enrichment, which affects the structure significantly after the star
becomes carbon rich (e.g., see also Marigo, 2002). These authors also examine
the effect of modelling uncertainties on the nucleosynthesis predictions and
note an uncertainty of $\pm 0.3$ dex in the abundances. This value does not
take uncertainties from nuclear physics into consideration, where many nuclei
are unstable and consequently have uncertain neutron capture cross sections.

An impact of the uncertainties in the (n,$\gamma$) reaction rates and
$\beta$-decay rates for 164 unstable isotopes, from 131I to 189Hf, on the
predicted i-process abundances was investigated by Denissenkov et al. (2021)
with their RAWD ([Fe/H] = $-2.6$) model for the elements from Ba to W.
Denissenkov et al. (2021) show that variations in the (n,$\gamma$) reaction
rates result in notable changes of up to 0.5 dex (see their Figs. 7, 10, 15)
in the predicted abundances.

For the RAWD model from Denissenkov et al. (2019), the main uncertainty is
that the mass accretion rate is held steady at a low value, which leads to
H-flashes that cause the envelope of the accreted matter to expand. In order
to stop the model from becoming a red giant, an artificial mass-loss is
assumed to bring the model back to a steady mass accretion. Secondly, the low-
metallicity WDs tend to grow significantly, nearing the Chandrasekhar limiting
mass and may explode as a type Ia supernova. While low-metallicity SNeIa
likely occurred, they must have been very rare in the early Universe or we
would see evidence of this in the chemical record (e.g., lower [$\alpha$/Fe],
higher Mn abundance, etc). Côté et al. (2018) studied the i-process
contribution from RAWDs to the chemical enrichment of the galaxy, and included
an exploration of the number of RAWD systems using binary population
synthesis. Interestingly in the binary population synthesis model, the systems
that evolved to become RAWDs avoided SN Ia altogether, as a consequence of
artificially suppressing H accumulation onto the WD. While this may avoid the
issues discussed above in regard to low metallicity SN Ia, it also suggests
that the WD growth seen in the models of Denissenkov et al. (2019) may not
happen.

In summary, the i-process may have occurred in multiple stellar sites in the
early Galaxy including low-metallicity AGB stars, rapidly rotating white
dwarfs, and even perhaps massive stars (Banerjee et al., 2018; Clarkson et
al., 2018). It is unclear as yet which of these sites polluted the stars with
neutron-capture elements, with all sites still showing considerable
uncertainties stemming from modelling and input physics.

####  6.3.2 Uncertainties in the abundance determinations

As already noted, we probably overestimate the Tm abundance. If we remove Tm
from the the heavy element abundance pattern, the AGB star of not only $M_{\rm
init}$ = 1.0, but also $M_{\rm init}$ = 2$M_{\odot}$ can be the site of an
i-process that enriched Pr184237. For the s-process models, removing Tm
decreases $\chi^{2}$, for example to $\chi^{2}$ = 93.7 for the 1.5$M_{\odot}$
shown in Fig. 9, but the best fit models remain the same.

The analysis of the enrichment scenarios was based on the LTE abundances of
Pr184237. When using the NLTE abundances for Sr, Ba, Eu, and Pb and the LTE
abundances for the remaining elements, we obtain slightly different $\chi^{2}$
= 140 ($d$ = 0.99995) for the iRAWD_2.6 model. In the case of the AGB
i-process, the best fit ($\chi^{2}$ = 110, $d$ = 0.99) is achieved for the
same mass (1.0$M_{\odot}$), but a higher metallicity of [Fe/H] = $-2.5$. The
best fits for the s-process models in the Stromlo grid are shifted to the
higher masses of 2 and 2.5$M_{\odot}$. We also made fitting the i-process
models of Choplin et al. (2022) to the NLTE abundances of Sr, Ba, Eu, and Pb
without including the other heavy elements and obtained the same solution for
the best fit model, that is M1.0z2.5. NLTE leads to a 0.66 dex higher Pb
abundance of Pr184237 compared with the LTE value, and none of the considered
models predicts such a high Pb abundance. The difference between the observed
and theoretical values amounts to 0.3 dex for the Choplin et al. (2022) model
and exceeds 0.5 dex for the iRAWD_2.6 and the s-process models.

When fitting the i-process models of Choplin et al. (2022) and Denissenkov et
al. (2019), we assume that the initial chemical composition has the r-process
pattern. We tested whether using the scaled solar abundances,
$\varepsilon_{\odot}$, can influence on the stellar parameters of the best
fits. The present solar system abundances were taken from Lodders (2021). The
obtained results are very similar to that for $\varepsilon_{\rm init}$ =
$\varepsilon_{\rm r}$. The least $\chi^{2}$ = 116 ($d$ = 0.97) was achieved
for the same AGB model of 1$M_{\odot}$ and [Fe/H] = $-3$. The $\chi^{2}$
slightly increases up to $\chi^{2}$ = 122 ($d$ = 0.9997) for the iRAWD_2.6
model.

Thus, an i-process nucleosynthesis in the AGB stars with a progenitor mass of
1.0-2.0$M_{\odot}$, as calculated by Choplin et al. (2022), is a promising
scenario to explain the CEMP-r/s phenomenon at the surface of Pr184237.

##  7 Conclusions

In a sample of 20 VMP stars selected from the PIGS survey for a high-
resolution spectroscopic follow-up with UVES/VLT2, a FERRE analysis identified
a CEMP star, Pr184237, which is strongly enhanced also in Ba, La, and Eu. Its
orbit is confined to within $\sim 2.6$ kpc of the Galactic centre, and it is
most likely an early Milky Way star rather than accreted. In this study, we
improved the star’s atmospheric parameters: $T_{\rm eff}$ = 5100 K, $\log{g}$
= 2.0, and [Fe/H] = $-2.60$ and determined abundances of 32 chemical elements,
including 15 heavy elements beyond the iron group. The atmospheric parameters
and abundances of 13 elements, from Na to Pb, were derived based on the NLTE
line formation.

We improved the C abundance, [C/Fe] = 1.77, and found that Pr184237 is also
enhanced in N and O. The star reveals a low carbon isotope ratio of 12C/13C =
7, which is a signature of nuclear-processed material. Pr184237 is far from
the AGB evolutionary phase, and the self-pollution hypothesis should be
rejected. The element abundance pattern in the Na-Zn range is typical of the
Galactic halo stars. However, the neutron-capture elements, in particular lead
with [Pb/Fe] = 2.12 (LTE) and 2.78 (NLTE), are strongly enhanced. With [Ba/Eu]
= 0.32, Pr184237 is the first star of the CEMP-r/s subclass in the inner
Galaxy. Compared with the halo CEMP-r/s stars, Pr184237 reveals very similar
abundance ratios among the representative elements of the first, second and
third s-process peaks, namely, Sr, Ba and Pb.

We considered two scenarios of chemical enrichment for Pr184237. In both
cases, the star is likely in a binary or higher order multiple system,
consistent with the detected radial velocity variability.

In the first scenario, the star would have been born from the r-process rich
material enriched by previous stellar generations, and would have been
polluted with the s-process elements by a binary companion (former AGB) – this
is the r+s scenario, where the different heavy elements come from different
sources. We applied the empirical r-process yields based on abundances of an
r-II star CS 31082-001 (Siqueira Mello et al., 2013) and the s-process
abundances from calculations of metal-poor AGB stars by Lugaro et al. (2012)
and fitted a combination of them. We find that some of the observed abundances
of the first s-process peak elements (Sr, Y) cannot be explained in the r + s
scenario. This scenario does not reproduce the subsolar 12C/13C isotope ratio
derived for Pr184237.

In the second scenario, the initial chemical composition could be polluted
with both s- and r-process elements produced by the i-process nucleosynthesis,
which can occur in the companion AGB star or rapidly accreting white dwarf. We
checked the i-process models from Hampel et al. (2019), Choplin et al. (2022)
– both for AGB stars and Denissenkov et al. (2019, for RAWDs). For each grid,
the best fit models reproduce reasonably well the heavy element abundance
pattern of Pr184237. The models from Choplin et al. (2022) and Denissenkov et
al. (2019) reproduce also the C enhancement and the C/O ratio. A subsolar
12C/13C isotope ratio is only obtained when using the AGB models of
1.0-2.0$M_{\odot}$ from calculations of Choplin et al. (2022).

This study supports the hypothesis that the CEMP-r/s stars, like all CEMP-s
stars, form in binary or higher order multiple system (see Abate et al., 2016;
Goswami et al., 2021, and references therein). We consider an i-process
nucleosynthesis in the AGB stars with a progenitor mass of 1.0-2.0$M_{\odot}$
and a transfer of the nuclear-processed material onto the companion surface as
a promising scenario for explaining the CEMP-r/s phenomenon of Pr184237. This
suggests that the i-process could contribute to the chemical evolution of
carbon and heavy elements in the inner Galaxy.

## Acknowledgements

This study is based on observations made with the Very Large Telescope of the
European South Observatory (Chile) under program ID 0105.B-0078(A). This
research has made use of the data from the European Space Agency (ESA) mission
Gaia777https://www.cosmos.esa.int/gaia, processed by the Gaia Data Processing
and Analysis Consortium
(DPAC888https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the
DPAC has been provided by national institutions, in particular the
institutions participating in the Gaia Multilateral Agreement. This research
has made use of the data from the VALD, SAGA, and
ADS999http://adsabs.harvard.edu/abstract_service.html databases. This research
was enabled in part by support provided by BC DRI (British Columbia Digital
Research Infrastructure) Group and the Digital Research Alliance of
Canada101010https://alliancecan.ca. AA acknowledges support from the Herchel
Smith Fellowship at the University of Cambridge and a Fitzwilliam College
research fellowship supported by the Isaac Newton Trust. DA acknowledges
financial support from the Spanish Ministry of Science and Innovation (MICINN)
under the 2021 Ramón y Cajal program MICINN RYC2021-032609. AIK was supported
by the Australian Research Council Centre of Excellence for All Sky
Astrophysics in 3 Dimensions (ASTRO 3D), through project number CE170100013.
NFM gratefully acknowledges support from the French National Research Agency
(ANR) funded project “Pristine” (ANR-18-CE31-0017) along with funding from the
European Research Council (ERC) under the European Unions Horizon 2020
research and innovation programme (grant agreement No. 834148). FS thanks the
Dr. Margaret "Marmie" Perkins Hess postdoctoral fellowship for funding his
work at the University of Victoria. JIGH acknowledges financial support from
the Spanish Ministry of Science and Innovation (MICINN) project
PID2020-117493GB-I00.

##  8 Data availability

All our main results are available in the tables in this article. The UVES
spectrum is publicly available on the ESO archive. Other data underlying this
article will be shared on reasonable request to the corresponding author.

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