# Chemical Analysis of the Brightest Star of the Cetus II Ultra-Faint Dwarf
Galaxy Candidate111This paper includes data gathered with the 6.5m Magellan
Telescopes located at Las Campanas Observatory, Chile.

K. B. Webber Mitchell Institute for Fundamental Physics and Astronomy and
Department of Physics and Astronomy, Texas A&M University, College Station, TX
77843-4242, USA  T. T. Hansen Department of Astronomy, Stockholm University,
AlbaNova University Center, SE-106 91 Stockholm, Sweden  J. L. Marshall
Mitchell Institute for Fundamental Physics and Astronomy and Department of
Physics and Astronomy, Texas A&M University, College Station, TX 77843-4242,
USA  J. D. Simon  Observatories of the Carnegie Institution for Science, 813
Santa Barbara St., Pasadena, CA 91101, USA  A. B. Pace  McWilliams Center for
Cosmology, Carnegie Mellon University, 5000 Forbes Ave, Pittsburgh, PA 15213,
USA  B. Mutlu-Pakdil  Department of Physics and Astronomy, Dartmouth College,
Hanover, NH 03755, USA  A. Drlica-Wagner  Fermi National Accelerator
Laboratory, P. O. Box 500, Batavia, IL 60510, USA  Kavli Institute for
Cosmological Physics, University of Chicago, Chicago, IL 60637, USA
Department of Astronomy and Astrophysics, University of Chicago, Chicago, IL
60637, USA  C. E. Martínez-Vázquez  Gemini Observatory, NSF’s NOIRLab, 670 N.
A’ohoku Place, Hilo, HI 96720, USA  Cerro Tololo Inter–American Observatory,
NSF’s NOIRLab, Casilla 603, La Serena, Chile  M. Aguena  Laboratório
Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, Rio de
Janeiro, RJ - 20921-400, Brazil  S. S. Allam Fermi National Accelerator
Laboratory, P. O. Box 500, Batavia, IL 60510, USA  O. Alves  Department of
Physics, University of Michigan, Ann Arbor, MI 48109, USA  E. Bertin  CNRS,
UMR 7095, Institut d’Astrophysique de Paris, F-75014, Paris, France  Sorbonne
Universités, UPMC Univ Paris 06, UMR 7095, Institut d’Astrophysique de Paris,
F-75014, Paris, France  D. Brooks Department of Physics & Astronomy,
University College London, Gower Street, London, WC1E 6BT, UK  A. Carnero
Rosell Instituto de Astrofisica de Canarias, E-38205 La Laguna, Tenerife,
Spain  Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José
Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil  Universidad de La Laguna,
Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain  J. Carretero Institut
de Física d’Altes Energies (IFAE), The Barcelona Institute of Science and
Technology, Campus UAB, 08193 Bellaterra (Barcelona) Spain  L. N. da Costa
Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino
77, Rio de Janeiro, RJ - 20921-400, Brazil  J. De Vicente Centro de
Investigaciones Energéticas, Medioambientales y Tecnológicas (CIEMAT), Madrid,
Spain  P. Doel  Department of Physics & Astronomy, University College London,
Gower Street, London, WC1E 6BT, UK  I. Ferrero  Institute of Theoretical
Astrophysics, University of Oslo. P.O. Box 1029 Blindern, NO-0315 Oslo, Norway
D. Friedel  Center for Astrophysical Surveys, National Center for
Supercomputing Applications, 1205 West Clark St., Urbana, IL 61801, USA  J.
Frieman Fermi National Accelerator Laboratory, P. O. Box 500, Batavia, IL
60510, USA  Kavli Institute for Cosmological Physics, University of Chicago,
Chicago, IL 60637, USA  J. García-Bellido Instituto de Fisica Teorica
UAM/CSIC, Universidad Autonoma de Madrid, 28049 Madrid, Spain  G. Giannini
Institut de Física d’Altes Energies (IFAE), The Barcelona Institute of Science
and Technology, Campus UAB, 08193 Bellaterra (Barcelona) Spain  D. Gruen
University Observatory, Faculty of Physics, Ludwig-Maximilians-Universität,
Scheinerstr. 1, 81679 Munich, Germany  R. A. Gruendl  Center for Astrophysical
Surveys, National Center for Supercomputing Applications, 1205 West Clark St.,
Urbana, IL 61801, USA  Department of Astronomy, University of Illinois at
Urbana-Champaign, 1002 W. Green Street, Urbana, IL 61801, USA  S. R. Hinton
School of Mathematics and Physics, University of Queensland, Brisbane, QLD
4072, Australia  D. L. Hollowood  Santa Cruz Institute for Particle Physics,
Santa Cruz, CA 95064, USA  K. Honscheid Center for Cosmology and Astro-
Particle Physics, The Ohio State University, Columbus, OH 43210, USA
Department of Physics, The Ohio State University, Columbus, OH 43210, USA  K.
Kuehn Australian Astronomical Optics, Macquarie University, North Ryde, NSW
2113, Australia  Lowell Observatory, 1400 Mars Hill Rd, Flagstaff, AZ 86001,
USA  J. Mena-Fernández Centro de Investigaciones Energéticas, Medioambientales
y Tecnológicas (CIEMAT), Madrid, Spain  F. Menanteau Center for Astrophysical
Surveys, National Center for Supercomputing Applications, 1205 West Clark St.,
Urbana, IL 61801, USA  Department of Astronomy, University of Illinois at
Urbana-Champaign, 1002 W. Green Street, Urbana, IL 61801, USA  R. Miquel
Institució Catalana de Recerca i Estudis Avançats, E-08010 Barcelona, Spain
Institut de Física d’Altes Energies (IFAE), The Barcelona Institute of Science
and Technology, Campus UAB, 08193 Bellaterra (Barcelona) Spain  R. L. C.
Ogando Observatório Nacional, Rua Gal. José Cristino 77, Rio de Janeiro, RJ -
20921-400, Brazil  M. E. S. Pereira  Hamburger Sternwarte, Universität
Hamburg, Gojenbergsweg 112, 21029 Hamburg, Germany  A. Pieres Laboratório
Interinstitucional de e-Astronomia - LIneA, Rua Gal. José Cristino 77, Rio de
Janeiro, RJ - 20921-400, Brazil  Observatório Nacional, Rua Gal. José Cristino
77, Rio de Janeiro, RJ - 20921-400, Brazil  A. A. Plazas Malagón Kavli
Institute for Particle Astrophysics & Cosmology, P. O. Box 2450, Stanford
University, Stanford, CA 94305, USA  SLAC National Accelerator Laboratory,
Menlo Park, CA 94025, USA  E. Sanchez Centro de Investigaciones Energéticas,
Medioambientales y Tecnológicas (CIEMAT), Madrid, Spain  B. Santiago
Instituto de Física, UFRGS, Caixa Postal 15051, Porto Alegre, RS - 91501-970,
Brazil  Laboratório Interinstitucional de e-Astronomia - LIneA, Rua Gal. José
Cristino 77, Rio de Janeiro, RJ - 20921-400, Brazil  J. Allyn Smith Austin
Peay State University, Dept. Physics, Engineering and Astronomy, P.O. Box 4608
Clarksville, TN 37044, USA  M. Smith School of Physics and Astronomy,
University of Southampton, Southampton, SO17 1BJ, UK  E. Suchyta Computer
Science and Mathematics Division, Oak Ridge National Laboratory, Oak Ridge, TN
37831  G. Tarle Department of Physics, University of Michigan, Ann Arbor, MI
48109, USA  C. To Center for Cosmology and Astro-Particle Physics, The Ohio
State University, Columbus, OH 43210, USA  N. Weaverdyck  Department of
Physics, University of Michigan, Ann Arbor, MI 48109, USA  Lawrence Berkeley
National Laboratory, 1 Cyclotron Road, Berkeley, CA 94720, USA  B. Yanny
Fermi National Accelerator Laboratory, P. O. Box 500, Batavia, IL 60510, USA
K. B. Webber<EMAIL_ADDRESS>
###### Abstract

We present a detailed chemical abundance analysis of the brightest star in the
ultra-faint dwarf (UFD) galaxy candidate Cetus II from high-resolution
Magellan/MIKE spectra. For this star, DES J011740.53-173053, abundances or
upper limits of 18 elements from Carbon to Europium are derived. Its chemical
abundances generally follow those of other UFD galaxy stars, with a slight
enhancement of the $\alpha$-elements (Mg, Si, and Ca) and low neutron-capture
element (Sr, Ba, Eu) abundances supporting the classification of Cetus II as a
likely UFD. The star exhibits lower Sc, Ti, and V abundances than Milky Way
(MW) halo stars with similar metallicity. This signature is consistent with
yields from a supernova (SN) originating from a star with a mass of $\sim$11.2
M⊙. In addition, the star has a Potassium abundance of $\mathrm{[K/Fe]}=0.81$
which is somewhat higher than the K abundances of MW halo stars with similar
metallicity, a signature which is also present in a number of UFD galaxies. A
comparison including globular clusters (GC) and stellar stream stars suggests
that high K is a specific characteristic for some UFD galaxy stars and can
thus be used to help classify objects as UFD galaxies.

††software: MOOG (Sneden, 1973; Sobeck et al., 2011), IRAF (Tody, 1986, 1993),
ATLAS9 (Castelli & Kurucz, 2003), linemake (Placco et al., 2021), NumPy (van
der Walt et al., 2011), Matplotlib (Hunter, 2007), AstroPy (Astropy
Collaboration et al., 2013, 2018), CarPy (Kelson, 2003), SMHR (Casey, 2014)

DES-2023-0757 FERMILAB-PUB-23-373-PPD

##  1 Introduction

Stars maintain in their atmospheres a fingerprint of the chemical composition
of their birth environment and thereby also contain information about the
stellar generations that were present before them. They can therefore be used
to study chemical enrichment, which is a useful probe to characterize the
nature of the stellar environment. Variations in chemical abundances across
different stellar associations can give insight into properties such as
different timescales for chemical enrichment and variations in stellar initial
mass functions.

In recent years, photometric surveys such as the Dark Energy Survey and Pan-
STARRS (DES Collaboration, 2005; Bechtol et al., 2015; Koposov et al., 2015;
Drlica-Wagner et al., 2015; Laevens et al., 2015) have detected a large number
of stellar overdensities around the Milky Way (MW). In the process of
characterizing these systems, chemical abundance analysis has proven to be a
valuable tool to discern the physical nature of the associations and sometimes
reveal unusual chemical abundance patterns. Although further observations of a
few of these overdensities have determined that they are not genuine physical
objects (Cantu et al., 2021), nearly all of the overdensities have been found
to be either star clusters or ultra-faint dwarf (UFD) galaxies (Laevens et
al., 2014; Luque et al., 2018; Simon, 2019).

Star clusters can be sorted into different categories: stellar associations,
open clusters, or GCs (Trumpler, 1930; Krumholz et al., 2019). Each of these
types of systems has distinguishing features in the abundances of their member
stars that can be used to identify the type of cluster. Open clusters are
homogeneous in the abundances of all elements, while GCs are not fully
homogeneous (Krumholz et al., 2019). GCs also are known to display specific
anti-correlations between certain elements such as Mg-Al and Na-O (Carretta et
al., 2010; Mucciarelli et al., 2018). Hence, studying the chemical enrichment
of star clusters can tell us useful information about the nature of these
objects and aid in their classification.

UFD galaxies are low mass, low luminosity, dark matter dominated galaxies
(Simon & Geha, 2007). The current stellar population of UFD galaxies is very
old ($\sim$10 Gyr) and metal-poor ($\mathrm{[Fe/H]}<-1.4$, Simon 2019), with
abundances reflecting the elements created by the first generation of stars.
Studying the abundances of stars in UFD galaxies, therefore, provides a window
to study the nucleosynthetic processes of the early Universe. Furthermore,
with the UFD galaxies being small, isolated systems, they reflect the chemical
signatures of few nucleosynthetic events (Ji et al., 2016c; Hansen et al.,
2020, 2017; Marshall et al., 2019).

Previous studies of stellar abundances in UFD galaxies have found generally
similar abundance patterns from system to system (Frebel & Norris, 2015), with
most of the stars exhibiting a slight enhancement of the $\alpha$-elements
(Mg, Si, Ca) (Simon, 2019; Frebel et al., 2010; Lai et al., 2011), as seen for
metal-poor stars in the MW halo (e.g., McWilliam, 1998; Cayrel et al., 2004),
reflecting enrichment by core-collapse SN. In addition, and specific to UFD
galaxies, most systems also display very low neutron-capture element
abundances (e.g., Ji et al., 2019). With the increase in the number of systems
being discovered and analysed, more variation in the abundances has begun to
appear. Examples of systems with unusual abundance patterns include Reticulum
II, where $\sim 72\%$ of stars analysed show enhancement in neutron-capture
elements likely due to a neutron star merger event occurring early in the
history of the galaxy (Ji et al., 2022, 2016; Roederer et al., 2016), and Grus
II, where all three stars analysed have a high $\mathrm{[Mg/Ca]}$ ratio
compared to other UFD galaxies, pointing to a top heavy initial mass function
governing the early star formation in this galaxy (Hansen et al., 2020).

In this paper, we present a detailed chemical abundance analysis of the
brightest star of the Cetus II (Cet II) UFD galaxy candidate, DES
J011740.53-173053.1 (hereafter J0117). Cet II was discovered as a stellar
overdensity in the Dark Energy Survey, having a heliocentric distance of 30
kpc (Drlica-Wagner et al., 2015). At the time of discovery, Cet II was the
faintest and smallest candidate UFD galaxy system detected, making it
difficult to determine the velocity and metallicity dispersion. The brightest
member was, therefore, a prime target for high-resolution follow-up aimed at
abundance analysis. As described above, both UFD galaxies and star clusters
have distinct chemical features that can be used to classify the system. Hence
a detailed chemical analysis is the natural next step in the efforts to
characterize the Cet II system.

The outline of the paper is as follows. In Section 2, the observations are
described, and in Section 3, the stellar parameters and analysis are detailed.
Section 4 presents our results that are further discussed in Section 5.
Section 6 provides a summary.

##  2 Observations

Based on medium-resolution spectroscopy of the Cet II field obtained with the
Magellan/IMACS spectrograph (Dressler et al., 2011) in 2016, we identified a
set of likely Cet II member stars centered at a velocity of $V_{\rm
helio}=-82$ km s-1. Subsequently, Pace & Li (2019) and Pace et al. (2022)
showed that these stars also share a common proper motion, confirming that
they are associated with Cet II. Surprisingly, given the low stellar mass of
Cet II, one of the member stars is located on the upper red giant branch at an
apparent magnitude of $g=16.44$, $\sim 3$ mag brighter than any other Cet II
stars. This bright star, J0117, was then an obvious target for high-resolution
spectroscopy to investigate the chemical abundances in Cet II.

High-resolution spectral data was obtained for J0117 with the MIKE echelle
spectrograph at Las Campanas Observatory in Chile (Bernstein et al., 2003) in
August and November 2017. A color-magnitude diagram of Cet II member stars
(Pace et al., 2022) is shown in Figure 1. J0117 is marked with a red dot and
is notably brighter than the other currently known members.

Right ascension, declination, and dereddened DES magnitudes for J0117 are
listed in Table 1, along with Heliocentric Julian Date (HJD), exposure times,
signal-to-noise ratio per pixel (S/N), and heliocentric radial velocities for
the spectra. The spectra were obtained using a 0.7” slit with 2x2 pixel
binning and cover a wavelength range of 3350-5000 Å in the blue and 4900-9500
Å in the red, with resolutions ($R$=$\lambda/\Delta\lambda$) of $R\sim$ 35,000
at blue wavelengths and $\sim 28,000$ at red wavelengths, respectively. The
data from each observing run were reduced using the CarPy MIKE pipeline
(Kelson et al., 2000; Kelson, 2003), and the spectra from the two runs were
subsequently co-added. The heliocentric radial velocity of the star was
determined by cross-correlation with a spectrum of the bright, metal-poor red
giant HD122563 ($V_{helio}$ = $-$26.17 km s-1; Gaia Collaboration et al.
2021). Thirty-six echelle orders were used for the correlation in each
spectrum, yielding mean radial velocities of $-$81.38 km s-1 and $-$81.89 km
s-1 for the two spectra. Since there is no appreciable variation in the radial
velocities between the two spectra, this star is likely not a short-period
binary.

Figure 1:  Color-magnitude diagram for Cetus II. Black dots are stars with
membership probability $p>0.01$ (Pace et al., 2022) and red dot is J0117. Blue
line shows a Dartmouth isochrone (Dotter et al., 2008) with
$\mathrm{[Fe/H]}=-2.3$ and age = 12.5 Gyr, transformed to the DES photometric
system (Drlica-Wagner et al., 2018) and shifted to the distance of Cetus II
($\sim$30 kpc, Drlica-Wagner et al. 2015).  Table 1: Observing Details and
Stellar Data

ID | RA | Dec | $g_{0}$ | $r_{0}$ | $i_{0}$ | $z_{0}$ | HJD | Exp Time | S/N @ 4500Å |  $V_{helio}$ $\pm$ $\sigma$ |
---|---|---|---|---|---|---|---|---|---|---|---
|  |  | (mag) | (mag) | (mag) | (mag) |  | sec |  | km s-1 |
J0117 | 01:17:40.9 | $-$17:30:54.5 | 16.443 | 15.716 | 15.818 | 16.014 | 2457988.70127 | 2x1800 | 15 |  $-$81.38 $\pm$ 0.52 |
2458059.49903 | 5x1800 | 21 |  $-$81.89 $\pm$ 0.51 |

Note. — HJD values at beginning of exposure

Table 2:  Data for atomic lines used in analysis and individual line EW
measurements.

Species | $\lambda$ | $\chi$ | $\log{gf}$ | EW | $\sigma_{\rm EW}$ | $\log{\epsilon}$ | ref |
---|---|---|---|---|---|---|---|---
| (Å) | (eV) |  | (mÅ) | (mÅ) |  |  |
Na I | 5889.95 | 0.00 | 0.11 | 197.33 | 2.83 | 4.04 | 1 |
Na I | 5895.92 | 0.00 |  $-$0.19 | 171.34 | 2.89 | 3.99 | 1 |
Mg I | 4167.27 | 4.35 |  $-$0.74 | 84.82 | 6.00 | 5.80 | 1 |
Mg I | 4702.99 | 4.33 |  $-$0.44 | 95.81 | 2.96 | 5.54 | 1 |
Mg I | 5172.68 | 2.71 |  $-$0.36 | 248.53 | 6.18 | 5.47 | 2 |

References. — (1) Kramida et al. (2018); (2) Pehlivan Rhodin et al. (2017);
(3) Yu & Derevianko (2018); (4) Lawler & Dakin (1989), using hfs from Kurucz &
Bell (1995); (5) Lawler et al. (2013); (6) = Wood et al. (2013); (7) Pickering
et al. (2001), with corrections given in Pickering et al. (2002); (8) Lawler
et al. (2014) for $\log{gf}$ values and HFS; (9) Wood et al. (2014a) for
$\log{gf}$ values and HFS, when available; (10) Sobeck et al. (2007); (11)
Lawler et al. (2017); (12) Den Hartog et al. (2011) for both $\log{gf}$ value
and hfs; (13) O’Brian et al. (1991); (14) Den Hartog et al. (2014); (15)
Belmonte et al. (2017); (16) Ruffoni et al. (2014); (17) Den Hartog et al.
(2019); (18) Meléndez & Barbuy (2009); (19) Lawler et al. (2015) for
$\log{gf}$ values and HFS; (20) Wood et al. (2014b); (21) Roederer & Lawler
(2012); (22) Kramida et al. (2018), using HFS/IS from McWilliam (1998) when
available.

Note. — The complete version of this Table is available online only. A short
version is shown here to illustrate its form and content.

##  3 Stellar Parameters and Abundance Analysis

Stellar parameters and abundances were derived from equivalent width (EW)
measurements and spectral synthesis using the program Spectroscopy Made
Hard(er) (SMHR 222https://github.com/andycasey/smhr), which runs the 2017
version of the radiative transfer code MOOG
333https://github.com/alexji/moog17scat (Sneden, 1973; Sobeck et al., 2011)
assuming local thermodynamical equilibrium (LTE). One dimensional (1D)
$\alpha$-enhanced ([$\alpha$/Fe]= +0.4) ATLAS model atmospheres (Castelli &
Kurucz, 2003) were used as input, with line lists generated from linemake
which include hyperfine structure and isotopic shifts where applicable
444https://github.com/vmplacco/linemake (Placco et al., 2021) and Solar
abundances were taken from Asplund et al. (2009). For the derivation of Ba, we
used the $r$-process isotopic ratio from Sneden et al. (2008).

The stellar parameters effective temperature ($T_{\mathrm{eff}}$), surface
gravity ($\log{g}$), metallicity ($\mathrm{[Fe/H]}$), and microturbulence
($\xi$) were determined spectroscopically using EW measurements of 95 Fe I and
7 Fe II lines (see Table 2). The EWs of the Fe I and Fe II lines were measured
by fitting Gaussian profiles to the absorption features in the continuum-
normalized spectra. Using these measurements, $T_{\mathrm{eff}}$ was derived
from the excitation equilibrium of the Fe I lines and then corrected for the
offset between spectroscopic and photometric temperature scales using the
method outlined in Frebel et al. (2013). Next, $\log{g}$ was determined from
the ionization balance between the Fe I and Fe II lines, and $\xi$ was
determined by removing any trend in line abundances with reduced equivalent
width for the Fe I lines. Final stellar parameters are listed in Table 3. For
comparison, a photometric temperature of 4592 $\pm$ 69 K for the star was also
derived by converting the de-reddened DES $g$, $r$, $i$, and $z$ colors to
$B$, $V$, $R$, and $I$ colors (Drlica-Wagner et al., 2018, R. Lupton
2005555http://www.sdss3.org/dr8/algorithms/sdssUBVRITransform.php) and using
the color-temperature relations from Casagrande et al. (2010). The photometric
temperature is in good agreement with the corrected spectroscopic temperature.
Following Frebel et al. (2013), we adopt a 150 K systematic uncertainty for
$T_{\rm eff}$, corresponding to an uncertainty of 0.3 dex in $\log{g}$, 0.2 km
s-1 in $\xi$, and 0.18 dex in $\mathrm{[Fe/H]}$. The statistical uncertainties
on $T_{\rm eff}$, $\log{g}$ and $\xi$ were derived by varying each parameter
to match the standard deviation of Fe I lines as listed in Table 3. Once the
stellar parameters were determined, elemental abundances were derived from EW
measurements or spectral synthesis. EWs were used for lines that are not
blended, while spectral synthesis was used for blended lines and/or lines that
are affected by isotopic and/or hyperfine splitting. Atomic data, wavelength,
excitation potential, and oscillator strength for individual lines used for
the abundance determination are listed in Table 2. The table also lists the
measured EWs, uncertainties on these, and the corresponding abundances for the
lines used. Final weighted mean abundances and associated uncertainties were
determined following the method outlined in Ji et al. (2020). This method uses
a mean that is weighted by the S/N of the individual lines to calculate the
final abundances and fully propagates statistical and systematic stellar
parameter uncertainties for individual line measurements, including stellar
parameter covariances, to determine the uncertainty.

Table 3: Stellar Model Atmosphere Parameters of J0117 | $T_{\rm eff}$ | $\log g$ | $\xi$ | $\mathrm{[Fe/H]}$
---|---|---|---|---
| (K) | (cgs) | (km s-1) |
Value | 4727 | 1.40 | 1.89 |  $-$2.09
Statistical uncertainties | 41 | 0.07 | 0.07 | 0.17
Systematic uncertainties | 150 | 0.3 | 0.2 | 0.12

Table 4: Weighted Average Abundance Summary for J0117

El. | N | $\log\epsilon$ | $\mathrm{[X/H]}$ | $\sigma_{\mathrm{[X/H]}}$ | $\mathrm{[X/Fe]}$ | $\sigma_{\mathrm{[X/Fe]}}$ | $\Delta_{T_{\rm eff}}$ | $\Delta_{\log g}$ | $\Delta_{\xi}$ | $\Delta_{\mathrm{[Fe/H]}}$ | $s_{X}$
---|---|---|---|---|---|---|---|---|---|---|---
CH | 4 |  $+$5.85 |  $-$2.61 | 0.18 |  $-$0.31 | 0.18 | 0.29 |  $-$0.08 | 0.01 | 0.10 | 0.10
Na I | 2 |  $+$4.02 |  $-$2.12 | 0.26 |  $+$0.17 | 0.25 | 0.26 |  $-$0.08 |  $-$0.14 |  $-$0.03 | 0.00
Mg I | 4 |  $+$5.69 |  $-$1.93 | 0.17 |  $+$0.36 | 0.18 | 0.15 |  $-$0.07 |  $-$0.07 | 0.00 | 0.18
Al I | 2 |  $+$3.79 |  $-$2.63 | 0.42 |  $-$0.33 | 0.41 | 0.25 |  $-$0.11 |  $-$0.11 |  $-$0.01 | 0.46
Si I | 3 |  $+$5.86 |  $-$1.71 | 0.24 |  $+$0.59 | 0.24 | 0.18 |  $-$0.03 |  $-$0.06 |  $-$0.00 | 0.28
K I | 2 |  $+$3.50 |  $-$1.49 | 0.26 |  $+$0.81 | 0.25 | 0.19 |  $-$0.03 |  $-$0.09 |  $-$0.02 | 0.09
Ca I | 20 |  $+$4.38 |  $-$2.03 | 0.10 |  $+$0.23 | 0.11 | 0.15 |  $-$0.03 |  $-$0.06 |  $-$0.01 | 0.11
Sc II | 6 |  $+$0.61 |  $-$2.61 | 0.19 |  $-$0.31 | 0.19 | 0.04 | 0.09 |  $-$0.07 | 0.03 | 0.10
Ti I | 14 |  $+$2.64 |  $-$2.52 | 0.23 |  $-$0.23 | 0.22 | 0.25 |  $-$0.03 |  $-$0.05 |  $-$0.02 | 0.23
Ti II | 23 |  $+$2.93 |  $-$2.23 | 0.17 |  $+$0.06 | 0.18 | 0.06 | 0.10 |  $-$0.11 | 0.04 | 0.28
V I | 2 |  $+$1.64 |  $-$2.30 | 0.19 |  $-$0.01 | 0.18 | 0.17 |  $-$0.02 |  $-$0.01 |  $-$0.04 | 0.00
V II | 3 |  $+$1.98 |  $-$1.93 | 0.17 |  $+$0.37 | 0.18 |  $-$0.02 | 0.12 |  $-$0.02 | 0.03 | 0.00
Cr I | 9 |  $+$3.13 |  $-$2.59 | 0.14 |  $-$0.30 | 0.14 | 0.24 |  $-$0.03 |  $-$0.08 |  $-$0.02 | 0.08
Cr II | 2 |  $+$3.29 |  $-$2.36 | 0.12 |  $-$0.07 | 0.13 |  $-$0.03 | 0.11 |  $-$0.02 | 0.02 | 0.00
Mn I | 5 |  $+$2.97 |  $-$2.48 | 0.23 |  $-$0.18 | 0.22 | 0.20 |  $-$0.02 |  $-$0.08 |  $-$0.04 | 0.10
Fe I | 95 |  $+$5.21 |  $-$2.29 | 0.08 |  $\cdots$ |  $\cdots$ | 0.07 |  $-$0.01 | 0.04 | 0.01 | 0.24
Fe II | 7 |  $+$5.39 |  $-$2.22 | 0.15 |  $\cdots$ |  $\cdots$ |  $-$0.01 | 0.12 | 0.01 | 0.03 | 0.10
Co I | 4 |  $+$2.73 |  $-$2.37 | 0.32 |  $-$0.07 | 0.31 | 0.25 | 0.00 |  $-$0.11 |  $-$0.04 | 0.00
Ni I | 12 |  $+$4.12 |  $-$2.19 | 0.13 |  $+$0.09 | 0.13 | 0.17 |  $-$0.01 |  $-$0.04 |  $-$0.01 | 0.17
Zn I | 1 |  $+$2.27 |  $-$2.29 | 0.15 |  $+$0.01 | 0.15 | 0.06 | 0.07 |  $-$0.02 | 0.02 | 0.00
Sr II | 1 |  $-$1.52 |  $-$4.39 | 0.43 |  $-$2.10 | 0.43 | 0.07 | 0.12 |  $-$0.18 | 0.01 | 0.00
Ba II | 1 |  $-$2.34 |  $-$4.52 | 0.39 |  $-$2.23 | 0.39 | 0.09 | 0.10 | 0.01 | 0.09 | 0.10
Eu II | 1 |  $<-$1.92 |  $<-$2.44 |  $\cdots$ |  $<-$0.33 |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$ |  $\cdots$

##  4 Results

Abundances or upper limits have been derived for 18 elements from C to Eu in
J0117. The LTE abundances and upper limits, along with the systematic
abundance uncertainty ($s_{X}$) and uncertainties arising from stellar
parameter uncertainties ($\Delta_{{\mathrm{T}}_{eff}}$, $\Delta_{\log{g}}$,
$\Delta_{\xi}$, and $\Delta_{[Fe/H]}$), are presented in Table 3, where N is
the number of lines for the given species. In Figure 2 and 3, we present a
subset of elemental abundances for J0117 compared to abundances for stars in
UFD galaxies (colored points) and metal-poor halo stars (grey points)
(Roederer et al., 2014). Data from the UFD galaxies are taken from: Boötes I
(Feltzing et al., 2009; Frebel et al., 2016; Gilmore et al., 2013; Ishigaki et
al., 2014; Norris et al., 2010; Waller et al., 2022), Boötes II (Ji et al.,
2016), Carina II (Ji et al., 2020a), Carina III (Ji et al., 2020a), Coma
Berenices (Frebel et al., 2010; Waller et al., 2022), Grus I (Ji et al.,
2019), Grus II (Hansen et al., 2020), Hercules (Koch et al., 2008), Horologium
I (Marshall et al. in prep.), Leo IV (Simon et al., 2010), Pisces II (Spite et
al., 2018), Reticulum II (Ji et al., 2016b, 2019; Hayes et al., 2023), Segue 1
(Frebel et al., 2014; Norris et al., 2010), Segue 2 (Roederer & Kirby, 2014),
Triangulum II (Ji et al., 2019), Tucana II (Ji et al., 2016c; Chiti et al.,
2018, 2023), Tucana III (Hansen et al., 2017; Marshall et al., 2019), and Ursa
Major II (Frebel et al., 2010).

Figure 2:  $\mathrm{[X/Fe]}$ derived abundances for J0117 (black square)
compared to abundances of UFD galaxies (colored dots, see text for references
4) and stars from the MW halo (gray dots Roederer et al. 2014). J0117 follows
the general abundance trends of what is seen in most UFD galaxy stars.  Figure
3:  $\mathrm{[X/Fe]}$ derived neutron-capture element abundances for J0117
(black square) compared to abundances of other UFD galaxies (colored dots, see
text for references and Fig. 2 for legend) and stars from the MW halo (gray
dots; Roederer et al. 2014). Upper limits are designated with downward-
pointing triangles. J0117 has low abundances of neutron-capture elements,
which is characteristic of most UFD galaxies.

###  4.1 Alpha elements

We derive abundances for the $\alpha$-elements Mg, Si, and Ca from equivalent
widths. We find a Mg I abundance of $\mathrm{[Mg~{}I/Fe]}=0.36\pm{0.18}$ from
four Mg I absorption features, a Si I abundance of
$\mathrm{[Si~{}I/Fe]}=0.59\pm{0.24}$ using three Si I lines, and a Ca I
abundance of $\mathrm{[Ca~{}I/Fe]}=0.23\pm{0.11}$ from 20 Ca I lines. J0117
shows a general enhancement in $\alpha$-elements as a result of enrichment by
core-collapse SN, following the trend seen in metal-poor MW halo and other UFD
galaxy stars (Frebel et al., 2010; Lai et al., 2011; Simon, 2019).

###  4.2 Carbon and odd-Z elements

Abundances for C, Na, Al, K, and Sc were derived from EW and spectral
synthesis analysis. A $\mathrm{[C/Fe]}$ of $-$$0.31\pm{0.18}$ was determined
from the CH G-band in regions around 4310 Å and 4230 Å via spectral synthesis.
The C and O abundances are coupled through the CO molecule, thus, a knowledge
of the O abundance of the star is needed to derive the C abundance. As we
could not derive an O abundance for the star, we assumed a standard
$\alpha$-enhanced O abundance of $\mathrm{[O/Fe]}=0.4$ for the spectral
synthesis of CH, since UFDs are known to be $\alpha$-enhanced and J0117 is
$\alpha$-enhanced. The C abundance of a star is also altered as the star
evolves. Using the tool from Placco et al. (2014), a carbon correction of
$\Delta\mathrm{[C/Fe]}$ = 0.57 dex was determined, resulting in an original C
abundance of the star of $\mathrm{[C/Fe]}=0.26$. A Na I abundance of
$\mathrm{[Na~{}I/Fe]}=0.17\pm{0.25}$ and an Al I abundance of
$\mathrm{[Al~{}I/Fe]}=-0.33\pm{0.41}$ were derived from EW analysis of the Na
D resonance lines and two Al lines, respectively. EW analysis of the two K
lines at 7664 and 7698 Å was used to derive a K I abundance of
$\mathrm{[K~{}I/Fe]}=0.81\pm{0.25}$. From spectral synthesis of six Sc II
lines, we derive a Sc II abundance of $\mathrm{[Sc~{}II/Fe]}=-0.31\pm{0.19}$.

As can be seen in Figure 2, the C and Na abundances for J0117 follow those of
other UFD galaxies and metal-poor MW halo stars. The Al abundance for J0117
may be slightly higher than what is seen in other UFD galaxies, and metal-poor
MW halo stars, but the uncertainty of $\pm 0.41$ is high. J0117 has an
elevated K abundance compared to the average K abundances found in MW halo
stars. In addition, J0117 exhibits a lower Sc abundance than what is found in
most UFD galaxies and metal-poor MW halo stars.

###  4.3 Iron peak elements

EW analysis was used to derive abundances for the iron peak elements Ti, Cr,
Ni, and Zn and spectral synthesis for the elements V, Mn, and Co. We identify
14 Ti I and 23 Ti II lines in the spectrum and find abundances of
$\mathrm{[Ti~{}I/Fe]}=-0.23\pm{0.22}$ and
$\mathrm{[Ti~{}II/Fe]}=0.06\pm{0.18}$. V I and V II abundances were determined
to be $\mathrm{[V~{}I/Fe]}=-0.01\pm{0.18}$ and
$\mathrm{[V~{}II/Fe]}=0.37\pm{0.18}$ from two and three lines respectively and
Cr I and Cr II to be $\mathrm{[Cr~{}I/Fe]}=-0.30\pm{0.14}$ and
$\mathrm{[Cr~{}II/Fe]}=-0.07\pm{0.13}$ from nine and two lines respectively.
There is a slight discrepancy between the abundances we derive from the
neutral and ionized lines of the elements mentioned above, which may be due to
the 1D LTE nature of the analysis. However, for most of the elements, the
neutral and ionized abundances agree within uncertainties. Finally, we derive
Ni I and Zn I abundances of $\mathrm{[Ni~{}I/Fe]}=0.09\pm{0.13}$ and
$\mathrm{[Zn~{}I/Fe]}=0.01\pm{0.15}$ using 12 Ni lines and one Zn line.
Comparing these abundances to those of other UFD galaxy and MW halo stars (see
Figure 2) reveals the Cet II star to have relatively low Ti and V abundances,
while the remainder of the iron peak element abundances for this star are
similar to other UFD galaxy and metal-poor MW halo stars.

###  4.4 Neutron-capture elements

Absorption features from the two neutron-capture elements Sr and Ba were
identified in the spectrum, and using spectral synthesis, we derive abundances
for these elements of $\mathrm{[Sr~{}II/Fe]}=-2.10\pm{0.43}$ and
$\mathrm{[Ba~{}II/Fe]}=-2.23\pm{0.39}$. The synthesis of the Sr 4077 Å and Ba
4554 Å lines are shown in Figure 3. The Sr abundance derived for J0117 is low
and similar to what is found in most other UFD galaxies, and the Ba abundance
is lower than the majority of other UFD galaxy stars, see Figure 3. No Eu
could be detected in the spectrum, so we derive a 3$\sigma$ upper limit of
$\mathrm{[Eu~{}II/Fe]}<-0.33$ from the 4129 Å line, plotted as a downwards
pointing triangle in Figure 3.

Figure 4: Spectral synthesis of the 4077 Å Sr II line (top) and 4554 Å Ba II
line (bottom). Black squares are the observed spectrum of J0117. The dotted
line shows a synthesis not including Sr or Ba, and the blue line shows a
synthesis with the $\mathrm{[Sr/Fe]}\pm\sigma_{\mathrm{[Sr/Fe]}}$ or
$\mathrm{[Ba/Fe]}\pm\sigma_{\mathrm{[Ba/Fe]}}$ value derived for J0117.

##  5 Discussion

###  5.1 Nature of Cetus II

The nature of the Cet II overdensity is still debated in the literature. The
DES discovery paper (Drlica-Wagner et al., 2015) and a more recent paper from
the DES based on analysis of more extensive photometry (Drlica-Wagner et al.,
2020) both classify it as a probable UFD galaxy, while Conn et al. (2018), who
presented GMOS-S photometry of the system, suggested it is more likely a part
of the Sagittarius (Sgr) tidal stream. However, the proper motion of the Cet
II system,
$(\mu_{\alpha\cos{\delta}},\mu_{\delta})=(+2.8\pm{0.06},0.5\pm{0.06})~{}{\rm
mas~{}yr^{-1}}$ (Pace & Li, 2019; Pace et al., 2022) does not agree with the
proper motion of the Sgr stream
$(\mu_{\alpha\cos{\delta}},\mu_{\delta})=(-1,-3)~{}{\rm mas~{}yr^{-1}}$
(Vasiliev et al., 2021) and therefore rules out an association with the Sgr
stream. As well as this, the velocity measured for J0117 does not agree with
the velocity of the Sgr stream (Vasiliev et al., 2021) which further confirms
that they are not associated. This will be explored more in forthcoming work
(Simon et al. in prep.).

The size of Cet II ($r_{1/2}\sim$17pc) (Drlica-Wagner et al., 2015) lies in an
ambiguous region in the size-luminosity plane where the star cluster and UFD
populations overlap (Willman & Strader, 2012). As discussed in Willman &
Strader (2012), the most effective classification for these systems is a
velocity dispersion followed by a metallicity dispersion. These have not yet
been measured for Cet II.

In this paper, we have derived the abundances of one star, J0117, in Cet II,
and the general pattern of this star is similar to other UFD galaxy stars (see
Figure 2) and hence compatible with a UFD galaxy classification of Cet II.
This is particularly supported by the very low abundance of the neutron-
capture elements Sr and Ba, which have been found to be characteristic of UFD
galaxies and can be used to classify a system (Ji et al., 2019), in
particular, this abundance feature distinguishes UFD galaxy stars from GC
stars. Furthermore, we suggest that a high K abundance is another potential
UFD abundance signature, which will be explored further in Section 5.2.2. No
abundance analysis exists for Sgr stream stars at metallicities similar to
J0117; however, there have been abundance analyses for stars in the Sgr dwarf
spheroidal galaxy (dSph) at similar metallicities to J0117. The Sgr dSph stars
show abundance signatures that we do not find in J0117, including higher
neutron capture element abundances
($-0.2<\mathrm{[Sr/Fe]}<0.6$,$-0.8<\mathrm{[Ba/Fe]}<0.8$,
$-0.17<\mathrm{[Eu/Fe]}<0.81$ Hansen et al. 2018; Reichert et al. 2020). Thus,
the chemical abundances of Cet II do not support an association with the Sgr
dSph. For the remainder of the discussion, we will adopt the classification of
Cet II as a UFD galaxy.

###  5.2 Abundance pattern of J0117

With its low metallicity, slight enhancement of the $\alpha$-elements, and low
abundances of neutron-capture elements, J0117 generally follows the same
abundance trends seen in the majority of UFD galaxy stars. However, the
derived abundances for some elements stand out and warrant further discussion.

####  5.2.1 Low Sc, Ti, and V abundances

One specific feature of J0117 is that the $\mathrm{[Sc/Fe]}$,
$\mathrm{[Ti/Fe]}$, and $\mathrm{[V/Fe]}$ ratios are $\sim 0.3,0.2,$ and $0.1$
dex lower, respectively, in this star than the average of what is seen in most
other UFD galaxy stars and metal-poor MW halo stars (see Figure 2). We plot
the 5657 Å Sc II and 5381 Å Ti II lines in the spectrum of J0117 (black)
compared to spectra of stars HD 26297 (green) and BD+29 2356 (blue) (I.
Roederer private comm.), in Figure 5. HD 26297 and BD+29 2356 were chosen
because they have similar model atmosphere parameters to J0117 but higher Sc
and Ti abundances. However, it should be noted that there is a difference in
$\mathrm{[Fe/H]}$ which amplifies the intrinsic difference in
$\mathrm{[Sc/Fe]}$ and $\mathrm{[Ti/Fe]}$. The parameters for each star and
abundances for each line are listed in Table 5. It is clear from Figure 5 that
the absorption features of Sc and Ti are weaker in the spectrum of J0117, thus
supporting the lower abundances derived for these elements in this star.

Figure 5:  Comparison of Sc II and Ti II absorption lines in J0117 to the MW
halos stars BD+29 2356, and HD 26297 from Roederer et al. (2014). The
absorption features for J0117 are weaker than what is seen for BD+29 2356 and
HD 26297. Model atmosphere parameters and Sc and Ti abundances for each star
are shown in Table 5. Table 5: Model Atmosphere Parameters and Sc and Ti
Abundances for Comparison Stars and J0117.

ID | $T_{\rm eff}$ | $\log g$ | $\xi$ |  Model $\mathrm{[Fe/H]}$ |  $\log\epsilon$(Sc II) |  $\log\epsilon$(Ti II)
---|---|---|---|---|---|---
| (K) | (cgs) |  (km s-1) |  |  |
HD 26297 | 4400 | 1.10 | 1.75 | -1.72 | 1.57 | 3.67
BD +29 2356 | 4710 | 1.75 | 1.50 | -1.62 | 1.65 | 3.61
J0117 | 4727 | 1.40 | 1.89 | -2.09 | 0.63 | 3.00

Note. — Abundances for BD +29 2356 and HD 26297 are taken from Roederer et al.
(2014).

The Sc, Ti, and V abundances for metal-poor stars have been investigated by
Sneden et al. (2016); Cowan et al. (2020), and Ou et al. (2020), who all found
that the abundances of these three elements are positively correlated in
metal-poor stars ($\mathrm{[Fe/H]}<-2$), indicating that these three elements
have linked nucleosynthetic origins. While the main focus of Cowan et al.
(2020) is the analysis of UV spectra taken with the Hubble Space Telescope of
three metal-poor stars, they also investigated the correlation among Sc, Ti,
and V abundances derived in several large spectroscopic studies (e.g., Cayrel
et al., 2004; Cohen et al., 2004, 2008; Barklem et al., 2005; Lai et al.,
2008; Yong et al., 2013; Roederer et al., 2014), and found that the abundances
reported in these studies also show correlations between the three elements.

Figure 6:  Best-fit SN yield model (Heger & Woosley, 2010) for derived
abundances in J0117 (top) and HD 84937 (bottom). Black points are the derived
abundances for J0117, gray points are the derived abundances for HD 84937
(Sneden et al., 2016; Spite et al., 2017), and the red line is the best-fit SN
yield model for each star.

Ti and V are produced via explosive Si- and O- burning in core-collapse SN
(Woosley & Weaver, 1995; Sneden et al., 2016), with models producing similar
V/Ti ratios for both the Si and O SN ejecta components (e.g., Pignatari et
al., 2016). Sneden et al. (2016) suggested that the correlation of the Ti and
V abundances in metal-poor stars points to co-production of these elements and
that their abundances can be used to explore properties of the progenitor and
its explosion, such as mass and explosion energy. The case of Sc is a bit more
complicated. Although it can be produced in the same explosive Si- and
O-burning layers (Woosley & Weaver, 1995), it is more efficiently produced via
neutrino feedback or during $\alpha$-rich freeze-out conditions (Woosley &
Weaver, 1995; Fröhlich et al., 2006). This fact led Sneden et al. (2016) to
argue that, although a correlation is seen, the Sc/Ti and Sc/V ratios are less
useful diagnostics than V/Ti ratios of the properties of the progenitor star.
Generally, theoretical SN models often underproduce all three elements
compared to the abundances derived from metal-poor stars (Kobayashi et al.,
2020, and references therein). However, it was suggested by both Cowan et al.
(2020) and Ou et al. (2020) that more energetic SN, like hypernovae, might
contribute to the abundances of these elements.

Contrary to the star analyzed in Sneden et al. (2016) and the three stars
presented in Cowan et al. (2020), which all display higher than normal Sc, Ti,
and V abundances, J0117 exhibits lower than average abundances for these
elements. Since the chemistry of these stars is a direct fingerprint of the
nucleosynthesis process of the first massive (Population III) stars to form in
the galaxies, they can be used to constrain properties such as mass range,
rotation, and explosion energies of the Pop III stars. To investigate the
signature of J0117, we used the STARFIT tool666https://starfit.org/ to match
SN yields from Heger & Woosley (2010) to the abundances of J0117. The STARFIT
code calculates a $\chi^{2}$ statistic using the derived abundances and upper
limits to determine the best-fit SN yields. For J0117, we obtained models with
$\chi^{2}_{red}<2$, which returned models in a narrow range of progenitor
masses, M⊙=10.6-13.6. The best fit ($\chi^{2}_{red}$= 1.12) was achieved with
a progenitor mass of 11.2 M⊙ and a modest explosion energy of 0.6 B 777B=1
Bethe = $10^{51}$ erg. Figure 6 shows the abundances for J0117 and best-fit
model yields. This model provides a good fit to most elements, with only a few
discrepancies. It should be noted that it is unlikely that this star was
polluted by a singular Pop III SN, but this tool gives an idea of the most
dominant source of metals. To demonstrate the range of masses that can come
from these abundances, we also derived the best-fit yields for the star from
Sneden et al. (2016), HD 84937. HD 84937 is used for comparison because it has
high Sc, Ti, and V abundances compared to the average halo stars and thus
represents the high end of the abundance distribution, with J0117 representing
the low end. Abundances for the light elements were taken from Spite et al.
(2017), who adopted similar stellar parameters for HD 84937 as Sneden et al.
(2016). For HD 84937, the best fit was achieved with a progenitor mass of 18.1
M⊙ and explosion energy of 1.8 B. This demonstrates how the abundances for
these elements can be used to gain information about the progenitor masses and
explosion energies as well as how these parameters can vary with different
abundances. Furthermore, the progenitor mass found for the Cet II star is also
somewhat lower than what has been found for stars in other UFD galaxies
(Hansen et al., 2020), pointing to a wide range of progenitor masses existing
in UFD galaxies.

Another possibility to explain the low Sc and V abundances that we see is
pollution from Type Ia SN. As seen in Bravo et al. (2019), the production of
Sc and V in Type Ia SN is metallicity dependent and decreases with decreasing
metallicity. Hence, it is possible that the low Sc and V abundances are a
result of increased Fe injection from Type Ia SN with low Sc and V production.
However, the $\alpha$-element abundance ratios in J0117 are slightly enhanced
and compatible with core-collapse SN production and don’t suggest extra Fe
injection from Type Ia SN.

####  5.2.2 High K abundance

Another peculiar feature of J0117 is the somewhat high $\mathrm{[K/Fe]}$
abundance of 0.81 derived for this star. But, as can be seen in Figure 2,
J0117 is not a complete outlier. High K has been found in several other UFD
galaxy stars. In fact, in Figure 2, it can be seen that stars in UFD galaxies
generally have higher $\mathrm{[K/Fe]}$ ratios than MW halo stars at similar
metallicity. However, K is known to suffer from non-local thermodynamic
equilibrium (NLTE) effects (de La Reza & Mueller, 1975). We, therefore,
corrected the K abundance derived for J0117 and the K abundances of the UFD
galaxy literature sample following the study of Reggiani et al. (2019) and got
a $\mathrm{[K/Fe]}$ abundance of $0.64$ for J0117. The NLTE-corrected K
abundances for J0117 and other UFD galaxy stars, along with NLTE-corrected K
abundances of halo stars (Roederer et al., 2014), are shown in the middle
panel of Figure 7, while the LTE abundances are shown in the panel on the
left. It can be seen that the offset between a subset of the UFD galaxies and
the halo stars remains after the correction is applied. Ivanova & Shimanskiĭ
(2000) also looked at the NLTE corrections for K and provided corrections
slightly larger than Reggiani et al. (2019). However, even with these
corrections, the offset between the MW halo stars and a subset of the UFD
galaxies remains.

Apart from the NLTE effects, the K lines also have hyperfine structure.
However, this is very weak and, therefore, generally not included in the
abundance analysis. But for completeness, we derived K abundances from the two
lines used above via spectral synthesis, including hyperfine structure. As
expected, the effect on the abundances was minimal (¡0.1). Hence, for
compatibility with the literature data, the discussion below is based on our K
abundance derived from the EW analysis.

Figure 7: Left: LTE $\mathrm{[K/Fe]}$ abundances for J0117 (black square),
UFDs (colored dots, see Fig.2 for legend), and MW halo (gray dots) stars.
Middle: NLTE corrected $\mathrm{[K/Fe]}$ abundances for J0117, UFDs, and MW
halo stars. Right: NLTE corrected $\mathrm{[K/Fe]}$ abundances for J0117, MW
halo stars, GC stars (colored x), and stars from stellar streams with GC
progenitor (colored diamonds). There is an offset in $\mathrm{[K/Fe]}$
abundance between some UFD galaxies and the average value for MW halo stars.
Also, the GC and stellar stream $\mathrm{[K/Fe]}$ abundances mostly overlap
with the MW halo stars and are lower than what is seen in UFD galaxy stars.

Although the abundances of J0117 are compatible with Cet II being a UFD
galaxy, we searched the literature for K abundances in GCs to explore if this
abundance signature is unique to stars in UFD galaxies or can be found in
other types of small stellar systems. The results of this exercise are plotted
in the right panel of Figure 7, where the GC NLTE-corrected K abundances for
stars in ten GCs (M54 (Carretta, 2022), NGC 2419 (Mucciarelli et al., 2012),
NGC 104, NGC 6809 (Mucciarelli et al., 2017), NGC 6752 (Carretta et al.,
2007), NGC 1904, NGC 5139, NGC 6121, NGC 6397, NGC 7099 (Carretta et al.,
2013)) are marked with crosses. As can be seen, for most GCs, the K abundances
overlap with the halo stars, suggesting that the high K abundances seen in UFD
galaxy stars may be a unique abundance signature of these systems. The
exception is NGC 2419, in which stars with a wide range of K abundances are
found, including K-rich ($\mathrm{[K/Fe]}>1$) stars. However, the K-rich stars
in this cluster are all Mg-poor ($\mathrm{[Mg/Fe]}$ ¡ 0) (Cohen & Kirby, 2012;
Mucciarelli et al., 2012), thus making them distinguishable from the UFD
galaxy stars, which usually exhibit a small enhancement in Mg (J0117, for
example, has $\mathrm{[Mg/Fe]}=0.36$). Although we find K abundances in a
number of GCs, most have metallicities higher than the UFD galaxy stars with K
abundances, hampering a direct comparison. However, in recent years a number
of stellar streams have been detected that are thought to be the remnants of
more metal-poor GCs accreted by the MW, such as the Phoenix, Aliqa Uma, and
ATLAS streams (Li et al., 2019; Wan et al., 2020; Casey et al., 2021; Li et
al., 2022). Ji et al. (2020) derived K abundances for stars in the three GC
streams listed above. These abundances are plotted as diamonds in Figure 7.
Similar to the present-day GCs, the stream stars also have K abundances
overlapping with the MW halo stars. Hence, high K abundances seem to be a
characteristic abundance signature for at least some UFD galaxies.

The nucleosynthetic origin of the high K abundances seen in UFD galaxy stars
is yet unknown. K is created through hydrostatic oxygen shell burning and
explosive oxygen burning with yields depending on the progenitor mass (Woosley
& Weaver, 1995). Recently, models have shown that K production is increased
when rotation is introduced (Prantzos et al., 2018). Hence, the high K
abundances in UFD galaxy stars could suggest that some of these stars were
enriched by massive rotating stars. The models also show low Sc production in
massive rotating stars (Prantzos et al., 2018), which is in agreement with the
Sc in J0117 ($\mathrm{[Sc/Fe]=-0.31}$). However, it is not a good fit for the
other UFD stars that do not have low Sc.

Since the high K abundances are mainly seen in the UFD galaxy stars, it is
possible that this signature is tied to the slow chemical evolution of these
systems. Although unlikely given the overall abundance signature of J0117, the
K could also come from a nucleosynthesis source with a time-delayed
contribution, like Type Ia SN. One candidate could be Ca-rich transients
(Kasliwal et al., 2012), the lightcurves of which Polin et al. (2021) recently
found to match models of sub-Chandrasekhar Type Ia SNe with low mass
progenitors. However, more modeling is needed to establish if these events can
contribute to the K abundances of UFD galaxy stars.

##  6 Summary

We have performed a detailed chemical abundance analysis of J0117, the
brightest star of the Cet II UFD galaxy candidate. Our analysis shows that
this star is a metal-poor, $\alpha$-enhanced ($\alpha$-elements $\gtrapprox
0.4$) star with low abundances of the neutron-capture elements
($\mathrm{[Sr/Fe]}=-2.10$ and $\mathrm{[Ba/Fe]}=-2.23$), following the trends
seen for chemical analysis of other UFD galaxy stars. Thus, although there is
still debate on the classification of this system, the results of this
chemical analysis suggest Cet II is a UFD galaxy. Further observations and
chemical analysis of more stars in this system will help to fully determine
the nature of Cet II. The analysis revealed that the star exhibits slightly
lower $\mathrm{[Sc/Fe]}$, $\mathrm{[Ti/Fe]}$, and $\mathrm{[V/Fe]}$ abundances
compared to other UFD galaxy and MW halo stars at similar metallicities. It
has been suggested that the abundances for $\mathrm{[Sc/Fe]}$,
$\mathrm{[Ti/Fe]}$, and $\mathrm{[V/Fe]}$ can be used as a diagnostic for
progenitor mass; we thus compared the abundances of J0117 with the Heger &
Woosley (2010) SNe yields and determined a best-fit with a progenitor mass of
11.2M⊙, somewhat lower than the progenitor masses found for stars in other UFD
galaxies (Hansen et al., 2020). Finally, we derive a K abundance of
$\mathrm{[K/Fe]}=0.81$ for J0117, which, even after it has been corrected for
NLTE effects, is somewhat higher than the K abundances derived for MW halo
stars at similar metallicities. We note that a number of UFD galaxies have
high K abundances compared to the MW halo stars, and by including K abundances
for stars in GCs and streams in our comparison, we find that this is a unique
signature of some UFD galaxy stars.

The authors would like to thank Dr. Maria Drout for collecting the spectra of
J0117 from November 2017 and Dr. Andrew McWilliam for atomic K line data and
insightful comments on K abundances in stars. This research made extensive use
of the SIMBAD database operated at CDS, Straasburg, France (Wenger et al.,
2000), arXiv.org, and NASA’s Astrophysics Data System for bibliographic
information. Funding for the DES Projects has been provided by the U.S.
Department of Energy, the U.S. National Science Foundation, the Ministry of
Science and Education of Spain, the Science and Technology Facilities Council
of the United Kingdom, the Higher Education Funding Council for England, the
National Center for Supercomputing Applications at the University of Illinois
at Urbana-Champaign, the Kavli Institute of Cosmological Physics at the
University of Chicago, the Center for Cosmology and Astro-Particle Physics at
the Ohio State University, the Mitchell Institute for Fundamental Physics and
Astronomy at Texas A&M University, Financiadora de Estudos e Projetos,
Fundação Carlos Chagas Filho de Amparo à Pesquisa do Estado do Rio de Janeiro,
Conselho Nacional de Desenvolvimento Científico e Tecnológico and the
Ministério da Ciência, Tecnologia e Inovação, the Deutsche
Forschungsgemeinschaft and the Collaborating Institutions in the Dark Energy
Survey. The Collaborating Institutions are Argonne National Laboratory, the
University of California at Santa Cruz, the University of Cambridge, Centro de
Investigaciones Energéticas, Medioambientales y Tecnológicas-Madrid, the
University of Chicago, University College London, the DES-Brazil Consortium,
the University of Edinburgh, the Eidgenössische Technische Hochschule (ETH)
Zürich, Fermi National Accelerator Laboratory, the University of Illinois at
Urbana-Champaign, the Institut de Ciències de l’Espai (IEEC/CSIC), the
Institut de Física d’Altes Energies, Lawrence Berkeley National Laboratory,
the Ludwig-Maximilians Universität München and the associated Excellence
Cluster Universe, the University of Michigan, NSF’s NOIRLab, the University of
Nottingham, The Ohio State University, the University of Pennsylvania, the
University of Portsmouth, SLAC National Accelerator Laboratory, Stanford
University, the University of Sussex, Texas A&M University, and the OzDES
Membership Consortium. Based in part on observations at Cerro Tololo Inter-
American Observatory at NSF’s NOIRLab (NOIRLab Prop. ID 2012B-0001; PI: J.
Frieman), which is managed by the Association of Universities for Research in
Astronomy (AURA) under a cooperative agreement with the National Science
Foundation. The DES data management system is supported by the National
Science Foundation under Grant Numbers AST-1138766 and AST-1536171. The DES
participants from Spanish institutions are partially supported by MICINN under
grants ESP2017-89838, PGC2018-094773, PGC2018-102021, SEV-2016-0588,
SEV-2016-0597, and MDM-2015-0509, some of which include ERDF funds from the
European Union. IFAE is partially funded by the CERCA program of the
Generalitat de Catalunya. Research leading to these results has received
funding from the European Research Council under the European Union’s Seventh
Framework Program (FP7/2007-2013) including ERC grant agreements 240672,
291329, and 306478. We acknowledge support from the Brazilian Instituto
Nacional de Ciência e Tecnologia (INCT) do e-Universo (CNPq grant
465376/2014-2). This manuscript has been authored by Fermi Research Alliance,
LLC under Contract No. DE-AC02-07CH11359 with the U.S. Department of Energy,
Office of Science, Office of High Energy Physics. T.T.H acknowledges support
from the Swedish Research Council (VR 2021-05556) Magellan:Clay

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