# Discovery of an Extremely $r$-process-enhanced Thin-disk Star with [Eu/H]
$=+0.78$

Xiao-Jin Xie CAS Key Laboratory of Optical Astronomy, National Astronomical
Observatories, Chinese Academy of Science, Beijing 100101, China  School of
Astronomy and Space Science, University of Chinese Academy of Sciences,
Beijing 100049, China  Jianrong Shi CAS Key Laboratory of Optical Astronomy,
National Astronomical Observatories, Chinese Academy of Science, Beijing
100101, China  School of Astronomy and Space Science, University of Chinese
Academy of Sciences, Beijing 100049, China  School of Physics and Technology,
Nantong University, Nantong 226019, China  Hong-Liang Yan CAS Key Laboratory
of Optical Astronomy, National Astronomical Observatories, Chinese Academy of
Science, Beijing 100101, China  School of Astronomy and Space Science,
University of Chinese Academy of Sciences, Beijing 100049, China  Institute
for Frontiers in Astronomy and Astrophysics, Beijing Normal University,
Beijing 102206, China  Tian-Yi Chen CAS Key Laboratory of Optical Astronomy,
National Astronomical Observatories, Chinese Academy of Science, Beijing
100101, China  School of Astronomy and Space Science, University of Chinese
Academy of Sciences, Beijing 100049, China  Carlos Allende Prieto Instituto de
Astrofísica de Canarias, Vía Láctea S/N, E-38205 La Laguna, Tenerife, Spain
Universidad de La Laguna, Departamento de Astrofísica, E-38206 La Laguna,
Tenerife, Spain  Timothy C. Beers Department of Physics and JINA Center for
the Evolution of the Elements, University of Notre Dame, Notre Dame, IN 46556,
USA  Shuai Liu CAS Key Laboratory of Optical Astronomy, National Astronomical
Observatories, Chinese Academy of Science, Beijing 100101, China  Chun-Qian Li
CAS Key Laboratory of Optical Astronomy, National Astronomical Observatories,
Chinese Academy of Science, Beijing 100101, China  School of Astronomy and
Space Science, University of Chinese Academy of Sciences, Beijing 100049,
China  Ming-Yi Ding CAS Key Laboratory of Optical Astronomy, National
Astronomical Observatories, Chinese Academy of Science, Beijing 100101, China
School of Astronomy and Space Science, University of Chinese Academy of
Sciences, Beijing 100049, China  Yao-Jia Tang CAS Key Laboratory of Optical
Astronomy, National Astronomical Observatories, Chinese Academy of Science,
Beijing 100101, China  School of Astronomy and Space Science, University of
Chinese Academy of Sciences, Beijing 100049, China  Ruizhi Zhang CAS Key
Laboratory of Optical Astronomy, National Astronomical Observatories, Chinese
Academy of Science, Beijing 100101, China  School of Astronomy and Space
Science, University of Chinese Academy of Sciences, Beijing 100049, China
Renjing Xie CAS Key Laboratory of Optical Astronomy, National Astronomical
Observatories, Chinese Academy of Science, Beijing 100101, China  School of
Astronomy and Space Science, University of Chinese Academy of Sciences,
Beijing 100049, China

(Accepted July 6, 2024)

###### Abstract

Highly $r$-process-enhanced stars are rare and usually metal-poor ([Fe/H]
$\textless$ $-$1.0), and mainly populate the Milky Way halo and dwarf
galaxies. This study presents the discovery of a relatively bright
($V=12.72$), highly $r$-process-enhanced ($r$-II) star ([Eu/Fe] = $+$1.32,
[Ba/Eu] = $-$0.95), LAMOST J020623.21$+$494127.9. This star was selected from
the Large Sky Area Multi-Object Fiber Spectroscopic Telescope (LAMOST) medium-
resolution ($R\sim 7500$) spectroscopic survey; follow-up high-resolution
($R\sim 25,000$) observations were conducted with the High Optical Resolution
Spectrograph (HORuS) installed on the Gran Telescopio Canarias (GTC). The
stellar parameters (${T_{\rm eff}}$ = 4130 K, $\rm log\,g$ = 1.52, $\rm[Fe/H]$
= $-$0.54, $\xi$ = 1.80 $\rm{km\,{s^{-1}}}$) have been inferred taking into
account non-local thermodynamic equilibrium (NLTE) effects. The abundances of
[Ce/Fe], [Pr/Fe], and [Nd/Fe] are $+$0.19, $+$0.65 and $+$0.64, respectively,
relatively low compared to the Solar $r$-process pattern normalized to Eu.
This star has a high metallicity ([Fe/H] = $-$0.54) compared to most other
highly $r$-process-enhanced stars, and has the highest measured abundance
ratio of Eu to H ([Eu/H] $=+0.78$). It is classified as a thin-disk star based
on its kinematics, and does not appear to belong to any known stream or dwarf
galaxy.

stars: abundances - stars: chemically peculiar - stars: fundamental parameters

††journal: ApJ††thanks: corresponding author<EMAIL_ADDRESS>
##  1 Introduction

The rapid neutron-capture process ($r$-process) is one of the fundamental
nucleosynthetic pathways for producing the heaviest elements. Several
astrophysical sites have been proposed to satisfy the physical conditions
required for the $r$-process, such as core-collapse supernovae (e.g., Sato,
1974; Witti et al., 1994; Farouqi et al., 2010; Mirizzi, 2015; Tsujimoto,
2021), magneto-rotational jet-driven supernovae (e.g., Symbalisty et al.,
1985; Fujimoto et al., 2008; Nishimura et al., 2015; Obergaulinger et al.,
2018; Reichert et al., 2023), collapsars (Siegel et al., 2019), and neutron
star mergers (NSMs, Lattimer & Schramm, 1974; Rosswog et al., 2014; Thielemann
et al., 2017). The detection of GW170817 (Abbott et al., 2017) and its
electromagnetic counterpart shows that NSMs can indeed produce $r$-process
elements (Drout et al., 2017; Kasen et al., 2017; Pian et al., 2017; Watson et
al., 2019). Research is still ongoing to determine whether NSMs are the sole
source of $r$-process elements; some studies suggest that the observed
$r$-process-element abundance patterns requires multiple sites (e.g., Côté et
al., 2018, 2019; Hotokezaka et al., 2018; Siegel, 2019; Brauer et al., 2021;
Farouqi et al., 2021; Tsujimoto, 2021; Yamazaki et al., 2021; Naidu et al.,
2022; Ekanger et al., 2023; Kobayashi et al., 2023).

Generally, $r$-process-enhanced (RPE) stars are metal-poor, with [Fe/H]
$<-$1.0 (e.g., McWilliam et al., 1995; Sneden et al., 2008; Holmbeck et al.,
2020), and they are mainly restricted to the Milky Way (MW) halo and dwarf
galaxies (e.g., Hirai et al., 2015; Ji et al., 2016; Jeon et al., 2021;
Matsuno et al., 2021). These stars are classified as $r$-I, $r$-II, or
limited-$r$, depending on the level of $r$-process enrichment (Beers &
Christlieb, 2005; Frebel, 2018; Holmbeck et al., 2020). The dominant
$r$-process patterns (Ba to Hf) of these stars are nearly identical (sometimes
referred to as “universaility”), although some star-to-star differences
between the lightest and heaviest elements remain (e.g., Sneden et al., 1994;
Aoki et al., 2017; Sakari et al., 2018; Roederer et al., 2022a). Information
on the elemental abundances and kinematics of these stars constrain their
early enrichment history, which is crucial to understanding the origin of the
$r$-process.

This work reports on the discovery of an RPE star with [Fe/H] $>-$0.6 that is
kinematically associated with the thin disk of the MW. Section 2 introduces
the observations, and Section 3 describes the determination of stellar
atmospheric parameters. A detailed abundance analysis and estimates of
uncertainties is presented in Section 4. A discussion of the implications of
the existence of this star, and an old (but uncertain) estimate of its
nucleosynthetic age based on the observed [Th/Eu] ratio, is provided in
Section 5, followed by a summary of our conclusions in Section 6.

##  2 Observations

The RPE candidate LAMOST J020623.21$+$494127.9 (hereafter J 0206$+$4941) was
selected from the Large Sky Area Multi-Object Fiber Spectroscopic Telescope
(LAMOST, Cui et al., 2012; Zhao et al., 2012; Liu et al., 2020; Yan et al.,
2022) medium-resolution ($R\sim$ 7500) spectroscopic survey by matching the
observed spectra to synthetic templates of the region around the Eu II line at
6645.1 Å (Chen et al., 2021). This star is relatively bright ($V<$ 13) and
substantially more metal-rich compared to previously observed $r$-II stars.

The high-resolution spectrum was taken with the High Optical Resolution
Spectrograph (HORuS) installed on the Gran Telescopio Canarias (GTC, Tabernero
et al., 2020; Allende Prieto, 2021). The spectral wavelength coverage is from
3800 to 6900 Å with a resolving power of $R\sim 25,000$. The multi-order
echelle spectra were extracted and wavelength calibrated based on Th-Ar
hollow-cathode lamp exposures, using the HORuS automated pipeline the
chain111The chain is written in the Interactive Data Language (IDL), runs in
GDL, and is available from github.com/callendeprieto/chain.. The signal-to-
noise ratio (S/N) at 5600 Å is estimated to be 119, using the code DER_SNR
(Stoehr et al., 2008). Table 1 lists the basic information of this target,
including the details of the observations, kinematics, and stellar atmospheric
parameters, as described in detail below.

Table 1: Target Information

Parameter | Value
---|---
Name | LAMOST J 020623.21$+$494127.9
R.A. (J2000) | 02:06:23.211
Decl. (J2000) | +49:41:27.914
$V$ mag | 12.72 $\pm$ 0.24
obsDate LAMOST | 2018-01-01
obsDate GTC | 2019-11-18
$\rm{{RV}_{helio}}$ (km s-1) |  $-$30.86 $\pm$ 0.67
$\rm{Parallax}$ (mas) | 0.3106 $\pm$ 0.0156 a
$\rm{{U}_{LSR}}$ (km s-1) | 31.11 $\pm$ 0.49
$\rm{{V}_{LSR}}$ (km s-1) |  $-$11.89 $\pm$ 0.50
$\rm{{W}_{LSR}}$ (km s-1) | 9.65 $\pm$ 0.35
${{r}_{apo}}$ (kpc) | 10.79 $\pm$ 0.17
${{r}_{peri}}$ (kpc) | 9.68 $\pm$ 0.06
${{Z}_{max}}$ (kpc) | 0.66 $\pm$0.03
$ecc$ | 0.054$\pm$0.004
$E$ ($10^{3}$ km2 s-2) |  $-$144.31 $\pm$ 0.60
$J_{r}$ (kpc km s-1) | 4.82 $\pm$ 0.85
$J_{\phi}$ (kpc km s-1) | 2353.05 $\pm$ 24.66
$J_{z}$ (kpc km s-1) | 8.27 $\pm$ 0.51
${T_{\rm eff}}$ (K) | 4078 $\pm$ 100 b
$\rm log\,g$ (cgs) |  1.62 $\pm$ 0.08b
${T_{\rm eff}}$ (K) | 4130 $\pm$ 100 c
$\rm log\,g$ (cgs) | 1.52 $\pm$ 0.20 c
$\rm[Fe/H]$ |  $-$0.54 $\pm$ 0.10 c
$\xi_{\rm t}\rm(km\,{s^{-1}})$ | 1.80$\pm$ 0.20 c

  * a

Gaia Collaboration et al. (2023)

  * b

Determined by a photometric method.

  * c

Determined by a spectroscopic approach.

The high-resolution spectrum of J 0206$+$4941 around 6645 Å is shown in Figure
1, along with the spectrum of HD 124897, a star with similar atmospheric
parameters (${T_{\rm eff}}$ = 4286 K, log g = 1.6 cgs, [Fe/H] = $-$0.52, and
$\xi$ = 1.58 km $\rm{s^{-1}}$, Blanco-Cuaresma et al. 2014). The slightly
higher ${T_{\rm eff}}$ of HD 124897 leads to stronger Fe-group lines. The
lines of heavy elements are labeled in red. HD 124897 has a moderate
$r$-process enhancement of [Eu/Fe] = +0.45 (Forsberg et al., 2019a), and is
classified as an $r$-I star. As shown in the figure, the Eu II line of our
candidate star is significantly stronger than that of HD 124897. Many lines of
heavy elements are weak or absent in the spectra of HD 124897, but they are
strong in the spectra of our candidate star, indicating that J 0206$+$4941 has
significant $r$-process enhancement.

Figure 1: Comparison of high-resolution spectra of J 0206$+$4941 and a
moderately $r$-process-enhanced ($r$-I) star, HD 124897 ([Eu/Fe] = +0.45,
Forsberg et al. 2019a) with similar stellar parameters. The lines of heavy
elements are labeled in red, and lines of lighter elements are labeled in
black. The stellar parameters ( ${T_{\rm eff}}$, log g, [Fe/H], $\xi$ ) are
shown in the legend.

##  3 Stellar Atmospheric Parameters

We first derive initial estimates of effective temperature and surface gravity
of J 0206+4941 using a photometric approach. For comparison, we also determine
the atmospheric parameters using a spectroscopic approach: effective
temperature ($T_{\rm eff}$), surface gravity (log g), metallicity ([Fe/H]),
and microturbulent velocity ($\xi_{\rm t}$).

###  3.1 Photometric Approach

We estimate the effective temperature of our target star based on the
photometric color index ($V-K$)0 with the empirical calibration relations
given by Alonso et al. (1999). An adopted reddening $E(B-V)$ of
$0.1518\,\pm\,0.0083$ is derived using the GALEXTIN online tool (Amôres et
al., 2021) with the Green et al. (2019) 3D dust map. The $E(V-K)$ reddening is
determined using the conversions from Ramírez & Meléndez (2005). The surface
gravity is calculated using the online code PARAM 1.3 (da Silva et al., 2006),
based on the relationships between bolometric flux, temperature, mass, and log
g. The calculation requires estimates of $T_{\rm eff}$, [Fe/H], $V$ magnitude,
and parallax, where we use [Fe/H] measured by the spectroscopic approach
described below, the $V$ magnitude from Zacharias et al. (2013), and the
parallax from Gaia DR3 (Gaia Collaboration et al., 2023).

###  3.2 Spectroscopic Approach

The line list and atomic data of Fe I and Fe II lines are taken from Yan et
al. (2018). Only lines with equivalent widths (EWs) between 20 mÅ and 120 mÅ,
and not heavily blended with other species, have been used. This selection
results in 45 Fe I and 7 Fe II lines, respectively. The abundances and EWs of
individual Fe lines are analyzed by the interactive IDL code Spectrum
Investigation Utility (SIU, Reetz, 1999) using the MARCS stellar atmospheric
models (Gustafsson et al., 2008). The effects of non-local thermodynamic
equilibrium (NLTE) are considered specifically for each Fe I line, based on
Mashonkina et al. (2011), following the method described in Sitnova et al.
(2015).

A spectroscopic estimate of $T_{\rm eff}$ is obtained from the excitation
equilibrium of Fe I lines with excitation energies (Eexc) higher than 2.0 eV
(Sitnova et al., 2015), log $g$ is estimated from the ionization balance of Fe
I and Fe II, and $\xi_{\rm t}$ is obtained by forcing the iron abundances from
Fe I lines to be independent of their EWs. We adjust $T_{\rm eff}$, log $g$,
and $\xi_{\rm t}$ by restricting the slope of log$\epsilon$(Fe I) versus Eexc
lower than 0.001 dex eV-1, the difference between [Fe II/H] and [Fe I/H]
within 0.1 dex, and the slope of log$\epsilon$(Fe I) versus log(EW) less than
0.01.

The results determined by both methods are listed in Table 1. The difference
between the photometric and spectroscopic effective temperatures is only 52 K,
and the difference in surface gravities between the two methods is 0.10 dex.
We adopt the spectroscopic parameters in the following analysis.

Figure 2: Spectral synthesis for the abundance determinations of Sr, Ba, La,
Ce, Eu and Th. The red line indicates the best-fit synthesis for the line(s)
of interest. The blue shaded regions indicate $\pm$0.30 dex deviations from
the best fit. The gray-dashed lines indicate synthetic spectra in the absence
of the element of interest.

##  4 Analysis of Elemental Abundances

###  4.1 Abundance Estimates

We use a spectrum-synthesis approach for the abundance determination. The line
list and atomic data are taken from Zhao et al. (2016) and Roederer et al.
(2018b); we have slightly revised the log $gf$ values by fitting the Solar
spectrum (Kurucz et al., 1984). The influence of hyperfine structure (HFS) is
considered for Sc (Zhang et al., 2008), Cu (Shi et al., 2014), Sr (Borghs et
al., 1983), Ba (Kurucz & Bell, 1995), and Eu (Lawler et al., 2001). We
calculate the theoretical line profiles by SIU, using the MARCS stellar model
atmospheres (Gustafsson et al., 2008). The Solar abundances are adopted from
Asplund et al. (2009).

The carbon abundance is estimated from the molecular CH $G$-band region around
4310 Å; the line list we adopted comes from Masseron et al. (2014). Among the
light odd-$Z$ elements, we can obtain abundance estimates of Na and Al. Four
Na I lines at 5682 Å, 5688 Å, 6154 Å, and 6160 Å, and two Al I lines at 6696 Å
and 6698 Å have been used. The $\alpha$-elements with detectable lines include
Mg, Si, Ca, and Ti, while the [O I] line at 6300 Å is too weak to derive a
reliable abundance. NLTE effects have been taken into consideration, based on
Gehren et al. (2004, Na), Baumueller & Gehren (1996, Al), Mashonkina (2013,
Mg), Shi et al. (2008, Si), Mashonkina et al. (2007, Ca), Sitnova et al.
(2016, Ti), and Shi et al. (2014, Cu). In addition, the NLTE corrections for
lines of Cr I, Mn I, and Co I are obtained from an online spectral tool
222https://nlte.mpia.de (Kovalev, 2019).

The neutron-capture elements frequently have their absorption lines blended
with other elements. This is especially the case for relatively metal-rich
stars, as strong Fe lines often lead to strong blends. Figure 2 shows a sample
of spectral synthesis for the Sr I, Ba II, La II, Ce II, Eu II, and Th II
lines. We use the Sr I line at 4607 Å to determine the Sr abundance. There are
numerous very strong absorption lines for elements, such as Fe and Ti lines
around 4077 Å and 4161 Å, making it difficult to distinguish the absorption
lines of Sr II. The abundance of Ba is derived from the Ba II lines at 4554 Å,
4934 Å, 6141 Å, and 6496 Å. The NLTE effects have been taken into account
(Mashonkina et al., 1999). The abundance of Eu is derived from the Eu II lines
at 4129 Å, 4205 Å, 6049 Å, and 6645 Å, with the NLTE effects considered
(Mashonkina & Gehren, 2000). As the lines at 4129 Å and 4205 Å are heavily
blended, only upper limits on Eu are obtained from these two lines. The
abundance of Ce is derived from the Ce II lines at 4418 Å, 4628 Å, 5247 Å. The
abundance of Th is estimated by using the Th II line at 4019 Å. The abundances
for a total of 15 neutron-capture elements, including Sr, Y, Zr, Ba, La, Ce,
Pr, Nd, Sm, Eu, Gd, Dy, Lu, Hf, and Th have been determined; the results are
listed in Table 2.

###  4.2 Uncertainties

To estimate the uncertainty of the abundance determinations, two types of
error sources are considered: systematic errors (${\sigma_{sys}}$, due to
errors in the stellar atmospheric parameters), and statistical errors
(${\sigma_{stat}}$, due to the line measurements themselves). Table 2 lists
the abundance differences when changing $T_{\rm eff}$ by $+100$ K, log $g$ by
$+0.2$ dex, [Fe/H] by $+0.1$ dex, and $\xi_{\rm t}$ by $+0.2$ km
$\rm{s^{-1}}$. The ${\sigma_{sys}}$ is calculated by taking the square root of
the quadratic sum of the errors associated with the atmospheric parameters.
The ${\sigma_{stat}}$ is the line-to-line dispersion divided by $\sqrt{N}$,
where $N$ is the number of lines used for a given element. Finally, we
computed the total error, represented as ${\sigma_{Total}}$, by summing
${\sigma_{sys}}$ and ${\sigma_{stat}}$ in quadrature.

Table 2: Elemental Abundances and Uncertainties Element | N | [X/Fe] | [X/Fe] | ${\sigma_{stat}}$ | $\Delta\,{T_{\rm eff}}$ |  $\Delta$ log g |  $\Delta$ [Fe/H] | $\Delta\,\xi$ | ${\sigma_{sys}}$ | ${\sigma_{Total}}$
---|---|---|---|---|---|---|---|---|---|---
LTE | NLTE | (dex) |  $+$100 K |  $+$0.2 dex |  $+$0.1 dex |  $+$0.2 $\rm{km}\,\rm{s^{-1}}$ | (dex) | (dex)
C (CH) | … |  $-$0.20 | … | … | … | … | … | … | … | …
Na I | 4 |  $+$0.26 |  $+$0.05 | 0.03 | 0.04 |  $-$0.01 | 0.00 |  $-$0.07 | 0.08 | 0.11
Mg I | 3 |  $+$0.31 |  $+$0.24 | 0.02 | 0.02 |  $-$0.02 | 0.00 |  $-$0.06 | 0.07 | 0.09
Al I | 2 |  $+$0.18 |  $+$0.08 | 0.01 | 0.07 | 0.01 |  $-$0.01 |  $-$0.03 | 0.08 | 0.09
Si I | 3 |  $+$0.23 |  $+$0.22 | 0.05 |  $-$0.08 | 0.03 | 0.01 |  $-$0.04 | 0.09 | 0.14
Ca I | 7 |  $-$0.02 |  $-$0.04 | 0.05 | 0.07 | 0.00 | 0.02 |  $-$0.09 | 0.12 | 0.17
Sc II | 4 |  $+$0.12 | … | 0.08 |  $-$0.01 | 0.04 | 0.02 |  $-$0.06 | 0.08 | 0.16
Ti II | 3 |  $+$0.09 |  $+$0.08 | 0.04 |  $-$0.03 | 0.04 | 0.02 |  $-$0.07 | 0.09 | 0.13
V I | 1 |  $+$0.10 | … | … | 0.02 | 0.03 | 0.01 |  $-$0.06 | 0.07 | 0.07
V II | 1 |  $+$0.06 | … | … | 0.02 | 0.04 | 0.00 |  $-$0.07 | 0.08 | 0.08
Cr I | 4 |  $-$0.02 |  $+$0.09 | 0.02 |  $-$0.01 | 0.04 | 0.02 |  $-$0.06 | 0.08 | 0.10
Mn I | 3 |  $+$0.10 |  $+$0.17 | 0.01 | 0.03 | 0.01 | 0.02 |  $-$0.09 | 0.10 | 0.11
Co I | 1 |  $+$0.00 |  $-$0.08 | … | 0.01 | 0.03 | 0.01 |  $-$0.07 | 0.08 | 0.08
Ni I | 5 |  $-$0.10 | … | 0.01 | 0.01 | 0.03 | 0.00 |  $-$0.06 | 0.07 | 0.08
Cu I | 3 |  $-$0.07 |  $-$0.04 | 0.02 |  $-$0.01 | 0.02 | 0.01 |  $-$0.13 | 0.13 | 0.15
Zn I | 2 |  $-$0.19 | … | 0.01 | 0.02 | 0.02 | 0.00 |  $-$0.10 | 0.10 | 0.11
Sr I | 1 |  $-$0.09 | … | … | 0.12 | 0.00 | 0.00 |  $-$0.12 | 0.17 | 0.17
Y II | 3 |  $+$0.41 | … | 0.05 |  $-$0.02 | 0.04 | 0.02 |  $-$0.01 | 0.05 | 0.10
Zr II | 3 |  $+$0.67 | … | 0.02 |  $-$0.03 | 0.04 | 0.03 |  $-$0.02 | 0.06 | 0.08
Ba II | 3 |  $+$0.42 |  $+$0.37 | 0.08 | 0.02 | 0.03 | 0.02 |  $-$0.12 | 0.13 | 0.21
La II | 2 |  $+$0.57 | … | 0.03 | 0.03 | 0.04 | 0.03 |  $-$0.03 | 0.07 | 0.10
Ce II | 3 |  $+$0.19 | … | 0.10 | 0.02 | 0.02 | 0.02 |  $-$0.09 | 0.10 | 0.20
Pr II | 2 |  $+$0.65 | … | 0.06 | 0.03 | 0.04 | 0.02 |  $-$0.01 | 0.05 | 0.11
Nd II | 3 |  $+$0.64 | … | 0.03 | 0.00 | 0.04 | 0.01 |  $-$0.10 | 0.11 | 0.14
Sm II | 3 |  $+$1.10 | … | 0.08 | 0.02 | 0.03 | 0.03 |  $-$0.07 | 0.08 | 0.16
Eu II | 2 |  $+$1.20 |  $+$1.32 | 0.07 | 0.00 | 0.04 | 0.02 |  $-$0.03 | 0.05 | 0.12
Gd II | 4 |  $+$1.14 | … | 0.18 | 0.01 | 0.04 | 0.02 |  $-$0.03 | 0.05 | 0.23
Dy II | 1 |  $+$1.30 | … | … | 0.02 | 0.03 | 0.00 |  $-$0.02 | 0.04 | 0.04
Lu II | 1 |  $+$0.20 | … | … |  $-$0.02 | 0.03 | 0.02 |  $-$0.01 | 0.04 | 0.04
Hf II | 1 |  $+$1.10 | … | … |  $-$0.02 | 0.05 | 0.03 |  $-$0.02 | 0.06 | 0.06
Th II | 1 |  $+$1.00 | … | … | 0.02 | 0.02 | 0.03 |  $-$0.15 | 0.16 | 0.16

Figure 3: Comparison of abundance ratios in J 0206$+$4941 (red star) with MW
field stars (gray circles). The abundance ratios of the MW stars come from
Venn et al. (2004), Bensby et al. (2014), Battistini & Bensby (2016), Forsberg
et al. (2019b), and Holmbeck et al. (2020). The black-dashed line is the Solar
reference line. The orange-dashed line in the [Eu/Fe] vs. [Fe/H] panel
indicates the $r$-II limit (+0.7) as defined by Holmbeck et al. (2020).
Figure 4: The heavy-element abundances of J 0206$+$4941 are shown in the top
panel, with the Solar $r$-process pattern, abundances of HD 222925 (Roederer
et al., 2022b), and 2MASS J22132050-5137385 (Roederer et al., 2024) normalized
to match the Eu abundance of J 0206$+$4941\. Residuals between J 0206$+$4941
and the Solar $r$-process pattern, HD 222925, and 2MASS J22132050-5137385 are
shown in the lower panels. Figure 5: Toomre Diagram for J 0206$+$4941\. The MW
thin-disk stars (pink triangles) and thick-disk stars (blue circles) are from
Bensby et al. (2003) and Venn et al. (2004), respectively.

##  5 Discussion

###  5.1 The Light and Fe-peak Elements

The observed carbon abundance of J 0206$+$4941 is deficient, relative to the
Solar value ([C/Fe] = $-$0.20). However, this star falls in the red giant
region, therefore, its carbon has been depleted during the normal course of
giant-branch evolution. Using the procedure described by Placco et al. (2014),
the carbon depletion for J 0206$+$4941 is $\sim$ 0.2 dex, yielding a natal
[C/Fe] $\approx$ 0.00, indicating that this star was likely never C-enhanced.
From inspection of Figure 3, the abundances of Mg, Si, and Ti appear similar
to those of the MW field stars. Although the [Ca/Fe] ratio is slightly lower,
it is still close to the boundary of other field stars. The [Na/Fe] and
[Al/Fe] ratios (not shown in the figure) are both slightly super-Solar.

Among the Fe-group elements, [V/Fe], [Cr/Fe], and [Mn/Fe] are slightly over-
abundant, while [Co/Fe], [Ni/Fe], [Cu/Fe], and [Zn/Fe] are slightly under-
abundant. All of these abundance ratios are normal compared to stars of
similar metallicity (e.g., Battistini & Bensby, 2015).

###  5.2 The $r$-process Enhancement and Pattern

The lower two panels of Figure 3 show the [Ba/Fe] and [Eu/Fe] ratios of J
0206$+$4941 compared to MW field stars. From inspection, this star exhibits a
very strong enhancement of [Eu/Fe] ($+$1.32), and a slight enhancement of
[Ba/Fe] ($+$0.37). The low [Ba/Eu] ratio ($-0.95$) indicates that its neutron-
capture elements are mainly due to the $r$-process.

In order to put J 0206$+$4941 in the context of other highly $r$-process-
enhanced stars, we have listed the known RPE stars with [Eu/Fe] $\geq$ +1.0 in
Table A1, ordered by decreasing [Eu/H]. As shown in this table, most of these
stars have low metallicity ([Fe/H] $\textless$ $-$1.0) and are identified as
members of the MW halo, while a few such stars are found in dwarf galaxies or
globular clusters. J 0206$+$4941 is the first strongly $r$-process-enhanced
star associated with the MW disk, and has a relatively high metallicity
([Fe/H] = $-$0.54). The [Eu/H] of J 0206$+$4941 is +0.78, which is the highest
among known RPE stars.

Many previous studies have found that the main $r$-process abundance patterns
of metal-poor RPE stars are very close to that of the Solar $r$-process. This
agreement is most noticeable for the elements from Ba to Hf, although there
are some differences in light neutron-capture elements (e.g., Sneden et al.,
1994; Roederer et al., 2014b; Sakari et al., 2018), and star-to-star
variations among the actinides. A comparison of the heavy-element abundances
in J 0206$+$4941 with the Solar $r$-process and two RPE “template” stars (HD
222925; Roederer et al. 2022b, and 2MASS J22132050-5137385; Roederer et al.
2024) is shown in the top panel of Figure 4, where the Solar $r$-process, HD
222925, and 2MASS J22132050-5137385 $r$-process patterns are normalized to the
Eu abundance of J 0206$+$4941\. The residuals are presented in the bottom
panels.

The abundance pattern of J 0206$+$4941 is not perfectly consistent with the
Solar $r$-process. The elements Sr, Y, Ce, Pr, and Nd are significantly lower
compared to the Solar pattern, while only Sr and Ce are significantly lower
compared to that of HD 222925. When compared to 2MASS J22132050-5137385, the
Ce abundance remains lower and the Y and Zr are higher. This difference
suggests that the production of $r$-process elements may require multiple
sites or multiple mechanisms, or that this star’s abundances have a different
contribution for these elements.

###  5.3 Kinematics

We measure the heliocentric radial velocity ($\rm{{RV}_{helio}}$ = $-$30.86
$\pm$ 0.67 km s-1) from the high-resolution GTC spectrum; it shows no obvious
variation from the LAMOST DR7 ($-$30.94 $\pm$ 1.00 km s-1; Luo et al., 2022),
and the Gaia DR3 measurements ($-$30.00 $\pm$ 0.40 km s-1; Gaia Collaboration
et al., 2023). As a result, J 0206$+$4941 is not likely to be in a binary
system.

We calculate the local standard of rest (LSR) Galactic space velocities
($\rm{U_{LSR}}$, $\rm{V_{LSR}}$, $\rm{W_{LSR}}$) for J 0206$+$4941, adopting
the coordinates, proper motions, and parallax with zero-point offset
corrected333The zero-point offset of the parallax is $-$0.014856 mas, as
obtained using the package from https://gitlab.com/icc-
ub/public/gaiadr3_zeropoint (Lindegren et al., 2021). from the Gaia DR3
catalog (Gaia Collaboration et al., 2023), as well as the radial velocity
derived from the GTC spectrum. The distance derived from the corrected
parallax is $d=3.072\pm 0.148\;\mathrm{kpc}$, which is consistent with the
photogeometric distance from Bailer-Jones et al. (2021, $d=3.128\pm
0.151\;\mathrm{kpc}$). We use a right-handed Galactocentric Cartesian
coordinate system and assume that the Sun is located at
($X_{\odot},Y_{\odot},Z_{\odot}$) = ($-$8.2, 0.0, 0.0208) $\mathrm{kpc}$
(McMillan, 2017; Bennett & Bovy, 2019), the circular velocity at the Solar
position is $v_{0}$ = 232.8 km s-1 (McMillan, 2017), and the Solar peculiar
motion is ($\rm{U_{\odot}}$, $\rm{V_{\odot}}$, $\rm{W_{\odot}}$) = (11.1,
12.24, 7.25) km s-1 (Schönrich et al., 2010). We adopt the MW potential from
McMillan (2017), and employ the AGAMA routine (Vasiliev, 2019) to derive the
orbital energy ($E$), three-dimensional action ($J=(J_{r},J_{\phi},J_{z})$),
pericentric radius (${r_{peri}}$), apocentric radius (${r_{apo}}$), maximum
height from the Galactic midplane (${{Z}_{\mathrm{max}}}$), and eccentricity
($ecc=\frac{r_{apo}-r_{peri}}{r_{apo}+r_{peri}}$). We assume Gaussian
distributions for its observables, and additionally account for the correlated
observational uncertainties in proper motions and parallax. We run a Monte
Carlo re-sampling procedure based on 100,000 draws to estimate the
uncertainties of the dynamical parameters.

J 0206$+$4941 exhibits typical thin-disk velocities (see Figure 5) with a very
low orbital eccentricity ($ecc=0.054\pm 0.004$) and limited orbital height
(${{Z}_{\mathrm{max}}=0.66\pm 0.03}$ kpc). Many previously discovered highly
$r$-process-enhanced stars have accretion origin (e.g., Roederer et al.,
2018a; Xing et al., 2019; Sakari et al., 2019). However, J 0206$+$4941 does
not appear to be clustered in dynamical phase space with any known accreted
structures (Naidu et al., 2020; Reichert et al., 2021; Hattori et al., 2023;
Shank et al., 2023), thus it may well have been formed in-situ in the MW’s
thin disk. This star may have formed from an ISM enriched in $r$-process
elements by a single binary neutron star merger and/or CCSNe, indicating that
mixing may not be isotropic, even in the thin disk. We also note that Hirai et
al. (2022) found, in their simulation of the origin of RPE stars in a MW-like
galaxy, that all simulated $r-$II stars with [Fe/H] $>-1$ were formed in-situ,
rather than via accretion from a parent dwarf satellite.

###  5.4 Age Estimation

The presence of the radioactive element Th allows us to estimate the age of
this star, albeit with a large error bar. The age can be estimated using the
following relation (Cayrel et al., 2001):

$t=46.67\,{\rm
Gyr\,[\,log\,\epsilon\,(Th/Eu)_{0}-log\,\epsilon\,(Th/Eu)_{obs}],}$

where $\rm log\epsilon(Th/Eu)_{0}$ is the initial production ratio (PR)
corresponding to the element formation at $t=$ 0, and $\rm
log\epsilon(Th/Eu)_{obs}$ is the observed ratio after the radioactive element
Th has decayed within a time $t$. We adopted an initial PR from Schatz et al.
(2002), and the calculated age is 12.32 Gyr. The uncertainty is estimated only
from the propagation of the abundance-measurement uncertainty. Only one Th
line can be measured, and the Th abundance has a systematic error of 0.16 dex,
which leads to a large age uncertainty of 9.08 Gyr.

##  6 Conclusions

We describe the discovery of an extremely $r$-process-enhanced thin-disk star,
J 0206$+$4941\. It was originally selected from the LAMOST medium-resolution
survey as a candidate RPE star, based on its unusually strong lines of Eu. A
high-S/N, moderately high-resolution ($R\sim 25,000$) follow-up observation
was obtained using HORuS on the GTC. We derive the abundances of 30 elements,
and determine the star’s kinematics. Our main findings can be summarized as
follows:

  1. 1.

J 0206$+$4941 is the most metal-rich ([Fe/H] = $-$0.54) highly $r$-process-
enhanced star found to date, and exhibits the highest known abundance ratio of
the $r$-process element Eu relative to H ([Eu/H] = +0.78).

  2. 2.

The abundances of the light elements in this star with $Z\leq$ 30 are
commensurate with other stars of similar metallicity and evolutionary phase,
with a strong enhancement of Eu, ([Eu/Fe] = +1.32) and modest enhancement of
Ba ([Ba/Fe] = +0.37). The observed [Ba/Eu] ratio ([Ba/Eu] = $-$0.95) indicates
that there is relatively little contribution from the $s$-process to the
enrichment of the neutron-capture elements.

  3. 3.

The abundance pattern of the elements with $Z\geq$ 30 is not perfectly
consistent with the scaled Solar $r$-process pattern, as demonstrated by the
lower abundances of Ce, Pr, and Nd. This difference provides some support for
multiple origins of the $r$-process.

  4. 4.

J 0206$+$4941 does not exhibit significant radial-velocity variations,
indicating it is not likely to be in a binary system.

  5. 5.

The Galactic space velocities and orbital parameters show that J 0206$+$4941
is restricted to the thin disk of the MW, and apparently was not accreted from
an external dwarf galaxy.

  6. 6.

The presence of the radioactive element Th allows us to estimate the age of J
0206$+$4941 as 12.32 $\pm$ 9.08 Gyr.

Clearly, a higher-resolution ($R>50,000$) spectrum of this star with better
S/N in the blue region would prove illuminating, and enable the measurement of
numerous other $r$-process elemental abundances, possibly including uranium,
enabling derivation of a more accurate age estimate. In addition to J
0206$+$4941, we have also found a few other candidate RPE stars with [Fe/H]
$>-$1.0 from the LAMOST-MRS spectra. These stars are excellent candidates for
high-resolution follow-up spectroscopic observations with HORuS and other
instruments in the near future.

We thank the anonymous referee for valuable suggestions that have improved the
manuscript. This research is supported by the National Natural Science
Foundation of China under Grant Nos. 12090040, 12090044, 12373036, 12022304,
11973052, the National Key R$\rm\&$D Program of China No.2019YFA0405502, the
Scientific Instrument Developing Project of the Chinese Academy of Sciences,
Grant No. ZDKYYQ20220009, and the International Partnership Program of the
Chinese Academy of Sciences, Grant No. 178GJHZ2022047GC. H.-L.Y. acknowledges
the support from the Youth Innovation Promotion Association, Chinese Academy
of Sciences. CAP is thankful for financial support from the Spanish Ministry
MICINN projects AYA2017-86389-P and PID2020-117493GB-I00. T.C.B. acknowledges
support from grant PHY 14-30152; Physics Frontier Center/JINA Center for the
Evolution of the Elements (JINA-CEE), and from OISE-1927130: The International
Research Network for Nuclear Astrophysics (IReNA), awarded by the US National
Science Foundation.

This paper is based on observations made with the Gran Telescopio Canarias
(GTC), installed at the Spanish Observatorio del Roque de los Muchachos of the
Instituto de Astrofísica de Canarias, on the island of La Palma. The
Guoshoujing Telescope (the Large Sky Area Multi-Object Fiber Spectroscopic
Telescope LAMOST) is a National Major Scientific Project built by the Chinese
Academy of Sciences. Funding for the project has been provided by the National
Development and Reform Commission. LAMOST is operated and managed by the
National Astronomical Observatories, Chinese Academy of Sciences.

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Here we assemble a list of highly r-process-enhanced stars, chosen to have
[Eu/Fe] $\geq+1.0$, sorted by decreasing [Eu/H]. Note that the abundances
listed in the table are LTE results. The [C/Fe] listed is the “as measured”
quantity, not including corrections for evolutionary state. The column labeled
LOC indicates the likely stellar population with which each star is
associated.

Table A1:  Highly r-process-enhanced Stars ([Eu/Fe] $\geq+1.0$ and [Ba/Eu]
$\textless$ 0.0).

Stellar ID | ${T_{\rm eff}}$ | $\rm log\,g$ | $\xi_{\rm t}$ | $\rm[Fe/H]$ | $\rm[Eu/H]$ | $\rm[Eu/Fe]$ | $\rm[Ba/Eu]$ | $\rm[C/Fe]$ | $\rm[Mg/Fe]$ | LOC | REF
---|---|---|---|---|---|---|---|---|---|---|---
| (K) | (cgs) | ($\rm km\,{s^{\rm-1}}$) |  |  |  |  |  |  |  |
LAMOST J020623.21+494127.9 | 4130 | 1.52 | 1.80 |  $-$0.54 | +0.78 | +1.32 |  $-$0.95 |  $-$0.20 | +0.31 | Disk | this work
$\rm{[LDH2014]}$Fnx-mem0556 | 4176 | 0.70 |  |  $-$0.79 | +0.46 | +1.25 |  $-$0.88 | … |  $-$0.15 | Fnx a afootnotemark:  | REI21
$\rm{[LDH2014]}$Fnx-mem0595 | 4223 | 0.66 |  |  $-$0.92 | +0.46 | +1.38 |  $-$0.85 | … |  $-$0.17 | Fnx a afootnotemark:  | REI21
BPS CS 31070-073 | 6190 | 3.86 | 1.50 |  $-$2.55 | +0.28 | +2.83 |  $-$0.41 | +1.34 | +0.64 | Halo | ALL12
2MASS J22132050-5137385 | 5509 | 2.28 | 1.25 |  $-$2.20 | +0.25 | +2.45 |  $-$0.73 | +0.38 | +0.49 | Halo | ROE24
$\rm{[LDH2014]}$Fnx-mem0546 | 4367 | 0.82 |  |  $-$1.33 | +0.12 | +1.45 |  $-$1.05 | … |  $-$0.11 | Fnx a afootnotemark:  | REI21
BPS CS 29526-110 | 6650 | 3.79 | 2.10 |  $-$2.19 | +0.09 | +2.28 |  $-$0.01 | +2.38 | +0.22 | Halo | ALL12
UMi COS 92 | 4325 | 0.30 | 1.95 |  $-$1.45 | +0.04 | +1.49 |  $-$0.72 | … | +0.02 | UMi b bfootnotemark:  | SHE01
HD 222925 | 5636 | 2.54 | 2.20 |  $-$1.47 |  $-$0.14 | +1.33 |  $-$0.78 |  $-$0.20 | +0.41 | Halo | ROE18
2MASS J07202253-3358518 | 5040 | 2.22 | 1.78 |  $-$1.60 |  $-$0.16 | +1.44 |  $-$0.76 | +0.11 | … | Halo | HOL20
LAMOST J112456.61+453531.3 | 5180 | 2.70 | 1.50 |  $-$1.27 |  $-$0.17 | +1.10 |  $-$0.86 |  $-$0.41 |  $-$0.31 | Halo | XIN19
RAVE J040618.2-030525 | 5260 | 2.75 | 1.80 |  $-$1.34 |  $-$0.17 | +1.17 |  $-$0.19 | +0.54 | … | Halo | RAS20
DES J033457.57-540531.4 | 5328 | 2.85 | 1.50 |  $-$2.08 |  $-$0.31 | +1.77 |  $-$0.41 |  |  | Ret II c cfootnotemark:  | JI16
2MASS J07103110-7121522 | 5167 | 2.94 | 1.75 |  $-$1.47 |  $-$0.42 | +1.05 |  $-$0.72 | 0.00 | … | Halo | HOL20
SMSS J175046.30-425506.9 | 4752 | 1.55 | 2.05 |  $-$2.17 |  $-$0.42 | +1.75 |  $-$0.90 |  $-$0.12 | +0.24 | Halo | JAC15
RAVEJ153830.9-180424 | 4995 | 2.00 | 1.85 |  $-$2.02 |  $-$0.50 | +1.52 |  $-$0.80 | +0.12 |  | Halo | RAS20
2MASS J18024226-4404426 | 4701 | 1.60 | 2.17 |  $-$1.55 |  $-$0.50 | +1.05 |  $-$0.10 | +0.35 | … | Halo | HAN18
2MASS J15213995-3538094 | 5850 | 2.10 | 2.65 |  $-$2.80 |  $-$0.57 | +2.23 |  $-$0.88 | +0.56 | +0.37 | Halo | CAI20
SMSS J183128.71-341018.4 | 4940 | 2.15 | 2.00 |  $-$1.83 |  $-$0.58 | +1.25 |  $-$0.72 | +0.01 | +0.22 | Bulge | HOW16
2MASS J19570805+5529491 | 5070 | 1.30 | 2.18 |  $-$1.72 |  $-$0.61 | +1.11 |  $-$1.26 | … | +0.31 | Halo | HAW18
2MASS J07114252-3432368 | 4767 | 1.33 | 1.74 |  $-$1.96 |  $-$0.66 | +1.30 |  $-$0.80 | +0.47 | +0.31 | Halo | SAK18
SMSS J183225.29-334938.4 | 5293 | 2.35 | 1.80 |  $-$1.74 |  $-$0.66 | +1.08 |  $-$0.58 |  $-$0.22 | +0.42 | Bulge | HOW16
2MASS J05241392-0336543 | 4430 | 2.25 | 3.00 |  $-$2.20 |  $-$0.70 | +1.50 |  $-$1.03 |  $-$0.40 |  $-$0.14 | Halo | EZZ20
2MASS J03073894-0502491 | 4360 | 1.75 | 2.90 |  $-$1.98 |  $-$0.71 | +1.27 |  $-$1.13 |  $-$0.57 | +0.09 | Halo | EZZ20
2MASS J13052137-1137220 | 4600 | 1.80 | 2.20 |  $-$1.82 |  $-$0.73 | +1.09 |  $-$0.89 |  $-$0.30 |  $-$0.59 | Halo | EZZ20
2MASS J20093393-3410273 | 4690 | 1.64 | 1.70 |  $-$2.10 |  $-$0.78 | +1.32 |  $-$1.09 |  $-$0.86 | … | Halo | HAN18
2MASS J00101758-1735387 | 5200 | 3.00 | 1.70 |  $-$2.19 |  $-$0.84 | +1.35 |  $-$0.92 | 0.00 |  $-$0.56 | Halo | EZZ20
2MASS J12170829+0415146 | 4540 | 1.65 | 2.90 |  $-$2.00 |  $-$0.90 | +1.10 |  $-$0.95 |  $-$0.40 |  $-$0.18 | Halo | EZZ20
2MASS J23362202-5607498 | 4630 | 1.28 | 2.15 |  $-$2.06 |  $-$0.92 | +1.14 |  $-$0.90 |  $-$0.05 | … | Halo | HAN18
2MASS J00512646-1053170 | 6440 | 4.02 | 1.59 |  $-$2.34 |  $-$0.97 | +1.37 |  $-$0.72 | +0.62 | +0.45 | Halo | SHA24
RAVE J012931.1-160046 | 4959 | 1.70 | 2.65 |  $-$2.77 |  $-$0.97 | +1.80 |  $-$0.60 | +0.44 | +0.47 | Halo | RAS20
2MASS J14534137+0040467 | 4550 | 1.55 | 2.30 |  $-$2.82 |  $-$1.02 | +1.80 |  $-$0.75 |  $-$0.41 | +0.92 | Halo | EZZ20
DES J033537.06-540401.2 | 5170 | 2.45 | 1.55 |  $-$2.73 |  $-$1.03 | +1.70 |  $-$0.30 | … | … | Ret II c cfootnotemark:  | JI16
DES J033447.93-540525.0 | 4900 | 1.70 | 1.90 |  $-$2.91 |  $-$1.04 | +1.87 |  $-$0.79 | … | … | Ret II c cfootnotemark:  | JI16
HE 0048-1109 | 6265 | 3.80 | 1.45 |  $-$2.35 |  $-$1.06 | +1.29 |  $-$0.67 | +0.61 | +0.30 | Halo | GUL21
CS 31082-001 | 4827 | 1.65 | 1.70 |  $-$2.79 |  $-$1.07 | +1.72 |  $-$0.50 | +0.04 | +0.46 | Halo | SAK18
2MASS J03270229+0132322 | 5100 | 2.95 | 1.80 |  $-$2.17 |  $-$1.10 | +1.07 |  $-$0.57 | +0.35 | +0.07 | Halo | EZZ20
HE 0007-1752 | 5215 | 2.65 | 1.45 |  $-$2.36 |  $-$1.11 | +1.25 |  $-$0.65 |  $-$0.29 |  $-$0.33 | Halo | GUL21
2MASS J03422816-6500355 | 4976 | 2.29 | 2.05 |  $-$2.16 |  $-$1.11 | +1.05 |  $-$0.65 | +0.05 | … | Halo | HOL20
2MASS J15260106-0911388 | 4499 | 0.76 | 2.41 |  $-$2.83 |  $-$1.13 | +1.70 |  $-$1.01 |  $-$0.82 | … | Halo | HAN18
DES J033556.28-540316.3 | 5305 | 2.95 | 1.65 |  $-$3.54 |  $-$1.14 |  $\textless$+2.40 |  $-$2.30 | … | … | Ret II c cfootnotemark:  | JI16
2MASS J21091825-1310062 | 4855 | 1.42 | 2.10 |  $-$2.40 |  $-$1.15 | +1.25 |  $-$1.13 |  $-$0.28 | … | Halo | HAN18
J0246-1518 | 4948 | 1.93 | 1.56 |  $-$2.45 |  $-$1.16 | +1.29 |  $-$0.64 |  $-$0.07 | +0.13 | Halo | SAK18
2MASS J14325334-4125494 | 5020 | 2.39 | 1.80 |  $-$2.79 |  $-$1.18 | +1.61 |  $-$0.90 |  $-$0.14 | … | Halo | HAN18
CS 29497-004 | 5013 | 2.23 | 1.62 |  $-$2.81 |  $-$1.19 | +1.62 |  $-$0.41 | +0.22 | +0.31 | Halo | BAR05
HE 0324-0122 | 5145 | 2.45 | 1.65 |  $-$2.41 |  $-$1.20 | +1.21 |  $-$0.75 | +0.37 | +0.36 | Halo | GUL21
2MASS J22182082-3827554 | 5000 | 2.70 | 1.60 |  $-$2.32 |  $-$1.22 | +1.10 |  $-$0.60 | +0.30 | +0.18 | Halo | EZZ20
DES J033607.74-540235.5 | 4833 | 1.55 | 2.15 |  $-$2.97 |  $-$1.23 | +1.74 |  $-$0.83 | … | … | Ret II c cfootnotemark:  | JI16
2MASS J05383296-5904280 | 5824 | 2.03 | 2.84 |  $-$2.53 |  $-$1.25 | +1.28 |  $-$0.52 |  $\textless$+0.80 | … | Halo | HOL20
2MASS J02462013-1518419 | 4879 | 1.80 | 2.35 |  $-$2.71 |  $-$1.26 | +1.45 |  $-$0.85 | +0.04 | … | Halo | HAN18
HE 1523-0901 | 4290 | 0.20 | 2.13 |  $-$3.09 |  $-$1.27 | +1.82 |  $-$2.37 | +0.39 | +0.47 | Halo | SAK18
2MASS J19161821-5544454 | 4450 | 0.65 | 2.50 |  $-$2.35 |  $-$1.27 | +1.08 |  $-$1.20 |  $-$0.80 | … | Halo | HAN18
RAVE J203843.2–002333 | 4630 | 1.20 | 2.15 |  $-$2.91 |  $-$1.27 | +1.64 |  $-$0.81 |  $-$0.44 | +0.36 | Halo | PLA17
DES J033523.85-540407.5 | 4608 | 1.00 | 2.40 |  $-$3.01 |  $-$1.33 | +1.68 |  $-$0.89 | … | … | Ret II c cfootnotemark:  | JI16
2MASS J17225742-7123000 | 5080 | 2.67 | 0.70 |  $-$2.42 |  $-$1.35 | +1.07 |  $-$0.37 |  $-$0.33 | … | Halo | HAN18
M15 31914 | 4338 | 0.58 | 2.20 |  $-$2.41 |  $-$1.35 | +1.06 |  $-$0.70 | … | … | M15 d dfootnotemark:  | GAR24
HE 1226-1149 | 5120 | 2.35 | 1.50 |  $-$2.91 |  $-$1.36 | +1.55 |  $-$0.65 | +0.42 | +0.32 | Halo | COH13
2MASS J02165716-7547064 | 4543 | 0.71 | 2.44 |  $-$2.50 |  $-$1.38 | +1.12 |  $-$0.87 |  $-$0.33 |  | Halo | HAN18
2MASS J21064294-6828266 | 5186 | 2.70 | 2.90 |  $-$2.76 |  $-$1.44 | +1.32 |  $-$0.80 | +0.53 | … | Halo | HAN18
SMSS J175738.37-454823.5 | 4617 | 1.20 | 2.30 |  $-$2.46 |  $-$1.44 | +1.02 |  $-$3.01 |  $-$0.43 | +0.39 | Halo | JAC15
SDSS J235718.91-005247.8 | 5000 | 4.80 | 0.00 |  $-$3.36 |  $-$1.44 | +1.92 |  $-$0.80 | +0.43 | +0.19 | Halo | AOK10
BPS CS 22892-052 | 4800 | 1.50 | 1.95 |  $-$3.10 |  $-$1.46 | +1.64 |  $-$0.65 | +0.88 | +0.30 | Halo | SNE03
SMSS J155430.57-263904.8 | 4783 | 1.20 | 2.20 |  $-$2.61 |  $-$1.47 | +1.14 |  $-$0.34 |  $-$0.26 | +0.33 | Halo | JAC15
HE 2224+0143 | 5198 | 2.66 | 1.67 |  $-$2.58 |  $-$1.53 | +1.05 |  $-$0.46 | +0.35 | +0.32 | Halo | BAR05
HE 0430-4901 | 5296 | 3.12 | 1.32 |  $-$2.72 |  $-$1.56 | +1.16 |  $-$0.66 | +0.09 | +0.17 | Halo | BAR05
HE 1219-0312 | 5060 | 2.30 | 1.60 |  $-$2.96 |  $-$1.58 | +1.38 |  $-$0.73 | +0.03 | +0.41 | Halo | HAY09
BPS CS 22958-052 | 6090 | 3.75 | 1.95 |  $-$2.62 |  $-$1.62 | +1.00 |  $-$1.00 | +0.32 | … | Halo | ROE14
BPS CS 31078-018 | 5257 | 2.75 | 1.50 |  $-$2.85 |  $-$1.62 | +1.23 |  $-$0.51 | +0.37 | +0.42 | Halo | LAI08
HE 1127-1143 | 5224 | 2.64 | 1.59 |  $-$2.73 |  $-$1.65 | +1.08 |  $-$0.45 | +0.54 | +0.22 | Halo | BAR05
SMSS J062609.83-590503.2 | 4960 | 1.75 | 1.75 |  $-$2.77 |  $-$1.71 | +1.06 |  $-$0.22 | +0.27 | +0.24 | Halo | JAC15
HE 2327-5642 | 5048 | 2.22 | 1.69 |  $-$2.95 |  $-$1.73 | +1.22 |  $-$0.56 | +0.43 | +0.14 | Halo | BAR05
CS 29491-069 | 5103 | 2.45 | 1.54 |  $-$2.81 |  $-$1.75 | +1.06 |  $-$0.72 | +0.18 | +0.28 | Halo | BAR05
BPS CS 22945-017 | 6080 | 3.70 | 1.25 |  $-$2.89 |  $-$1.76 | +1.13 |  $-$0.64 | +1.78 | … | Halo | ROE14
DES J033531.13-540148.2 | 4925 | 1.90 | 1.80 |  $-$3.34 |  $-$1.84 |  $\textless$+1.50 |  $-$2.30 | … | … | Ret II c cfootnotemark:  | JI16
BPS CS 22945-058 | 5990 | 3.65 | 1.55 |  $-$2.98 |  $-$1.84 | +1.14 |  $-$0.85 | +0.68 | … | Halo | ROE14
HE 0432-0923 | 5131 | 2.64 | 1.54 |  $-$3.19 |  $-$1.94 | +1.25 |  $-$0.53 | +0.24 | +0.34 | Halo | BAR05
BPS BS 16929-005 | 5250 | 3.10 | 2.30 |  $-$3.15 |  $-$2.14 | +1.01 |  $-$0.74 | +1.50 | +0.49 | Halo | ALL12
SPLUS J142445.34-254247.1 | 4762 | 1.58 | 1.60 |  $-$3.82 |  $-$2.20 | +1.62 |  $-$0.37 |  $-$0.21 | +0.54 | Halo | PLA23
LAMOST J110901.22+075441.8 | 4440 | 0.70 | 1.98 |  $-$3.41 |  $-$2.25 | +1.16 |  $-$0.85 |  $-$0.57 | +0.41 | Halo | LIH15
BPS CS 22885-096 | 5050 | 2.00 | 1.80 |  $-$3.60 |  $-$2.60 | +1.00 |  $-$1.74 | … | +0.26 | Halo | RYA96
SMSS J024858.41-684306.4 | 4977 | 1.60 | 1.80 |  $-$3.71 |  $-$2.71 | +1.00 |  $-$0.41 | +0.66 | +0.57 | Halo | JAC15

aafootnotetext: The Fornax dwarf galaxy.

bbfootnotetext: The Ursa Minor dwarf galaxy.

ccfootnotetext: The Reticulum II dwarf galaxy.

ddfootnotetext: The M15 globular cluster.

Note. — RYA96: Ryan et al. (1996); SHE01: Shetrone et al. (2001); SNE03:
Sneden et al. (2003); BAR05: Barklem et al. (2005); LAI08: Lai et al. (2008);
HAY09: Hayek et al. (2009); AOK10: Aoki et al. (2010); ALL12: Allen et al.
(2012); COH13: Cohen et al. (2013); ROE14: Roederer et al. (2014a); JAC15:
Jacobson et al. (2015); LIH15: Li et al. (2015); HOW16: Howes et al. (2016);
JI16: Ji et al. (2016); PLA17: Placco et al. (2017); HAN18: Hansen et al.
(2018); HAW18: Hawkins & Wyse (2018); ROE18: Roederer et al. (2018b); SAK18:
Sakari et al. (2018); XIN19: Xing et al. (2019); CAI20: Cain et al. (2020);
EZZ20: Ezzeddine et al. (2020); HOL20: Holmbeck et al. (2020); RAS20:
Rasmussen et al. (2020); GUL21: Gull et al. (2021); REI21: Reichert et al.
(2021); PLA23: Placco et al. (2023); GAR24: Cabrera Garcia et al. (2024);
ROE24: Roederer et al. (2024); SHA24: Shah et al. (2024).