#  JWST Identification of Extremely Low C/N Galaxies with [N/O]$\gtrsim 0.5$
at $z\sim 6-10$
Evidencing the Early CNO-Cycle Enrichment and a Connection with Globular
Cluster Formation

Yuki Isobe Institute for Cosmic Ray Research, The University of Tokyo, 5-1-5
Kashiwanoha, Kashiwa, Chiba 277-8582, Japan  Department of Physics, Graduate
School of Science, The University of Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo
113-0033, Japan<EMAIL_ADDRESS>Masami Ouchi National Astronomical
Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan  Institute
for Cosmic Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa,
Chiba 277-8582, Japan  Department of Astronomical Science, SOKENDAI (The
Graduate University for Advanced Studies), Osawa 2-21-1, Mitaka, Tokyo,
181-8588, Japan  Kavli Institute for the Physics and Mathematics of the
Universe (WPI), University of Tokyo, Kashiwa, Chiba 277-8583, Japan  Nozomu
Tominaga National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka,
Tokyo 181-8588, Japan  Department of Physics, Faculty of Science and
Engineering, Konan University, 8-9-1 Okamoto, Kobe, Hyogo 658-8501, Japan
Department of Astronomical Science, SOKENDAI (The Graduate University for
Advanced Studies), Osawa 2-21-1, Mitaka, Tokyo, 181-8588, Japan  Kuria
Watanabe  Department of Astronomical Science, SOKENDAI (The Graduate
University for Advanced Studies), Osawa 2-21-1, Mitaka, Tokyo, 181-8588, Japan
Kimihiko Nakajima National Astronomical Observatory of Japan, 2-21-1 Osawa,
Mitaka, Tokyo 181-8588, Japan  Hiroya Umeda Institute for Cosmic Ray Research,
The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8582, Japan
Department of Physics, Graduate School of Science, The University of Tokyo,
7-3-1 Hongo, Bunkyo, Tokyo 113-0033, Japan  Hidenobu Yajima Center for
Computational Sciences, University of Tsukuba, Ten-nodai, 1-1-1 Tsukuba,
Ibaraki 305-8577, Japan  Yuichi Harikane Institute for Cosmic Ray Research,
The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8582, Japan
Hajime Fukushima Center for Computational Sciences, University of Tsukuba,
Ten-nodai, 1-1-1 Tsukuba, Ibaraki 305-8577, Japan  Yi Xu Institute for Cosmic
Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba
277-8582, Japan  Department of Astronomy, Graduate School of Science, The
University of Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo 113-0033, Japan  Yoshiaki Ono
Institute for Cosmic Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha,
Kashiwa, Chiba 277-8582, Japan  Yechi Zhang Institute for Cosmic Ray Research,
The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8582, Japan
Department of Physics, Graduate School of Science, The University of Tokyo,
7-3-1 Hongo, Bunkyo, Tokyo 113-0033, Japan

(Received X, 2023; Revised Y, 2023; Accepted Z, 2023)

###### Abstract

We present chemical abundance ratios of 70 star-forming galaxies at $z\sim
4$–10 observed by the JWST/NIRSpec ERO, GLASS, and CEERS programs. Among the
70 galaxies, we have pinpointed 2 galaxies, CEERS_01019 at $z=8.68$ and
GLASS_150008 at $z=6.23$, with extremely low C/N ([C/N]$\lesssim-1$),
evidenced with C iii]$\lambda\lambda$1907,1909, N iii]$\lambda$1750, and N
iv]$\lambda\lambda$1483,1486, which show high N/O ratios ([N/O]$\gtrsim 0.5$)
comparable with the one of GN-z11. CEERS_01019 (GLASS_150008) has (does not
have) a strong high ionization line that is explained by AGN (star-formation)
photoionization models, which is (not) similar to GN-z11. Such low C/N and
high N/O ratios found in CEERS_01019 and GLASS_150008 (additionally identified
in GN-z11) are close to the equilibrium of the CNO cycle, suggesting that
these 3 galaxies are enriched by metals processed by the CNO cycle. On the C/N
vs. O/H plane, these 3 galaxies do not coincide with Galactic Hii regions,
normal star-forming galaxies, and nitrogen-loud quasars with AGB stars, but
globular-cluster (GC) stars, indicating a connection with GC formation. We
compare C/O and N/O of these 3 galaxies with those of theoretical models, and
find that these 3 galaxies are explained by scenarios with dominant CNO-cycle
materials, i.e. Wolf-Rayet stars, supermassive ($10^{3}-10^{5}\ M_{\odot}$)
stars, and tidal disruption events, interestingly with a requirement of
frequent direct collapses. For all the 70 galaxies, we present measurements of
Ne/O, S/O, and Ar/O, together with C/O and N/O. We identify 4 galaxies with
very low Ne/O, $\log(\rm Ne/O)<-1.0$, indicating abundant massive ($\gtrsim
30\ M_{\odot}$) stars.

High-redshift galaxies (734); Galaxy chemical evolution (580); Galaxy
formation (595); Star formation (1569)

††journal: ApJ††software: astropy (Astropy Collaboration et al., 2013; The
Astropy Collaboration, 2018), PyNeb (Luridiana et al., 2015), Cloudy (Ferland
et al., 2013)

##  1 Introduction

Chemical abundance ratios of the inter-stellar medium (ISM) in early galaxies
are crucial to understanding stellar nucleosynthesis. Local dwarf star-forming
galaxies have $\alpha$-to-oxygen ($\alpha$/O) ratios such as neon-to-oxygen
(Ne/O) that are around solar abundances and remain mostly constant for their
gas-phase metallicity (e.g., Izotov et al., 2006; Kojima et al., 2021; Isobe
et al., 2022). These findings suggest that the majority of $\alpha$ elements
are produced by massive stars evolving into core-collapse supernovae (CCSNe)
and/or hypernovae (HNe; e.g., Nomoto et al. 2013). On the other hand,
nitrogen-to-oxygen (N/O) ratios of star-forming galaxies increase with
metallicity (e.g., Izotov et al., 2006; Pilyugin et al., 2012; Kojima et al.,
2021; Isobe et al., 2022), thought to originate from primary oxygen production
by massive stars and secondary nitrogen production by low and intermediate-
mass metal-rich stars evolving to asymptotic giant branch (AGB) stars (e.g.,
Vincenzo et al., 2016). Similarly, carbon-to-oxygen (C/O) ratios of star-
forming galaxies increase with metallicity due to massive stars and AGB stars
(e.g., Berg et al., 2019).

Before the arrival of the James Webb Space Telescope (JWST), these kinds of
chemical abundance studies were restricted to at most intermediate redshifts,
even with stacking analysis (Steidel et al., 2016) or lensed objects (e.g.,
Kojima et al., 2017). However, Near Infrared Spectrograph (NIRSpec) on JWST
can spectroscopically observe a near-infrared (NIR) wavelength range of 1–5
$\mu$m 10–1000 times more deeply than other spectrographs. This great
advancement has led to several emission line detections of high-redshift
($z\gtrsim 4$) galaxies in restframe ultraviolet (UV) to optical ranges, which
are vital for chemical abundance measurements. Shortly after Early Release
Observations (ERO; Pontoppidan et al. 2022; PID: 2736) became public,
Arellano-Córdova et al. (2022) have reported Ne/O ratios of 3 ERO galaxies and
the C/O ratio of one ERO galaxy at $z>7$. The 3 and 1 galaxies have Ne/O and
C/O ratios comparable to those of local galaxies at a given metallicity,
respectively.

However, Jones et al. (2023) have found that a $z=6.23$ galaxy observed by the
GLASS-JWST Early Release Science (ERS) program (hereafter GLASS; Treu et al.
2022; PID: 1324) has a very low $\log(\rm C/O)$ value of $-1.01\pm 0.22$. As
the low C/O value is consistent with those expected from yields of pure core-
collapse supernovae (Nomoto et al., 2013), Jones et al. (2023) have
interpreted the galaxy as having a young chemical composition where carbon
production by AGB stars is negligible.

Moreover, GN-z11 at $z=10.6$, photometricly identified by Oesch et al. (2016)
and spectroscopically confirmed by the JWST Advanced Deep Extragalactic Survey
(JADES; Bunker et al. 2023), has a super-solar value of $\log(\rm
N/O)\gtrsim-0.5$ (Cameron et al., 2023; Senchyna et al., 2023), which is
significantly higher than those of local galaxies at a given metallicity.
Given the ultra-high redshift of $z=10.6$ corresponding to 440 Myr after the
Big Bang, nitrogen production by AGB stars is difficult to explain the N/O
enhancement of GN-z11 (e.g., Cameron et al., 2023). Since the report, several
studies have been trying to explain the high N/O ratio of GN-z11 (see Section
4.2 for more details). However, it remains unknown whether there is another
galaxy with super-solar N/O ratios at high redshift. The carbon-nitrogen (C/N)
ratio has also not been studied, although it is usually assumed that both
carbon and nitrogen originate from AGB stars.

Furthermore, given that chemical evolution models (Watanabe et al., 2023)
predict significant changes of Ne/O ratios from the solar abundances at the
early formation phase (more details in Section 5.2), Ne/O ratios potentially
evolve toward higher redshifts.

The aim of this paper is to study the neon, carbon, nitrogen, and oxygen
abundances of high-$z$ galaxies using data from the NIRSpec public surveys to
characterize the nature of star formation at high redshifts. This paper is
organized as follows. Section 2 explains the data and sample we use. Our
analysis is described in Section 3. We report and discuss our results of
nitrogen, carbon, and oxygen abundances in Sections 4. Results and discussions
of Ne/O ratios are described in 5. Our findings are summarized in Section 6.
We assume a standard $\Lambda$CDM cosmology with parameters of ($\Omega_{\rm
m}$, $\Omega_{\rm\Lambda}$, $H_{0}$) = (0.3, 0.7, 70 km ${\rm s}^{-1}$ ${\rm
Mpc}^{-1}$). Throughout this paper, we use the solar abundance ratios of
Asplund et al. (2021). The notation [A/B] is defined as $\log(\rm A/B)$
subtracted by the solar abundance $\log(\rm A/B)_{\odot}$.

##  2 Data and Sample

We utilize JWST/NIRSpec data from the Early Release Observations (ERO;
Pontoppidan et al. 2022), taken in the SMACS 0723 lensing cluster field
(hereafter referred to as ERO data), the GLASS survey (Treu et al. 2022;
hereafter referred to as GLASS data), and the CEERS survey (Finkelstein et al.
2022; hereafter referred to as CEERS data). The ERO data were acquired using
medium-resolution ($R\sim 1000$) filter-grating pairs of F170LP-G235M and
F290LP-G395M, covering the wavelength ranges of 1.7–3.1 and 2.9–5.1 $\mu$m,
respectively. The total exposure time for the ERO data is 4.86 hours for each
filter-grating pair. The GLASS data were collected using high-resolution
($R\sim 2700$) filter-grating pairs of F100LP-G140H, F170LP-G235H, and
F290LP-G395H, covering the wavelength ranges of 1.0–1.6, 1.7–3.1, and 2.9–5.1
$\mu$m, respectively. The total exposure time for the GLASS data is 4.9 hours
for each filter-grating pair. The CEERS data were obtained using medium-
resolution filter-grating pairs of F100LP-G140M, F170LP-G235M, and
F290LP-G395M, covering the wavelength ranges of 1.0–1.6, 1.7–3.1, and 2.9–5.1
$\mu$m, respectively. The total exposure time for the CEERS data is 0.86 hours
for each filter-grating pair.

We use spectroscopic data that has been reduced by Nakajima et al. (2023).
Nakajima et al. (2023) have extracted the raw data from the MAST archive and
performed level-2 and 3 calibrations using the JWST Science Calibration
Pipeline with the reference file jwst_1028.pmap, whose flux calibration is
based on in-flight flat data. Checking the data, we identify 5, 15, and 50
galaxies at $z>4$ in the ERO, GLASS, and CEERS data, respectively. We refer to
these 70 ($=5+15+50$) galaxies as our sample galaxies, hereafter.

##  3 Analysis

###  3.1 Flux Measurement

We measure emission-line fluxes by fitting a Gaussian function convolved by
line-spread functions that Isobe et al. (2023) have derived from in-flight
NIRSpec data of a planetary nebula (PI: J. Muzerolle; PID: 1125). To measure
weak emission lines accurately, we fix redshift and velocity dispersion of the
Gaussian function to those of [O iii]$\lambda\lambda$4959,5007. To derive
Ne/O, C/O, and N/O ratios, we explore the following line detections: [Ne
iii]$\lambda$3869 (for neon abundance), C iii]$\lambda\lambda$1907,1909 (for
carbon abundance), N iii]$\lambda$1750111We refer to the quintet of N iii]
lines at restframe 1747, 1749, 1750, 1752, and 1754 Å as N iii]$\lambda$1750
in this paper. and N iv]$\lambda\lambda$1483,1486 (for nitrogen abundance), [O
ii]$\lambda\lambda$3727,3729, [O iii]$\lambda$5007, and O
iii]$\lambda\lambda$1661,1666 (for oxygen abundance), [O iii]$\lambda$4363
(for electron temperature $T_{\rm e}$), and Balmer lines of H$\alpha$,
H$\beta$, and H$\gamma$ (for hydrogen abundance and dust correction).
Hereafter, we just write [Ne iii], C iii], N iii], N iv], [O ii], [O iii], and
O iii] as meaning [Ne iii]$\lambda$3869, C iii]$\lambda\lambda$1907,1909, N
iii]$\lambda$1750, N iv]$\lambda\lambda$1483,1486, [O
ii]$\lambda\lambda$3727,3729, [O iii]$\lambda$5007, and O
iii]$\lambda\lambda$1661,1666, respectively, for simplicity. Thirty-five and 7
of our sample galaxies have $S/N>3$ detections of [Ne iii] and C iii],
respectively. We find that one of our sample galaxies, GLASS_150008, has a
detection of N iii] with $S/N=4.2$. We also identify the detection of N iv]
($S/N=4.3$) from one of our sample galaxies, CEERS_01019, which has also been
reported by Larson et al. (2023).

###  3.2 Emission-line Diagnostics

Figure 1: N iii]/O iii] (top left), N iv]/O iii] (bottom left), C iii]/N iii]
(top right), and C iii]/N iv] (bottom right) ratios as a function of N iv]/N
iii]. Observed emission-line ratios of GLASS_150008, CEERS_01019, and GN-z11
are represented by the double red circle, the red square, and the magenta
square, respectively. The emission-line ratios of GN-z11 are measured by
Maiolino et al. (2023) and Bunker et al. (2023). Cloudy (Ferland et al., 2013)
photoionization models with [N/O$]=0$ and $\log(Z/Z_{\odot})=-0.5$ based on
the stellar and AGN radiations are shown by the black solid and dashed curves,
respectively. These models with [N/O$]=0$ and different metallicities of
$\log(Z/Z_{\odot})=-1$ and 0.25 are shown by the semi-transparent blue and red
lines, respectively. Detailed measurements of the N/O and C/N ratios are made
in Sections 3.4 and 3.5.

The left panels of Figure 1 compare N iii]/O iii] and N iv]/O iii] ratios of
GLASS_150008 (double red circle) and CEERS_01019 (red square) with those of
GN-z11 (magenta square) as a function of N iv]/N iii], while we use N iii] and
N iv] fluxes measured by Maiolino et al. (2023) and the upper limit of O iii]
(Bunker et al., 2023)222Originally reported as $2\sigma$ upper limit in Bunker
et al. (2023) but converted to $3\sigma$ upper limit for consistency with our
flux measurements in Section 3.1.. The N iii] and N iv] fluxes from Maiolino
et al. (2023) and the O iii] upper limit from Bunker et al. (2023) are scaled
by C iii] lines. We find that GLASS_150008 and CEERS_01019 have high N iii]/O
iii] and N iv]/O iii] ratios, respectively, both of which are comparable to
those of GN-z11, which is reported to have a super-solar N/O ratio (Section
1). The right panels of Figure 1 also show that GLASS_150008 and CEERS_01019
have low C iii]/N iii] and C iii]/N iv] ratios, respectively, both of which
are comparable to those of GN-z11.

To compare the observed emission-line ratios, we construct photoionization
models using Cloudy (Ferland et al., 2013) for both stellar and AGN
radiations. To model young star-forming galaxies, we use BPASS (Stanway &
Eldridge, 2018) binary stellar radiations under the assumptions of the
instantaneous star-formation history with the stellar age of 10 Myr, upper
star mass cut of 100 $M_{\odot}$ with the Salpeter (1955) initial mass
function (IMF), and the hydrogen density $n_{\rm H}$ of 300 cm-3. The $n_{\rm
H}$ value is inferred from a typical value of the electron density ($n_{\rm
e}$) of $\gtrsim 300$ cm-3 in the ISM of $z>4$ galaxies (Isobe et al., 2023).
We also fix He/H and metal-to-oxygen ratios to be the solar abundances. We
vary O/H and ionization parameter ($U$) values within the ranges of
$-2\leq[$O/H$]\leq 1$ and $-3.5\leq\log(U)\leq-0.5$ in 0.25 and 0.25
increments, respectively. We set the stellar metallicity equal to the nebular
metallicity defined by the O/H ratio. We refer to this model as the young
stellar model, hereafter. We also construct the photoionization models for
extremely young star-forming galaxies with very massive stars (extremely young
massive stellar model, hereafter), assuming the stellar age of 1 Myr and upper
star mass cut of 300 $M_{\odot}$. The other parameters are the same as those
of the young stellar model. We also compute the stellar photoionization models
with [N/O$]=1$, whose other parameters are the same as those of the young
stellar model and extremely young massive stellar model.

We also construct AGN photoionization models whose incident radiation is
parameterized by the following 4 parameters: the big-bump temperature $T_{\rm
BB}$, the X-ray to UV ratio $\alpha_{\rm OX}$, the low-energy slope of the
big-bump component $\alpha_{\rm UV}$, and the X-ray component slope
$\alpha_{\rm X}$. Typical AGNs have $\alpha_{\rm OX}\sim-1.4$ (Zamorani et
al., 1981) and $\alpha_{\rm UV}\sim-0.5$ (Francis, 1993; Elvis et al., 1994).
We refer to the AGN model with the parameter set of $T_{\rm BB}=1.5\times
10^{5}$ K, $\alpha_{\rm OX}=-1.4$, $\alpha_{\rm UV}=-0.5$, and $\alpha_{\rm
X}=-1.0$ as the default AGN model. The flux density per frequency $f_{\nu}$ is
expressed by:

$f_{\nu}=\nu^{\alpha_{\rm UV}}\exp(-h\nu/kT_{\rm BB})\exp(-kT_{\rm
IR}/h\nu)+a\nu^{\alpha_{\rm X}},$ (1)

where $a$ is a coefficient corresponding to $\alpha_{\rm OX}$ and $T_{\rm IR}$
is an infrared cutoff temperature at $kT_{\rm IR}=0.01$ Ryd. The continuum
above 100 keV is also assumed to cut off by $\nu^{-2}$. Using the AGN
radiation, we calculate emission line ratios under the assumption of $n_{\rm
H}=300$ cm-3. We vary O/H, $U$, and N/O within the ranges of
$-2\leq[$O/H$]\leq 1$, $-3.5\leq\log(U)\leq-0.5$, $0\leq[$N/O$]\leq 2$, and
$-1\leq[$C/O$]\leq 0$ in 0.25, 0.25, 1, and 1 increments, respectively. We
also fix He/H and metal-to-oxygen ratios other than N/O to be the solar
abundances. Note that different assumptions of $n_{\rm H}$ from 10 to $10^{6}$
cm-3 can change key emission-line ratios of N iii]/O iii], C iii]/O iii]. and
C iii]/N iii] by only $\lesssim 0.1$ dex.

In Figure 1, we plot the young stellar models (solid) and the default AGN
models (dashed) with [N/O$]=0$ and 1 (corresponding to [C/N$]=0$ and $-1$). We
find that GLASS_150008, and CEERS_01019, and GN-z11 have high N iii]/O iii]
and N iv]/O iii] ratios close to those predicted by both the young stellar
model and the default AGN model with [N/O$]=1$. Similarly, the low C iii]/N
iii] and C iii]/N iv] ratios in the 3 galaxies are comparable to those of the
models with [C/N$]=-1$. These results indicate that the 3 galaxies have very
high N/O and low C/N ratios significantly above and below the solar abundance,
respectively. These results are not likely affected by the metallicity changes
from $\log(\rm Z/Z_{\odot})=-1$ (blue) to 0.25 (red) as shown in Figure 1,
while for accuracy, we derive metallicities first in Sections 3.3 and 3.5.

###  3.3 Nebular Property

We derive color excesses $E(B-V)$ from Balmer line ratios of
H$\beta$/H$\alpha$, H$\gamma$/H$\alpha$, and H$\gamma$/H$\beta$ as many as
possible by assuming the dust attenuation curve of Calzetti et al. (2000) and
the case B recombination. Intrinsic values of the Balmer line ratios are
calculated by PyNeb (Luridiana et al. 2015; v1.1.15). If none of the above
Balmer line ratios are available, we use a median value of $E(B-V)=0.07$ for
the other galaxies in our sample. Using the $E(B-V)$ values and Calzetti et
al. (2000)’s curve, we correct emission-line ratios for dust attenuation.

For the 9 galaxies with the measurements of electron temperatures of O iii
($T_{\rm e}(\rm O\,\textsc{iii})$) by Nakajima et al. (2023), we use the
$T_{\rm e}$ values and their errors. In addition, we derive $T_{\rm e}$(O iii)
values of 5 galaxies with [O iii]$\lambda$4363 and/or O iii]$\lambda$1666\.
The other galaxies in our sample are assumed to have $T_{\rm e}$(O
iii$)=15000$ K with 1$\sigma$ uncertainty of 5000 K. We calculate $T_{\rm
e}$(O ii) from $T_{\rm e}$(O iii) and the empirical relation of Garnett
(1992). We assume $n_{\rm e}$ values to be 300 cm-3.

We derive O+/H+ from [O ii]/H$\beta$ and $T_{\rm e}$(O ii) and O2+/H+ from [O
iii]$\lambda\lambda$4959,5007/H$\beta$ and $T_{\rm e}$(O iii). In our sample
galaxies where only [O ii] upper limits are present, we presume the smallest
values of [O ii] fluxes to be the ones obtained from [O iii] under the
assumption of $\log(U)=-1$. For these galaxies, we also suppose the true and
maximum values of [O ii] fluxes to be their 1$\sigma$ and 3$\sigma$ upper
limits, respectively. For our sample galaxies where [O ii] falls outside the
wavelength coverage of NIRSpec, we assume the true [O ii] fluxes of the
galaxies to be [O iii] divided by a median [O iii]/[O ii] ratio of the other
galaxies in our sample. The minimum and maximum values are assumed to be [O
iii] divided by [O iii]/[O ii] at $\log(U)=-1$ and $-3$, respectively. We then
obtain gas-phase metallicity $12+\log(\rm O/H)$ from the equation ${\rm
O/H}={\rm O}^{+}/{\rm H}^{+}+{\rm O}^{2+}/{\rm H}^{+}$ by ignoring neutral
oxygen and O3+ and higher-order oxygen ions as in e.g., Izotov et al. (2006).

We include errors of the used fluxes as well as uncertainties of the
assumptions of $T_{\rm e}$ and [O ii] for some of our sample galaxies into
$12+\log(\rm O/H)$ errors by Monte Carlo simulations. We derive 1000 values of
$12+\log(\rm O/H)$ with $T_{\rm e}$ and flux values such as [O ii] randomly
fluctuated by their errors under the assumption of the normal distribution.
Then, we derive the 16th and 84th percentiles of the distribution of the 1000
$12+\log(\rm O/H)$ values as the $\pm 1\sigma$ confidence interval of
$12+\log(\rm O/H)$. We note that we cannot measure $12+\log(\rm O/H)$ of 7
galaxies in our sample due to the lack of H$\beta$. The derived $T_{\rm e}$,
$E(B-V)$, and $12+\log(\rm O/H)$ values are summarized in Table A.

###  3.4 Abundance Ratio

Table 1: Atomic data Ion | Emission process | Transition probability | Collision Strength
---|---|---|---
(1) | (2) | (3) | (4)
H0 | Re | Storey & Hummer (1995) | —
C2+ | CE | Wiese et al. (1996) | Berrington et al. (1985)
N2+ | CE | Galavis et al. (1998) | Blum & Pradhan (1992)
O+ | CE | Froese Fischer & Tachiev (2004) | Kisielius et al. (2009)
O2+ | CE | Froese Fischer & Tachiev (2004) |  Storey et al. (2014)†
Ne2+ | CE | Froese Fischer & Tachiev (2004) | McLaughlin & Bell (2000)
Ar2+ | CE | Munoz Burgos et al. (2009) | Munoz Burgos et al. (2009)
S+ | CE | Rynkun et al. (2019) | Tayal & Zatsarinny (2010)
Table 2: Abundance Ratio of GLASS_150008, CEERS_01019, and GN-z11 ID | $12+\log(\rm O/H)$ | $\log(\rm N/O)$ | $\log(\rm C/O)$ | $\log(\rm C/N)$ | $\log(\rm Ne/O)$ | Rad.
---|---|---|---|---|---|---
(1) | (2) | (3) | (4) | (5) | (6) | (7)
GLASS_150008 | $7.65^{+0.14}_{-0.08}$ | $-0.40^{+0.05}_{-0.07}$ | $-1.08^{+0.06}_{-0.14}$ | $-0.68^{+0.09}_{-0.15}$ | $\cdots$ | stellar
CEERS_01019 | 8.23–8.50 | $>-0.01$ | $<-0.36$ | $<-0.35$ |  $(-0.72)$–$(-0.48)$ | AGN
GN-z11 | 8.58–9.23 | $>-0.09$ | $>-0.78$ |  $(-0.80)$–$(-0.54)$ | $>-1.85$ | AGN

We derive Ne/O, C/O, and N/O ratios from ionic abundance ratios with similar
ionization energies of Ne2+/O2+, C2+/O2+, and N2+/O2+ to minimize systematics
of ionization correction factors (ICFs). We calculate Ne2+/O2+ ratios from [Ne
iii]/[O iii] ratios. For GLASS_150008 with the O iii] detection, we use C
iii]/O iii] and N iii]/O iii] ratios to obtain C2+/O2+ and N2+/O2+ ratios with
less systematics. For the other galaxies in our sample, we calculate C2+/O2+
and N2+/O2+ ratios from [O iii] lines. We calculate these ionic abundance
ratios using PyNeb with the same transition probabilities and collision
strengths listed in Table 1.

To derive chemical abundance ratios (Ne/O, C/O, and N/O) from the ionic
abundance ratios (Ne2+/O2+, C2+/O2+, and N2+/O2+), we extract ICFs from the
young stellar model (Section 3.2). We have checked that our ICFs provide Ne/O
and C/O ratios consistent with those based on Izotov et al. (2006) and Berg et
al. (2019)’s ICFs, respectively. We have also confirmed that ICF(Ne2+/O2+),
ICF(C2+/O2+), and ICF(N2+/O2+) of all our sample galaxies are $\sim 1$. In
addition, the extremely young massive stellar model provides Ne/O, C/O, and
N/O values similar to those of the young stellar model. We derive the errors
of the Ne/O, C/O, and N/O ratios by Monte Carlo simulations in the same way as
for $12+\log(\rm O/H)$ (Section 3.3). For our sample galaxies without the
detections of [Ne iii], [C iii], or [N iii], we use $3\sigma$ upper limits of
these lines to obtain upper limits of Ne/O, C/O, or N/O ratios. We summarize
the abundance ratios in Table A.

We have confirmed that our derived values of C/O and $12+\log(\rm O/H)$ for
GLASS_150008 are consistent with those measured by Jones et al. (2023) within
$1\sigma$ level. We find that GLASS_150008 has a super-solar N/O value of
$\log(\rm N/O)=-0.40^{+0.05}_{-0.07}$ as predicted in Figure 1. The derived
properties of GLASS_150008 are listed in Table 2.

###  3.5 Abundance Ratio Re-estimation Based on AGN Radiation

Figure 2: Emission-line ratios of CEERS_01019 (red square and red shaded
region) and GN-z11 (magenta square and magenta shaded region). (Top) Cloudy
photoionization models of young massive stellar population and AGN are shown
by the solid and dashed lines, respectively, color-coded by metallicity. The
observed emission-line ratios are not reproduced by the stellar models but the
AGN models. (Second top, second bottom, and bottom) Comparison with the
observed line ratios and the AGN photoionization models with different
$12+\log(\rm O/H)$ (second top), N/O (second bottom), C/O (bottom left), and
C/N (bottom right). Ionization parameters of the models are determined by [O
iii]/[O ii] and N iv]/N iii] for CEERS_01019 and GN-z11, respectively. The
gray shaded regions represent uncertainties of the observed ionization
parameters and $12+\log(\rm O/H)$ (for the bottom middle and bottom panels).
Figure 3: Metallicity as a function of redshift. The metallicities of
CEERS_01019 (red square) and GN-z11 (magenta square) are based on their O/H
ratios. The black and blue symbols are FOREVER22 (Yajima et al., 2022a, b)
simulated galaxies of First0 and First1, respectively, while the dots and the
solid lines denote metallicities measured within the half-mass radii ($r_{\rm
half}$) and the total gas of the galaxies, respectively.

It should be noted that GLASS_150008 does not have very high ionization lines
that require the presence of AGN. We have also checked that GLASS_150008 does
not have a clear broad component in the H$\beta$ line profile. We thus
conclude that GLASS_150008 does not have AGN signatures at least in the
current data set.

Contrary to GLASS_150008, Larson et al. (2023) have reported that CEERS_01019
with the N iv] detection has an H$\beta$ broad component with the full-width
half maximum of $\sim 1200$ km s-1, which can be an evidence of AGN. Maiolino
et al. (2023) have also reported that GN-z11 has a detection of [Ne
iv]$\lambda\lambda$2422,2424 lines that require high-energy photons beyond
63.5 eV, which can also suggest that GN-z11 is an AGN.

Here, we check if emission line ratios of CEERS_01019 and GN-z11 cannot be
reproduced by stellar radiations. The solid curves in the top panels of Figure
2 represent predicted emission line ratios of the extremely young massive
stellar model. The model contains very massive stars up to 300 $M_{\odot}$,
which are expected to produce harder radiation than normal stellar
populations. However, we confirm that even the extremely young massive stellar
model cannot reproduce the observed N iv]/N iii] and [O iii]/[O ii] relation
of CEERS_01019 (red circle) or the observed [Ne
iv]$\lambda\lambda$2422,2424/[Ne iii] and N iv]/N iii] relation of GN-z11
(magenta square; Maiolino et al. 2023). These result suggest that CEERS_01019
and GN-z11 require hard emissions such as AGN as discussed by Larson et al.
(2023) and Maiolino et al. (2023), respectively.

Then, we revisit metallicity and N/O measurements of CEERS_01019 and GN-z11
with AGN radiations. We find that CEERS_01019 has the lower limit of N iv]/N
iii] and the upper limit of He ii/H$\delta$ simultaneously reproduced by the
default AGN model. We also find that both [Ne iv]/[Ne iii] and He ii/H$\delta$
ratios of GN-z11 measured by Maiolino et al. (2023) are reproduced by the
default AGN model except for $\alpha_{\rm OX}=-1.5$ (low-$\alpha_{\rm OX}$ AGN
model, hereafter). These AGN models are shown by the color-coded dashed lines
in the top panels of Figure 2, reproducing the observed emission-line ratios
of both CEERS_01019 and GN-z11.

We then re-estimate $12+\log(\rm O/H)$ and N/O of CEERS_01019 and GN-z11 by
using these AGN models. The middle left panel of Figure 2 shows the observed
[O iii]/H$\beta$ ratio of CEERS_01019 (red solid line; 1$\sigma$ uncertainty
is shown by the red shaded region) and those predicted by the default AGN
models with the $U$ value derived from the observed [O iii]/[O ii] ratio (gray
dashed line; 1$\sigma$ uncertainty based on the observed [O iii]/[O ii]
uncertainty is shown by the gray shaded region) as a function of $12+\log(\rm
O/H)$. Note that [O iii]/H$\beta$ values slightly decrease with increasing
N/O. With the default AGN models with [N/O$]=1$, we find that the observed and
model-predicted [O iii]/H$\beta$ ratios agree with each other at $12+\log(\rm
O/H)=8.23$–8.50. This value is slightly higher than that based on the stellar
radiation of $12+\log(\rm O/H)=7.94^{+0.56}_{-0.28}$. We then derive the N/O
ratio of CEERS_01019 by comparing the observed lower limit of N iv]/O iii]
with those predicted by the default AGN models with the metallicity range of
$12+\log(\rm O/H)=8.23$–8.50, $U$ based on the observed [O iii]/[O ii] ratio,
and different N/O ratios. We find that the observed lower limit of N iv]/O
iii] favors the model with the super-solar value of $[\rm N/O]>0.85$ (i.e.,
$\log(\rm N/O)>-0.01$) as predicted in Figure 1.

We derive C/O and Ne/O ratios of CEERS_01019 from C iii]/[O iii] and [Ne
iii]/[O iii] ratios in the same way as for the N/O ratio, using the default
AGN models with $12+\log(\rm O/H)=8.23$–8.50 and [N/O$]=1$. The bottom left
panel of Figure 4 describes the result of the C/O measurement as an example.
We obtain [C/O$]<-0.13$ and [Ne/O$]=(-0.09)$–0.15, and our derived C/O and N/O
ratios provide a low value of [C/N$]<-0.98$ (i.e., $\log(\rm C/N)<-0.35$).
Note that we can obtain a more strict upper limit of [C/N$]<-1.8$ based on the
C iii]/N iv] ratio. The observed ratios among N iii], N iv], and C iii] lines
are simultaneously reproduced by the default AGN model with $12+\log(\rm
O/H)=8.23$–8.50, $[\rm N/O]>0.85$, and [C/N$]<-1.8$.

We derive $12+\log(\rm O/H)$ and N/O of GN-z11 in the similar way to
CEERS_01019, while we use the low-$\alpha_{\rm OX}$ AGN model, N iv]/N iii]
(Maiolino et al., 2023) as an indicator of $U$, the [O iii]$\lambda$4363/[O
ii] (Bunker et al., 2023) for the metallicity measurement, and the lower limit
of N iii]/O iii] (Maiolino et al., 2023; Bunker et al., 2023) for the N/O
measurement. Using the H$\delta$/H$\gamma$ ratio of Maiolino et al. (2023), we
derive $E(B-V)=0.00$ under the assumption of $T_{\rm e}=15000$ K and Calzetti
et al. (2000)’s dust attenuation curve. We obtain $12+\log(\rm
O/H)=8.63$–9.21, which is higher than those based on the stellar radiation and
[O iii]$\lambda$4363 of $12+\log(\rm O/H)=7.03$–8.60 (Cameron et al., 2023)
and $12+\log(\rm O/H)=7.84^{+0.06}_{-0.05}$ (Senchyna et al., 2023). This is
because the AGN radiation can enhance [O iii]$\lambda$4363, which results in
overestimation of $T_{\rm e}$. Our derived metallicity is consistent with that
inferred from the deep blueshifted absorption of C iv
$\lambda\lambda$1548,1550 (Maiolino et al., 2023).

In addition, Figure 3 compares the metallicity of GN-z11 (magenta square)
based on the AGN radiation with those of First0 (black) and First1 (blue)
galaxies of the cosmological hydrodynamics zoom-in simulation FOREVER22
(Yajima et al., 2022a, b). Although the metallicities of the simulated
galaxies calculated for the total gas (solid) are lower than that of GN-z11,
we find that the simulated galaxies can instantaneously have super-solar
metallicities within their half-mass radii ($\sim 100$ pc at $z=10.6$, where
GN-z11 is located). This suggests that it is not very strange even for
high-$z$ galaxies such as GN-z11 (and CEERS_01019 at $z=8.679$; red square) to
have high metallicities around the solar value.

We derive C/N, C/O, and Ne/O ratios of GN-z11 from C iii]/N iii] (Maiolino et
al., 2023), C iii]/O iii] (Maiolino et al., 2023; Bunker et al., 2023), and
[Ne iii]/O iii] (Maiolino et al., 2023; Bunker et al., 2023) ratios in the
same way as for the N/O ratio, using the low-$\alpha_{\rm OX}$ AGN models with
$12+\log(\rm O/H)=8.23$–8.50 and [N/O$]=1$. The bottom right panel of Figure 4
describes the result of the C/N measurement as an example. We obtain a low
value of [C/N$]=(-1.43)$–($-1.17$), [C/O$]>-0.53$, and [Ne/O$]>-1.22$ (i.e.,
$\log(\rm C/N)=(-0.80)$–$(-0.54)$, $\log(\rm C/O)>-0.78$, and $\log(\rm
Ne/O)>-1.85$). We have confirmed that the C/N value based on C iii]/N iii] is
consistent with that based on C iii]/N iv] within the error level. We
summarize the metallicities and the abundance ratios of GLASS_150008,
CEERS_01019, and GN-z11 in Table 2.

##  4 Nitrogen, Carbon, and Oxygen Abundances

###  4.1 Result

Figure 4: N/O (top left), C/O (top right), C/N (bottom left) as a function of
$12+\log(\rm O/H)$. N/O as a function of C/O (bottom right). The measurements
of CEERS_01019 (red square) and GN-z11 (magenta square) are based on the AGN
photoionization models (see Section 3.5 and Figure 2), while those of
GLASS_150008 are based on the stellar photoionization model (Section 3.4). In
the top right panel, C/O ratios of the other galaxies in our sample are shown
by the smaller red circles (with C iii] detection) and the white circles (with
strong upper limits of C iii]), respectively. The double circles correspond to
the measurements with $T_{\rm e}$ determinations. The magenta dashed lines and
the magenta pentagon indicate the measurements by Cameron et al. (2023) and
Senchyna et al. (2023; gas only), respectively. Mrk996, a Wolf-Rayet galaxy,
is represented by the black diamonds (Senchyna et al., 2023; Berg et al.,
2016). Local galaxies (Berg et al., 2016, 2019; Izotov et al., 2006) and
Galactic H ii regions (García-Rojas & Esteban, 2007) are shown by the gray
dots, while the gray curves are the empirical relation of the Galactic stars
(Nicholls et al., 2017). The purple circles are dwarf stars in a globular
cluster (GC) NGC6752 (Carretta et al., 2005), while those O/H values are taken
from Senchyna et al. (2023). The pink crosses are carbon-enhanced metal-poor
(CEMP) stars (Norris et al., 2013) and nitrogen-enhanced metal-poor (NEMP)
stars (Beveridge & Sneden, 1994). The cyan and brown shaded regions indicate
yields of CCSN (Watanabe et al., 2023) and the equilibrium value of the CNO
cycle (Maeder et al., 2015), respectively. The orange dashed line shows a
chemical evolution model that reproduces emission-line ratios of quasars
(Hamann & Ferland, 1993), while nitrogen enrichment in the model is mainly
caused by asymptotic giant branch (AGB) stars. The blue dashed, green solid,
and yellow dashed-dotted lines illustrate chemical evolution models of Wolf-
Rayet (WR) stars, supermassive stars (SMS), and tidal disruption events (TDE),
respectively, where most of massive stars are assumed to undergo the direct
collapse (Watanabe et al. 2023; Watanabe et al. in prep.). The potential
change of N/O and C/O by dust depletion (Ferland et al., 2013) is indicated by
the length of the small black arrow at the bottom left corner of the bottom
right panel.

The top left panel of Figure 4 shows N/O ratios as a function of metallicity.
We find that CEERS_01019 (red square) and GLASS_150008 (double red circle)
show [N/O$]\gtrsim 0.5$. These 2 galaxies are the second and third examples of
super-solar N/O galaxies at $z>6$ after the report of GN-z11. Hereafter we
refer to these 3 galaxies (CEERS_01019, GLASS_150008, and GN-z11) as JWST
N-rich galaxies. The super-solar N/O values are significantly higher than
those of typical local dwarf galaxies, Galactic H ii regions (gray dots), and
typical Galactic stars (gray curve) at a given $12+\log(\rm O/H)$, but
comparable to that of a Wolf-Rayet galaxy, Mrk996.

Contrary to the high N/O ratios, the top right panel of Figure 4 illustrates
that CEERS_01019 has a C/O ratio comparable to those of the typical local
galaxies and stars (gray). GLASS_150008 (larger double red circle) shows a C/O
ratio even lower than those of the typical local galaxies and stars but
comparable to some of our sample galaxies (white circle).

This discrepancy between the high N/O and low C/O ratios is illustrated by the
bottom left panel of Figure 4, which shows the C/N ratios as a function of
metallicity. We find that all the JWST N-rich galaxies have [C/N] values less
than $\sim-1$, which are significantly lower than those of the typical local
galaxies and stars.

Such low C/N ratios are expected to be observed in nitrogen-loud quasars
(e.g., Batra & Baldwin, 2014), while those O/H ratios have not fully been
investigated. Alternatively, we plot a chemical evolution model of Hamann &
Ferland (1993) reproducing observed emission-line ratios of quasars. Nitrogen
in the model is mainly enriched by AGB stars. Although the quasar model with
AGB stars can produce low C/N ratios down to [C/N$]\sim-1$, it does not
reproduce the (relatively-)low O/H ratios of the JWST N-rich galaxies
simultaneously. This is because AGB stars require much delay time ($\gtrsim 1$
Gyr) to decrease the C/N ratio to [C/N$]\sim-1$, which increases O/H ratios
too much. This result also suggests that the JWST N-rich galaxies are not
likely to be nitrogen-loud quasars whose nitrogen is enriched by AGB stars.

The bottom right panel of Figure 4 shows the relations between N/O and C/O. We
confirm that the JWST N-rich galaxies have N/O ratios higher than those of the
typical local galaxies and stars at a given C/O ratios, which indicate that
only nitrogen is selectively enriched in the JWST N-rich galaxies with respect
to the typical local galaxies and stars.

###  4.2 Origin of Low C/N

Figure 5: Ne/O ratio as a function of $12+\log(\rm O/H)$ (left) and redshift
(right). The symbols are the same as those in the top right panel of Figure 4,
while our sample galaxies with moderately high Ne/O ratios are shown in semi-
transparent for presentation purposes. The gray shaded region in the right
panel corresponds to the 16th and 84th percentiles of the Ne/O distribution of
the local dwarf galaxies (Izotov et al., 2006; Kojima et al., 2021; Isobe et
al., 2022). The blue curves represent chemical evolution models (Watanabe et
al., 2023) with different metallicities. The cyan lines denote isochrones of
the chemical evolution models at the model ages of $10^{6.7}$, $10^{6.8}$,
$10^{6.9}$, and $10^{7.0}$ year that correspond to the lifetime of progenitor
stars with 40, 30, 25, and 20 $M_{\odot}$, respectively (Portinari et al.,
1998).

In the bottom panels of Figure 4, we plot the abundance ratios processed by
the CNO cycle at the equilibrium state (Maeder et al., 2015) by the brown
shaded regions, while core-collapse supernova yields (Watanabe et al., 2023)
are shown by the cyan shaded regions. In contrast to the typical local
galaxies and stars, we find that the N-rich galaxies have the abundance ratios
similar to those of the equilibrium state of the CNO cycle rather than those
of the CCSN yields. Thus, we need scenarios that can enrich gaseous materials
originating from the CNO cycle while suppressing oxygen supply by CCSNe in the
relatively low-metallicity environments.

In fact, these requirements are met by the scenarios presented to explain
nitrogen enrichment of GN-z11 such as Wolf-Rayet (WR) stars (Senchyna et al.,
2023; Watanabe et al., 2023), supermassive stars (SMS; Charbonnel et al. 2023;
Nagele & Umeda 2023), and tidal disruption events (TDE; Cameron et al. 2023).
In all the 3 scenarios, CNO-cycle material can be ejected from the outermost
hydrogen-burning layer of stars via stellar winds (for the WR star and SMS
scenarios) or gravitational interactions with black holes (for the TDE
scenario). Low-metallicity massive stars can also undergo direct collapse
(i.e., no ejecta) due to the low opacity of the hydrogen envelope, resulting
in a heavier helium core during collapse (Heger et al., 2003). The blue
dashed, green solid, and yellow dashed-dotted lines in the bottom panel of
Figure 4 represent chemical evolution models of WR stars, SMSs, and TDEs,
respectively, all of which are constructed by Watanabe et al. (2023) and
Watanabe et al. (in prep.). The WR star and SMS models include the yields of
WR stars with 25–120 $M_{\odot}$ (Limongi & Chieffi, 2018) and SMSs with
$10^{5}\ M_{\odot}$ (Nagele & Umeda, 2023), respectively. The TDE model
includes TDE yields calculated by Watanabe et al. (in prep.) with the
nucleosynthesis code of Tominaga et al. (2007), assuming the destruction of
stars with 9–40 $M_{\odot}$. Note that 100%, 99.99%, and 90% of massive stars
are assumed to undergo direct collapse (i.e., no ejecta) in the WR star, SMS,
and TDE models, respectively, and that the remaining massive stars explode as
CCSNe whose yields are taken from Nomoto et al. (2013). We confirm that all
the 3 models can reproduce the observed N/O and C/O ratios of the JWST N-rich
galaxies. We have also checked that the N/O and C/O ratios of the wind yields
from the $10^{3}\ M_{\odot}$ SMS (Nagele & Umeda, 2023) are also comparable to
those of the JWST N-rich galaxies. These findings suggest the presence of
high-$z$ galaxies affected by WR stars, SMSs, and/or TDEs with frequent direct
collapses.

###  4.3 Local Object with Similar C, N, O Abundances

We find that Mrk996, a local metal-poor dwarf star-forming galaxy with a high
value of $\log(\rm N/O)\sim 0$ (Senchyna et al., 2023), has a low C/N value
similar to the JWST N-rich galaxies (bottom left panel of Figure 4). The
bottom right panel of Figure 4 also shows that Mrk996 coincides with the JWST
N-rich galaxies on the N/O vs. C/O plane. Mrk996 is known to host WR stars
evidenced by the presence of the WR bumps (e.g., Telles et al., 2014). Similar
to Mrk996, the JWST N-rich galaxies may also host WR stars.

We also identify some of carbon-enhanced metal-poor (CEMP) stars (Norris et
al., 2013) and nitrogen-enhanced metal-poor (NEMP) stars (Beveridge & Sneden,
1994) in the Galactic halo with the C/N vs. O/H and N/O vs. C/O relations
comparable to those of the JWST N-rich galaxies as shown by the pink crosses
in Figure 4. It should be noted that some of such stars are giants undergoing
the conversion of C to N in their hydrogen layers.

Moreover, the purple circles in the bottom panels of Figure 4 indicate dwarf
stars in a globular cluster (GC) NGC6752, which have been reported to have
high N/O ratios similar to GN-z11 (Senchyna et al., 2023). We find that the GC
dwarf stars also have low C/N ratios and N/O vs. C/O relations similar to
those of the JWST N-rich galaxies. As surface abundance of dwarf stars are
likely unevovled, the gas composition at the time of GC formation may also
have similar abundance ratios to the JWST N-rich galaxies. This implies that
the JWST N-rich galaxies may be progenitors of GCs.

##  5 Neon Abundance

###  5.1 Result

The left panel of Figure 5 shows Ne/O ratios of our sample galaxies (red) as a
function of $12+\log(\rm O/H)$. Although the majority of our sample galaxies
have Ne/O ratios comparable to those of local dwarf galaxies (gray dot; Izotov
et al. 2006; Kojima et al. 2021; Isobe et al. 2022), 4 of our sample galaxies
(GLASS_100003, CEERS_00698, CEERS_01143, and CEERS_01149) have $\log(\rm
Ne/O)$ values of $<-1.0$, which are significantly lower than those of the
local dwarf galaxies and the other our sample galaxies. The right panel of
Figure 5 illustrates the Ne/O ratios as a function of redshift. Interestingly,
all the 4 galaxies with the low Ne/O ratios are located at $z>6$.

###  5.2 Origin of Low Ne/O

To explore the origins of the low Ne/O, we compare these data points with
(Watanabe et al., 2023)’s chemical evolution models accumulating the CCSN
yields calculated by the nucleosynthesis code of Tominaga et al. (2007) under
the assumption of Kroupa (2001) IMF. We assume the metallicity evolution of
the Milky Way (Suzuki & Maeda, 2018), and convert the maximum stellar ages of
the models (Agemod) to $12+\log(\rm O/H)$. The blue curves in Figure 5
correspond to the chemical evolution models that accumulate CCSN ejecta with
fixed metallicities of $Z=0$, 0.001, and 0.004. We also illustrate isochrones
of the models at $\log(\rm Age_{\rm mod}/yr)=6.7$, 6.8, 6.9, and 7.0, which
correspond to lifetimes of stars with 40, 30, 25, and 20 $M_{\odot}$,
respectively (Portinari et al., 1998). As shown in the top left panel of
Figure 5, the models predict an increase in Ne/O with Agemod. This increase
can originate from the prediction that CCSNe with more massive progenitors
have higher temperatures in the carbon-burning layer (e.g., Woosley & Janka,
2005), which reduces neon abundance.

As shown in the top left panel of Figure 5, we find that the models with
$\log(\rm Age_{\rm mod}/yr)\lesssim 6.8$ (i.e., Age${}_{\rm mod}\lesssim 6$
Myr) reproduce the low Ne/O ratios of the 4 galaxies. This suggests that the 4
galaxies have an abundance population of massive stars beyond $\sim 30\
M_{\odot}$. This implication agrees with the fact that the 4 galaxies are
located at higher redshifts beyond 6, because galaxies at higher redshifts are
expected to be younger and/or have top-heavy IMFs due to the metal-poorer
environment.

##  6 Summary

We present N/O, C/O, C/N, and Ne/O ratios of 70 star-forming galaxies at
$z\sim 4$–10, observed by the JWST/NIRSpec ERO, GLASS, and CEERS programs. We
derive these abundance ratios from emission-line ratios of similar ionization
energies, accounting for both stellar and AGN radiation, particularly for
CEERS_01019 and GN-z11, which exhibit AGN signatures of high-ionization lines.
Our findings are summarized below:

  * •

Among the 70 galaxies, we have pinpointed 2 galaxies with unique
characteristics: CEERS_01019 at $z=8.68$ and GLASS_150008 at $z=6.23$. We find
the low C/N and high N/O ratios these galaxies, which are also found in
GN-z11, closely align with the equilibrium of the CNO cycle. This suggests
that these 3 galaxies have the gas component dominated by metals processed by
the CNO cycle.

  * •

The C/O and N/O ratios of these 3 galaxies are reproduced by chemical
evolution models dominated by CNO-cycle materials involving Wolf-Rayet stars,
supermassive stars, and tidal disruption events. Interestingly, these
scenarios would require frequent direct collapses.

  * •

On the C/N vs. O/H plane, these 3 galaxies do not align with Galactic H ii
regions, typical Galactic stars, typical star-forming galaxies, or nitrogen-
loud quasars whose nitrogen is assumed to originate from AGB stars. They do,
however, coincide with GC dwarf stars. We may see the site of GC formation in
these 3 galaxies.

  * •

We identify 4 galaxies with very low Ne/O, with $\log(\rm Ne/O)<-1.0$. This
low ratio indicates the abundance of massive ($\gtrsim 30\ M_{\odot}$) stars.

We thank T. Nagao, A.K. Inoue, and S. Aoyama for having useful discussions. We
are grateful to staff of the James Webb Space Telescope Help Desk for letting
us know useful information. This work is based on observations made with the
NASA/ESA/CSA James Webb Space Telescope. The data were obtained from the
Mikulski Archive for Space Telescopes at the Space Telescope Science
Institute, which is operated by the Association of Universities for Research
in Astronomy, Inc., under NASA contract NAS 5-03127 for JWST. These
observations are associated with programs 1125, 2736, 1324, and 1345. The
authors acknowledge the ERO, GLASS, and CEERS teams led by Klaus M.
Pontoppidan, Tommaso Treu, and Steven L. Finkelstein, respectively, for
developing their observing programs with a zero-exclusive-access period. This
work is based on observations taken by the CANDELS Multi-Cycle Treasury
Program with the NASA/ESA HST, which is operated by the Association of
Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555.
This work was supported by the joint research program of the Institute for
Cosmic Ray Research (ICRR), University of Tokyo. This publication is based
upon work supported by the World Premier International Research Center
Initiative (WPI Initiative), MEXT, Japan. Y.I., M.O., K.N., and Y.H. are
supported by JSPS KAKENHI Grant Nos. 21J20785, 20H00180/21H04467, 20K22373,
and 21K13953, respectively. H.Y. is supported by MEXT/JSPS KAKENHI (21H04489)
and JST FOREST Program (JP-MJFR202Z). This research was supported by a grant
from the Hayakawa Satio Fund awarded by the Astronomical Society of Japan.

##  Appendix A Table of Metallicity and Abundance Ratios

Metallicity and Abundance Ratios of $z\sim 4$–9 Galaxies Based on the Stellar Radiation (Sections 3.4 and B) |  |  |  |  |  |  |  |  |
---|---|---|---|---|---|---|---|---|---
Table 2 continued.
ID | Redshift | $T_{\rm e}$ | $E(B-V)$ | $12+\log(\rm O/H)$ | $\log(\rm Ne/O)$ | $\log(\rm C/O)$ | $\log(\rm N/O)$ | $\log(\rm Ar/O)$ | $\log(\rm S/O)$
|  |  $10^{4}$ K | mag |  |  |  |  |  |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10)
Notes. (1) ID. (2) Redshift. (3) Electron temperature. a: Based on [O
iii]$\lambda$4363 derived by Nakajima et al. (2023); b: Based on [O
iii]$\lambda$4363 but not reported by Nakajima et al. (2023); c: Based on O
iii]$\lambda$1666; d: Assumption. (4) Color excess. e: Assumed to be the
median $E(B-V)$ of the other our sample galaxies. (5) Metallicity. f: [O ii]
$S/N\leq 3$; g: [O ii] $S/N<3$; h: [O ii] out of wavelength coverage. (6)–(10)
Abundance ratios. We include uncertainties of the assumptions of $T_{\rm e}$
and [O ii] into the errors of $12+\log(\rm O/H)$ and the abundance ratios by
performing Monte Carlo simulations (Section 3.4).
ID | Redshift | $T_{\rm e}$ | $E(B-V)$ | $12+\log(\rm O/H)$ | $\log(\rm Ne/O)$ | $\log(\rm C/O)$ | $\log(\rm N/O)$ | $\log(\rm Ar/O)$ | $\log(\rm S/O)$
|  |  $10^{4}$ K | mag |  |  |  |  |  |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10)
ERO_04590 | 8.4956 | $2.08\pm 0.26^{\rm a}$ | 0.12 | $7.19^{+0.13\,{\rm g}}_{-0.10}$ | $<-0.22$ | $-0.53^{+0.17}_{-0.15}$ | $\cdots$ | $\cdots$ | $\cdots$
ERO_05144 | 6.3784 | $1.64\pm 0.16^{\rm a}$ | 0.00 | $7.79^{+0.11\,{\rm f}}_{-0.10}$ | $-0.74^{+0.05}_{-0.06}$ | $\cdots$ | $\cdots$ | $\cdots$ | $<-0.87$
ERO_06355 | 7.6651 | $1.13\pm 0.08^{\rm a}$ | 0.13 | $8.31^{+0.10\,{\rm f}}_{-0.09}$ | $-0.74^{+0.02}_{-0.02}$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$
ERO_08140 | 5.2751 | $1.5\pm 0.5^{\rm d}$ | 0.05 | $7.77^{+0.45\,{\rm f}}_{-0.31}$ | $<-0.90$ | $\cdots$ | $\cdots$ | $<-1.75$ | $-1.64^{+0.13}_{-0.10}$
ERO_10612 | 7.6601 | $2.04\pm 0.16^{\rm a}$ | 0.00 | $7.62^{+0.08\,{\rm f}}_{-0.06}$ | $-0.74^{+0.02}_{-0.02}$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$
GLASS_10000 | 7.8809 | $1.83^{+0.32\,{\rm b}}_{-0.32}$ | 0.00 | $7.63^{+0.18\,{\rm g}}_{-0.13}$ | $-0.71^{+0.04}_{-0.06}$ | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$
GLASS_100001 | 7.8737 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.90^{+0.52\,{\rm f}}_{-0.30}$ | $<-0.83$ | $\cdots$ | $<0.46$ | $\cdots$ | $\cdots$
GLASS_100003 | 7.8773 | $1.98\pm 0.37^{\rm a}$ | 0.00 | $7.67^{+0.20\,{\rm f}}_{-0.15}$ | $-1.05^{+0.09}_{-0.10}$ | $\cdots$ | $<-0.30$ | $\cdots$ | $\cdots$
GLASS_10021 | 7.2863 | $1.66\pm 0.17^{\rm a}$ | -0.00 | $7.87^{+0.12\,{\rm f}}_{-0.10}$ | $-0.87^{+0.04}_{-0.03}$ | $\cdots$ | $<-0.40$ | $\cdots$ | $\cdots$
GLASS_110000 | 5.7631 | $1.5\pm 0.5^{\rm d}$ | 0.13 | $7.96^{+0.50\,{\rm f}}_{-0.29}$ | $-0.56^{+0.11}_{-0.07}$ | $<-0.66$ | $<0.25$ | $<-1.92$ | $<-1.57$
GLASS_150008 | 6.2291 | $1.93^{+0.12\,{\rm c}}_{-0.12}$ | 0.07e | $7.65^{+0.14\,{\rm h}}_{-0.08}$ | $\cdots$ | $-1.08^{+0.06}_{-0.14}$ | $-0.40^{+0.05}_{-0.07}$ | $\cdots$ | $\cdots$
GLASS_150029 | 4.5837 | $1.76\pm 0.14^{\rm a}$ | 0.16 | $7.73^{+0.09\,{\rm f}}_{-0.08}$ | $\cdots$ | $-0.84^{+0.10}_{-0.09}$ | $\cdots$ | $<-2.63$ | $-1.69^{+0.09}_{-0.11}$
GLASS_160122 | 5.3312 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.92^{+0.54\,{\rm f}}_{-0.28}$ | $-0.62^{+0.10}_{-0.06}$ | $-0.72^{+0.63}_{-0.31}$ | $<-0.03$ | $<-1.96$ | $<-1.26$
GLASS_160133 | 4.0151 | $1.48\pm 0.07^{\rm a}$ | 0.13 | $7.95^{+0.05\,{\rm f}}_{-0.05}$ | $-0.78^{+0.02}_{-0.02}$ | $\cdots$ | $\cdots$ | $\cdots$ | $<-2.09$
GLASS_40066 | 4.0197 | $1.66^{+0.14\,{\rm b}}_{-0.14}$ | 0.23 | $7.84^{+0.17\,{\rm h}}_{-0.10}$ | $-0.69^{+0.02}_{-0.15}$ | $\cdots$ | $\cdots$ | $<-2.76$ | $-1.94^{+0.05}_{-0.60}$
GLASS_50002 | 5.1333 | $1.5\pm 0.5^{\rm d}$ | 0.52 | $8.32^{+0.51\,{\rm f}}_{-0.29}$ | $<-0.37$ | $\cdots$ | $\cdots$ | $<-1.95$ | $<-1.78$
GLASS_50038 | 5.7720 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.94^{+0.48\,{\rm f}}_{-0.29}$ | $-0.87^{+0.13}_{-0.13}$ | $<-1.00$ | $<-0.15$ | $\cdots$ | $\cdots$
GLASS_80029 | 3.9510 | $1.5\pm 0.5^{\rm d}$ | 0.09 | $7.91^{+0.54\,{\rm h}}_{-0.33}$ | $-0.72^{+0.08}_{-0.14}$ | $\cdots$ | $\cdots$ | $-2.19^{+0.12}_{-0.38}$ | $-1.62^{+0.08}_{-0.62}$
GLASS_80070 | 4.7969 | $1.5\pm 0.5^{\rm d}$ | 0.11 | $7.79^{+0.55\,{\rm f}}_{-0.29}$ | $-0.73^{+0.11}_{-0.06}$ | $-0.71^{+0.65}_{-0.32}$ | $<0.18$ | $<-2.26$ | $<-1.80$
CEERS_00323 | 5.6657 | $1.5\pm 0.5^{\rm d}$ | 0.07 | $7.61^{+0.46\,{\rm g}}_{-0.30}$ | $-0.63^{+0.09}_{-0.10}$ | $<-0.42$ | $<0.42$ | $<-1.28$ | $<-0.55$
CEERS_00355 | 6.0998 | $1.5\pm 0.5^{\rm d}$ | 0.01 | $7.88^{+0.59\,{\rm f}}_{-0.27}$ | $<-0.94$ | $<-0.72$ | $<0.22$ | $<-1.60$ | $<-1.29$
CEERS_00397 | 6.0005 | $1.52^{+0.16\,{\rm b}}_{-0.17}$ | 0.10 | $7.99^{+0.12\,{\rm f}}_{-0.10}$ | $-0.76^{+0.03}_{-0.04}$ | $-0.65^{+0.14}_{-0.12}$ | $<-0.20$ | $<-2.18$ | $<-1.80$
CEERS_00403 | 5.7609 | $1.5\pm 0.5^{\rm d}$ | 0.53 | $8.12^{+0.46\,{\rm f}}_{-0.28}$ | $<-0.62$ | $<0.40$ | $<1.46$ | $<-2.03$ | $-1.76^{+0.14}_{-0.14}$
CEERS_00407 | 7.0291 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.83^{+0.52\,{\rm g}}_{-0.30}$ | $<-0.42$ | $<-0.27$ | $<0.58$ | $\cdots$ | $\cdots$
CEERS_00515 | 5.6644 | $1.5\pm 0.5^{\rm d}$ | 0.27 | $7.68^{+0.47\,{\rm g}}_{-0.30}$ | $<-0.35$ | $<0.25$ | $<1.26$ | $<-1.21$ | $<-0.72$
CEERS_00603 | 6.0568 | $1.5\pm 0.5^{\rm d}$ | 0.04 | $7.92^{+0.49\,{\rm f}}_{-0.28}$ | $-0.80^{+0.13}_{-0.12}$ | $<-0.65$ | $<0.18$ | $<-1.50$ | $<-1.09$
CEERS_00618 | 6.0499 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.55^{+0.51\,{\rm g}}_{-0.30}$ | $<-0.31$ | $<-0.14$ | $<0.67$ | $<-0.86$ | $<-0.39$
CEERS_00670 | 5.8045 | $1.5\pm 0.5^{\rm d}$ | 0.34 | $7.87^{+0.55\,{\rm g}}_{-0.29}$ | $<-0.39$ | $<0.29$ | $<1.31$ | $<-1.29$ | $\cdots$
CEERS_00689 | 7.5458 | $1.5\pm 0.5^{\rm d}$ | 0.07e | $7.95^{+0.47\,{\rm f}}_{-0.30}$ | $<-0.80$ | $<-0.18$ | $<0.32$ | $\cdots$ | $\cdots$
CEERS_00698 | 7.4708 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.97^{+0.50\,{\rm f}}_{-0.30}$ | $-1.01^{+0.11}_{-0.09}$ | $<-0.91$ | $<-0.16$ | $\cdots$ | $\cdots$
CEERS_00707 | 4.8964 | $1.5\pm 0.5^{\rm d}$ | 0.24 | $8.03^{+0.56\,{\rm h}}_{-0.34}$ | $-0.76^{+0.10}_{-0.12}$ | $<-0.82$ | $<0.25$ | $<-2.42$ | $-2.00^{+0.09}_{-0.59}$
CEERS_00717 | 6.9315 | $1.5\pm 0.5^{\rm d}$ | 0.21 | $7.95^{+0.59\,{\rm f}}_{-0.28}$ | $-0.77^{+0.15}_{-0.12}$ | $<-0.31$ | $<0.60$ | $\cdots$ | $\cdots$
CEERS_00792 | 6.2572 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $8.10^{+0.58\,{\rm f}}_{-0.29}$ | $<-0.87$ | $<-0.67$ | $<0.16$ | $<-1.54$ | $<-1.18$
CEERS_01019 | 8.6791 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.94^{+0.46\,{\rm f}}_{-0.31}$ | $-0.77^{+0.11}_{-0.10}$ | $<-1.04$ | $<0.14$ | $\cdots$ | $\cdots$
CEERS_01023 | 7.7764 | $1.5\pm 0.5^{\rm d}$ | 0.07e | $7.67^{+0.55\,{\rm g}}_{-0.27}$ | $<-0.41$ | $\cdots$ | $<0.81$ | $\cdots$ | $\cdots$
CEERS_01025 | 8.7140 | $1.5\pm 0.5^{\rm d}$ | 0.07e | $7.99^{+0.53\,{\rm f}}_{-0.28}$ | $-0.79^{+0.14}_{-0.14}$ | $\cdots$ | $<0.63$ | $\cdots$ | $\cdots$
CEERS_01027 | 7.8197 | $1.79^{+0.28\,{\rm b}}_{-0.27}$ | 0.25 | $7.67^{+0.16\,{\rm f}}_{-0.12}$ | $-0.57^{+0.03}_{-0.04}$ | $<-0.60$ | $<0.12$ | $\cdots$ | $\cdots$
CEERS_01029 | 8.6100 | $1.5\pm 0.5^{\rm d}$ | 0.07e | $7.92^{+0.54\,{\rm f}}_{-0.29}$ | $<-0.64$ | $<-0.16$ | $<0.71$ | $\cdots$ | $\cdots$
CEERS_01038 | 7.1945 | $1.5\pm 0.5^{\rm d}$ | 0.07e | $7.64^{+0.48\,{\rm g}}_{-0.28}$ | $<-0.40$ | $<0.00$ | $<0.80$ | $\cdots$ | $\cdots$
CEERS_01115 | 6.3002 | $1.5\pm 0.5^{\rm d}$ | 0.06 | $7.70^{+0.57\,{\rm g}}_{-0.28}$ | $<-0.66$ | $<-0.47$ | $<0.46$ | $\cdots$ | $<-0.34$
CEERS_01143 | 6.9266 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.99^{+0.49\,{\rm f}}_{-0.28}$ | $-1.02^{+0.10}_{-0.09}$ | $<-1.05$ | $<-0.19$ | $\cdots$ | $\cdots$
CEERS_01149 | 8.1750 | $1.5\pm 0.5^{\rm d}$ | 0.02 | $7.87^{+0.51\,{\rm g}}_{-0.29}$ | $-1.06^{+0.12}_{-0.12}$ | $-0.61^{+0.60}_{-0.36}$ | $<0.14$ | $\cdots$ | $\cdots$
CEERS_01160 | 6.5667 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.97^{+0.47\,{\rm f}}_{-0.28}$ | $<-0.62$ | $<-0.62$ | $<0.25$ | $\cdots$ | $\cdots$
CEERS_01173 | 4.9957 | $1.5\pm 0.5^{\rm d}$ | 0.31 | $8.12^{+0.53\,{\rm f}}_{-0.27}$ | $-0.75^{+0.14}_{-0.09}$ | $<-0.27$ | $<0.89$ | $<-2.15$ | $-1.83^{+0.12}_{-0.11}$
CEERS_01236 | 4.4842 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.83^{+0.56\,{\rm f}}_{-0.27}$ | $<-0.40$ | $<-0.09$ | $\cdots$ | $<-1.43$ | $<-1.33$
CEERS_01289 | 4.8798 | $1.5\pm 0.5^{\rm d}$ | 0.07 | $8.02^{+0.52\,{\rm f}}_{-0.28}$ | $<-0.77$ | $<-0.56$ | $<0.43$ | $<-1.81$ | $<-1.33$
CEERS_01294 | 4.9987 | $1.5\pm 0.5^{\rm d}$ | 0.15 | $8.12^{+0.46\,{\rm f}}_{-0.30}$ | $-0.82^{+0.13}_{-0.12}$ | $<-0.45$ | $<0.56$ | $<-1.84$ | $<-1.50$
CEERS_01324 | 5.0074 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.68^{+0.54\,{\rm g}}_{-0.33}$ | $<-0.29$ | $<-0.19$ | $<0.79$ | $<-1.10$ | $<-0.69$
CEERS_01358 | 5.5038 | $1.5\pm 0.5^{\rm d}$ | 0.38 | $8.02^{+0.59\,{\rm g}}_{-0.28}$ | $-0.41^{+0.11}_{-0.09}$ | $<0.31$ | $<1.25$ | $<-1.43$ | $<-0.47$
CEERS_01365 | 4.3074 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.71^{+0.49\,{\rm g}}_{-0.28}$ | $<-0.49$ | $\cdots$ | $\cdots$ | $<-1.26$ | $<-0.76$
CEERS_01374 | 4.9999 | $1.5\pm 0.5^{\rm d}$ | 0.13 | $8.06^{+0.50\,{\rm f}}_{-0.27}$ | $-0.86^{+0.10}_{-0.05}$ | $<-1.17$ | $<-0.20$ | $\cdots$ | $-1.84^{+0.08}_{-0.06}$
CEERS_01401 | 5.3749 | $1.5\pm 0.5^{\rm d}$ | 0.07 | $7.97^{+0.53\,{\rm f}}_{-0.29}$ | $-0.87^{+0.14}_{-0.10}$ | $<-0.76$ | $<0.17$ | $<-1.86$ | $<-1.73$
CEERS_01449 | 4.7519 | $1.5\pm 0.5^{\rm d}$ | 0.20 | $7.92^{+0.52\,{\rm f}}_{-0.28}$ | $-0.70^{+0.12}_{-0.06}$ | $<-0.91$ | $<0.11$ | $-2.25^{+0.11}_{-0.25}$ | $-1.78^{+0.07}_{-0.06}$
CEERS_01465 | 5.2692 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.65^{+0.52\,{\rm f}}_{-0.31}$ | $-0.48^{+0.15}_{-0.12}$ | $<-0.22$ | $<0.69$ | $<-1.32$ | $<-1.21$
CEERS_01536 | 5.0337 | $2.25\pm 0.43^{\rm a}$ | 0.08 | $7.62^{+0.22\,{\rm f}}_{-0.12}$ | $-0.84^{+0.07}_{-0.07}$ | $<-1.16$ | $<-0.33$ | $<-1.75$ | $<-1.29$
CEERS_01539 | 4.8841 | $1.5\pm 0.5^{\rm d}$ | 0.27 | $7.90^{+0.50\,{\rm f}}_{-0.29}$ | $-0.70^{+0.12}_{-0.07}$ | $<-0.27$ | $<0.74$ | $<-2.18$ | $-1.85^{+0.09}_{-0.07}$
CEERS_01544 | 4.1901 | $1.5\pm 0.5^{\rm d}$ | 0.05 | $7.76^{+0.49\,{\rm g}}_{-0.30}$ | $<-0.31$ | $<-0.01$ | $\cdots$ | $<-0.41$ | $<-0.76$
CEERS_01561 | 6.1967 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.65^{+0.52\,{\rm g}}_{-0.28}$ | $<-0.69$ | $<-0.51$ | $<0.33$ | $<-1.09$ | $<-0.33$
CEERS_01605 | 4.6309 | $1.5\pm 0.5^{\rm d}$ | 0.07 | $7.86^{+0.48\,{\rm f}}_{-0.30}$ | $<-0.75$ | $\cdots$ | $<0.55$ | $<-1.79$ | $<-1.71$
CEERS_01658 | 4.6038 | $1.5\pm 0.5^{\rm d}$ | 0.24 | $7.98^{+0.51\,{\rm f}}_{-0.30}$ | $\cdots$ | $<-0.42$ | $\cdots$ | $<-2.03$ | $\cdots$
CEERS_01677 | 5.8670 | $1.5\pm 0.5^{\rm d}$ | 0.00 | $7.91^{+0.56\,{\rm f}}_{-0.31}$ | $-0.23^{+0.12}_{-0.06}$ | $<-0.21$ | $<0.36$ | $<-1.42$ | $<-1.33$
CEERS_02174 | 5.3030 | $1.5\pm 0.5^{\rm d}$ | 0.21 | $8.08^{+0.53\,{\rm f}}_{-0.29}$ | $<-0.39$ | $\cdots$ | $<0.95$ | $<-1.60$ | $<-1.54$
CEERS_02782 | 5.2408 | $1.5\pm 0.5^{\rm d}$ | 0.32 | $7.74^{+0.55\,{\rm g}}_{-0.28}$ | $-0.59^{+0.11}_{-0.10}$ | $<-0.01$ | $<0.99$ | $<-1.49$ | $<-0.69$

##  Appendix B Argon and Sulfur Abundances

We obtain Ar/O and S/O ratios in the same way as the Ne/O and N/O
measurements. We derive ionic abundance ratios of Ar2+/H+ and S+/H+ from [Ar
iii]$\lambda$7136 and [S ii]$\lambda\lambda$6716,6731 lines respectively. We
use $T_{\rm e}$(O ii) for S+/H+ and $T_{\rm e}$(S iii) for Ar2+/H+, where
$T_{\rm e}$(S iii) is calculated from $T_{\rm e}$(O iii) and the empirical
relation of Garnett (1992). Because the ionization energy of S+ ion is lower
than that of H+, we derive ICF(S+) calculated until the photo-dissociation
region (PDR). Figure 6 illustrates Ar/O and S/O ratios of our sample galaxies
as a function of metallicity. In addition to the CCSN models (blue), we plot
Watanabe et al. (2023)’s hypernova (HN) models (green). The observed Ar/O and
S/O ratios can be explained by the combination of CCSN and HN yields (see
Watanabe et al. 2023 for more details).

Figure 6: Ar/O and S/O as a function of metallicity. The symbols are the same
as those in Figure 4, while we add Watanabe et al. (2023)’s chemical evolution
models of hypernovae with different metallicities (green) and isochrones
(light green).

## References

  * Aggarwal & Keenan (1999) Aggarwal, K. M., & Keenan, F. P. 1999, ApJS, 123, 311, doi: 10.1086/313232
  * Arellano-Córdova et al. (2022) Arellano-Córdova, K. Z., Berg, D. A., Chisholm, J., et al. 2022, ApJ, 940, L23, doi: 10.3847/2041-8213/ac9ab2
  * Asplund et al. (2021) Asplund, M., Amarsi, A. M., & Grevesse, N. 2021, A&A, 653, A141, doi: 10.1051/0004-6361/202140445
  * Astropy Collaboration et al. (2013) Astropy Collaboration, Robitaille, T. P., Tollerud, E. J., et al. 2013, A&A, 558, A33, doi: 10.1051/0004-6361/201322068
  * Batra & Baldwin (2014) Batra, N. D., & Baldwin, J. A. 2014, MNRAS, 439, 771, doi: 10.1093/mnras/stu007
  * Berg et al. (2019) Berg, D. A., Erb, D. K., Henry, R. B. C., Skillman, E. D., & McQuinn, K. B. W. 2019, ApJ, 874, 93, doi: 10.3847/1538-4357/ab020a
  * Berg et al. (2016) Berg, D. A., Skillman, E. D., Henry, R. B. C., Erb, D. K., & Carigi, L. 2016, ApJ, 827, 126, doi: 10.3847/0004-637X/827/2/126
  * Berrington et al. (1985) Berrington, K. A., Burke, P. G., Dufton, P. L., & Kingston, A. E. 1985, Atomic Data and Nuclear Data Tables, 33, 195, doi: 10.1016/0092-640X(85)90001-4
  * Beveridge & Sneden (1994) Beveridge, R. C., & Sneden, C. 1994, AJ, 108, 285, doi: 10.1086/117068
  * Blum & Pradhan (1992) Blum, R. D., & Pradhan, A. K. 1992, ApJS, 80, 425, doi: 10.1086/191670
  * Bunker et al. (2023) Bunker, A. J., Saxena, A., Cameron, A. J., et al. 2023, arXiv e-prints, arXiv:2302.07256, doi: 10.48550/arXiv.2302.07256
  * Calzetti et al. (2000) Calzetti, D., Armus, L., Bohlin, R. C., et al. 2000, ApJ, 533, 682, doi: 10.1086/308692
  * Cameron et al. (2023) Cameron, A. J., Katz, H., Rey, M. P., & Saxena, A. 2023, MNRAS, doi: 10.1093/mnras/stad1579
  * Carretta et al. (2005) Carretta, E., Gratton, R. G., Lucatello, S., Bragaglia, A., & Bonifacio, P. 2005, A&A, 433, 597, doi: 10.1051/0004-6361:20041892
  * Charbonnel et al. (2023) Charbonnel, C., Schaerer, D., Prantzos, N., et al. 2023, A&A, 673, L7, doi: 10.1051/0004-6361/202346410
  * Elvis et al. (1994) Elvis, M., Wilkes, B. J., McDowell, J. C., et al. 1994, ApJS, 95, 1, doi: 10.1086/192093
  * Ferland et al. (2013) Ferland, G. J., Porter, R. L., van Hoof, P. A. M., et al. 2013, Rev. Mexicana Astron. Astrofis., 49, 137.  https://arxiv.org/abs/1302.4485
  * Finkelstein et al. (2022) Finkelstein, S. L., Bagley, M. B., Ferguson, H. C., et al. 2022, arXiv e-prints, arXiv:2211.05792.  https://arxiv.org/abs/2211.05792
  * Francis (1993) Francis, P. J. 1993, ApJ, 407, 519, doi: 10.1086/172533
  * Froese Fischer & Tachiev (2004) Froese Fischer, C., & Tachiev, G. 2004, Atomic Data and Nuclear Data Tables, 87, 1, doi: 10.1016/j.adt.2004.02.001
  * Galavis et al. (1998) Galavis, M. E., Mendoza, C., & Zeippen, C. J. 1998, A&AS, 131, 499, doi: 10.1051/aas:1998435
  * García-Rojas & Esteban (2007) García-Rojas, J., & Esteban, C. 2007, ApJ, 670, 457, doi: 10.1086/521871
  * Garnett (1992) Garnett, D. R. 1992, AJ, 103, 1330, doi: 10.1086/116146
  * Hamann & Ferland (1993) Hamann, F., & Ferland, G. 1993, ApJ, 418, 11, doi: 10.1086/173366
  * Heger et al. (2003) Heger, A., Fryer, C. L., Woosley, S. E., Langer, N., & Hartmann, D. H. 2003, ApJ, 591, 288, doi: 10.1086/375341
  * Isobe et al. (2023) Isobe, Y., Ouchi, M., Nakajima, K., et al. 2023, arXiv e-prints, arXiv:2301.06811, doi: 10.48550/arXiv.2301.06811
  * Isobe et al. (2022) Isobe, Y., Ouchi, M., Suzuki, A., et al. 2022, ApJ, 925, 111, doi: 10.3847/1538-4357/ac3509
  * Izotov et al. (2006) Izotov, Y. I., Stasińska, G., Meynet, G., Guseva, N. G., & Thuan, T. X. 2006, A&A, 448, 955, doi: 10.1051/0004-6361:20053763
  * Jones et al. (2023) Jones, T., Sanders, R., Chen, Y., et al. 2023, arXiv e-prints, arXiv:2301.07126, doi: 10.48550/arXiv.2301.07126
  * Kisielius et al. (2009) Kisielius, R., Storey, P. J., Ferland, G. J., & Keenan, F. P. 2009, MNRAS, 397, 903, doi: 10.1111/j.1365-2966.2009.14989.x
  * Kojima et al. (2017) Kojima, T., Ouchi, M., Nakajima, K., et al. 2017, PASJ, 69, 44, doi: 10.1093/pasj/psx017
  * Kojima et al. (2021) Kojima, T., Ouchi, M., Rauch, M., et al. 2021, ApJ, 913, 22, doi: 10.3847/1538-4357/abec3d
  * Kroupa (2001) Kroupa, P. 2001, MNRAS, 322, 231, doi: 10.1046/j.1365-8711.2001.04022.x
  * Larson et al. (2023) Larson, R. L., Finkelstein, S. L., Kocevski, D. D., et al. 2023, arXiv e-prints, arXiv:2303.08918, doi: 10.48550/arXiv.2303.08918
  * Limongi & Chieffi (2018) Limongi, M., & Chieffi, A. 2018, ApJS, 237, 13, doi: 10.3847/1538-4365/aacb24
  * Luridiana et al. (2015) Luridiana, V., Morisset, C., & Shaw, R. A. 2015, A&A, 573, A42, doi: 10.1051/0004-6361/201323152
  * Maeder et al. (2015) Maeder, A., Meynet, G., & Chiappini, C. 2015, A&A, 576, A56, doi: 10.1051/0004-6361/201424153
  * Maiolino et al. (2023) Maiolino, R., Scholtz, J., Witstok, J., et al. 2023, arXiv e-prints, arXiv:2305.12492, doi: 10.48550/arXiv.2305.12492
  * McLaughlin & Bell (2000) McLaughlin, B. M., & Bell, K. L. 2000, JPhB, 33, 597, doi: 10.1088/0953-4075/33/4/301
  * Munoz Burgos et al. (2009) Munoz Burgos, J. M., Loch, S. D., Ballance, C. P., & Boivin, R. F. 2009, A&A, 500, 1253, doi: 10.1051/0004-6361/200911743
  * Nagele & Umeda (2023) Nagele, C., & Umeda, H. 2023, ApJ, 949, L16, doi: 10.3847/2041-8213/acd550
  * Nakajima et al. (2023) Nakajima, K., Ouchi, M., Isobe, Y., et al. 2023, arXiv e-prints, arXiv:2301.12825, doi: 10.48550/arXiv.2301.12825
  * Nicholls et al. (2017) Nicholls, D. C., Sutherland, R. S., Dopita, M. A., Kewley, L. J., & Groves, B. A. 2017, MNRAS, 466, 4403, doi: 10.1093/mnras/stw3235
  * Nomoto et al. (2013) Nomoto, K., Kobayashi, C., & Tominaga, N. 2013, ARA&A, 51, 457, doi: 10.1146/annurev-astro-082812-140956
  * Norris et al. (2013) Norris, J. E., Yong, D., Bessell, M. S., et al. 2013, ApJ, 762, 28, doi: 10.1088/0004-637X/762/1/28
  * Oesch et al. (2016) Oesch, P. A., Brammer, G., van Dokkum, P. G., et al. 2016, ApJ, 819, 129, doi: 10.3847/0004-637x/819/2/129
  * Pilyugin et al. (2012) Pilyugin, L. S., Vílchez, J. M., Mattsson, L., & Thuan, T. X. 2012, MNRAS, 421, 1624, doi: 10.1111/j.1365-2966.2012.20420.x
  * Pontoppidan et al. (2022) Pontoppidan, K. M., Barrientes, J., Blome, C., et al. 2022, ApJ, 936, L14, doi: 10.3847/2041-8213/ac8a4e
  * Portinari et al. (1998) Portinari, L., Chiosi, C., & Bressan, A. 1998, A&A, 334, 505.  https://arxiv.org/abs/astro-ph/9711337
  * Rynkun et al. (2019) Rynkun, P., Gaigalas, G., & Jönsson, P. 2019, A&A, 623, A155, doi: 10.1051/0004-6361/201834931
  * Salpeter (1955) Salpeter, E. E. 1955, ApJ, 121, 161, doi: 10.1086/145971
  * Senchyna et al. (2023) Senchyna, P., Plat, A., Stark, D. P., & Rudie, G. C. 2023, arXiv e-prints, arXiv:2303.04179, doi: 10.48550/arXiv.2303.04179
  * Stanway & Eldridge (2018) Stanway, E. R., & Eldridge, J. J. 2018, MNRAS, 479, 75, doi: 10.1093/mnras/sty1353
  * Steidel et al. (2016) Steidel, C. C., Strom, A. L., Pettini, M., et al. 2016, ApJ, 826, 159, doi: 10.3847/0004-637X/826/2/159
  * Storey & Hummer (1995) Storey, P. J., & Hummer, D. G. 1995, MNRAS, 272, 41, doi: 10.1093/mnras/272.1.41
  * Storey et al. (2014) Storey, P. J., Sochi, T., & Badnell, N. R. 2014, MNRAS, 441, 3028, doi: 10.1093/mnras/stu777
  * Suzuki & Maeda (2018) Suzuki, A., & Maeda, K. 2018, MNRAS, 478, 110, doi: 10.1093/mnras/sty999
  * Tayal & Zatsarinny (2010) Tayal, S. S., & Zatsarinny, O. 2010, ApJS, 188, 32, doi: 10.1088/0067-0049/188/1/32
  * Telles et al. (2014) Telles, E., Thuan, T. X., Izotov, Y. I., & Carrasco, E. R. 2014, A&A, 561, A64, doi: 10.1051/0004-6361/201219270
  * The Astropy Collaboration (2018) The Astropy Collaboration. 2018, astropy v3.1: a core python package for astronomy, 3.1, Zenodo, Zenodo, doi: 10.5281/zenodo.4080996
  * Tominaga et al. (2007) Tominaga, N., Umeda, H., & Nomoto, K. 2007, ApJ, 660, 516, doi: 10.1086/513063
  * Treu et al. (2022) Treu, T., Roberts-Borsani, G., Bradac, M., et al. 2022, ApJ, 935, 110, doi: 10.3847/1538-4357/ac8158
  * Vincenzo et al. (2016) Vincenzo, F., Belfiore, F., Maiolino, R., Matteucci, F., & Ventura, P. 2016, MNRAS, 458, 3466, doi: 10.1093/mnras/stw532
  * Watanabe et al. (2023) Watanabe, K., Ouchi, M., Nakajima, K., et al. 2023, arXiv e-prints, arXiv:2305.02078, doi: 10.48550/arXiv.2305.02078
  * Watanabe et al. (in prep.) Watanabe et al., K. in prep.
  * Wiese et al. (1996) Wiese, W. L., Fuhr, J. R., & Deters, T. M. 1996, Atomic transition probabilities of carbon, nitrogen, and oxygen : a critical data compilation
  * Woosley & Janka (2005) Woosley, S., & Janka, T. 2005, Nature Physics, 1, 147, doi: 10.1038/nphys172
  * Yajima et al. (2022a) Yajima, H., Abe, M., Fukushima, H., et al. 2022a, arXiv e-prints, arXiv:2211.12970, doi: 10.48550/arXiv.2211.12970
  * Yajima et al. (2022b) Yajima, H., Abe, M., Khochfar, S., et al. 2022b, MNRAS, 509, 4037, doi: 10.1093/mnras/stab3092
  * Zamorani et al. (1981) Zamorani, G., Henry, J. P., Maccacaro, T., et al. 1981, ApJ, 245, 357, doi: 10.1086/158815