# RIOJA. Complex Dusty Starbursts in a Major Merger B14-65666 at z=7.15

Yuma Sugahara National Astronomical Observatory of Japan, 2-21-1 Osawa,
Mitaka, Tokyo 181-8588, Japan  Waseda Research Institute for Science and
Engineering, Faculty of Science and Engineering, Waseda University, 3-4-1
Okubo, Shinjuku, Tokyo 169-8555, Japan<EMAIL_ADDRESS>Javier Álvarez-
Márquez Centro de Astrobiología (CAB), CSIC-INTA, Ctra. de Ajalvir km 4,
Torrejón de Ardoz, E-28850, Madrid, Spain  Takuya Hashimoto Division of
Physics, Faculty of Pure and Applied Sciences, University of Tsukuba,Tsukuba,
Ibaraki 305-8571, Japan  Tomonaga Center for the History of the Universe
(TCHoU), Faculty of Pure and Applied Sciences, University of Tsukuba, Tsukuba,
Ibaraki 305-8571, Japan  Luis Colina Centro de Astrobiología (CAB), CSIC-INTA,
Ctra. de Ajalvir km 4, Torrejón de Ardoz, E-28850, Madrid, Spain  Akio K.
Inoue Waseda Research Institute for Science and Engineering, Faculty of
Science and Engineering, Waseda University, 3-4-1 Okubo, Shinjuku, Tokyo
169-8555, Japan  Department of Pure and Applied Physics, School of Advanced
Science and Engineering, Faculty of Science and Engineering, Waseda
University, 3-4-1 Okubo, Shinjuku, Tokyo 169-8555, Japan  Luca Costantin
Centro de Astrobiología (CAB), CSIC-INTA, Ctra. de Ajalvir km 4, Torrejón de
Ardoz, E-28850, Madrid, Spain  Yoshinobu Fudamoto National Astronomical
Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan  Waseda
Research Institute for Science and Engineering, Faculty of Science and
Engineering, Waseda University, 3-4-1 Okubo, Shinjuku, Tokyo 169-8555, Japan
Center for Frontier Science, Chiba University, 1-33 Yayoi-cho, Inage-ku, Chiba
263-8522, Japan  Ken Mawatari Division of Physics, Faculty of Pure and Applied
Sciences, University of Tsukuba,Tsukuba, Ibaraki 305-8571, Japan  Tomonaga
Center for the History of the Universe (TCHoU), Faculty of Pure and Applied
Sciences, University of Tsukuba, Tsukuba, Ibaraki 305-8571, Japan  Yi W. Ren
Department of Pure and Applied Physics, School of Advanced Science and
Engineering, Faculty of Science and Engineering, Waseda University, 3-4-1
Okubo, Shinjuku, Tokyo 169-8555, Japan  Santiago Arribas Centro de
Astrobiología (CAB), CSIC-INTA, Ctra. de Ajalvir km 4, Torrejón de Ardoz,
E-28850, Madrid, Spain  Tom J. L. C. Bakx Department of Space, Earth and
Environment, Chalmers University of Technology, Onsala Space Observatory,
SE-439 92 Onsala, Sweden  Carmen Blanco-Prieto Centro de Astrobiología (CAB),
CSIC-INTA, Ctra. de Ajalvir km 4, Torrejón de Ardoz, E-28850, Madrid, Spain
Daniel Ceverino Universidad Autonoma de Madrid, Ciudad Universitaria de
Cantoblanco, E-28049 Madrid, Spain  CIAFF, Facultad de Ciencias, Universidad
Autonoma de Madrid, E-28049 Madrid, Spain  Alejandro Crespo Gómez Centro de
Astrobiología (CAB), CSIC-INTA, Ctra. de Ajalvir km 4, Torrejón de Ardoz,
E-28850, Madrid, Spain  Masato Hagimoto Department of Physics, Graduate School
of Science, Nagoya University, Nagoya, Aichi 464-8602, Japan  Takeshi
Hashigaya  Department of Astronomy, Kyoto University Sakyo-ku, Kyoto 606-8502,
Japan  Rui Marques-Chaves Geneva Observatory, Department of Astronomy,
University of Geneva, Chemin Pegasi 51, CH-1290 Versoix, Switzerland  Hiroshi
Matsuo National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo
181-8588, Japan  The Graduate University for Advanced Studies (SOKENDAI),
2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan  Yurina Nakazato Department of
Physics, The University of Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo 113-0033, Japan
Miguel Pereira-Santaella Instituto de Física Fundamental (IFF), CSIC, Serrano
123, E-28006, Madrid, Spain  Yoichi Tamura Department of Physics, Graduate
School of Science, Nagoya University, Nagoya, Aichi 464-8602, Japan  Mitsutaka
Usui  Division of Physics, Faculty of Pure and Applied Sciences, University of
Tsukuba,Tsukuba, Ibaraki 305-8571, Japan  Naoki Yoshida Department of Physics,
The University of Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo 113-0033, Japan  Kavli
Institute for the Physics and Mathematics of the Universe (WPI), UT Institute
for Advanced Study, The University of Tokyo, Kashiwa, Chiba 277-8583, Japan
Research Center for the Early Universe, School of Science, The University of
Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo 113-0033, Japan

###### Abstract

We present JWST NIRCam imaging of B14-65666 (“Big Three Dragons”), a bright
Lyman-break galaxy system ($M_{\text{UV}}=-22.5$ mag) at $z=7.15$. The high
angular resolution of NIRCam reveals the complex morphology of two galaxy
components: galaxy E has a compact core (E-core), surrounded by diffuse,
extended, rest-frame optical emission, which is likely to be tidal tails; and
galaxy W has a clumpy and elongated morphology with a blue UV slope
($\beta_{\text{UV}}=-2.2\pm 0.1$). The flux excess, F356W$-$F444W, peaks at
the E-core ($1.05^{+0.08}_{-0.09}$ mag), tracing the presence of strong [O
iii]$\lambda\lambda$4960,5008 Å emission. ALMA archival data show that the
bluer galaxy W is brighter in dust continua than the redder galaxy E, while
the tails are bright in [O iii] 88 $\mathrm{\mu m}$. The UV/optical and sub-mm
SED fitting confirms that B14-65666 is a major merger in a starburst phase as
derived from the stellar mass ratio (3:1 to 2:1) and the star-formation rate,
$\simeq 1$ dex higher than the star-formation main sequence at the same
redshift. The galaxy E is a dusty ($A_{\text{V}}=1.2\pm 0.1$ mag) starburst
with a possible high dust temperature ($\geq 63\mbox{--}68$ K). The galaxy W
would have a low dust temperature ($\leq 27\mbox{--}33$ K) or patchy stellar-
and-dust geometry, as suggested from the infrared excess (IRX) and
$\beta_{\text{UV}}$ diagram. The high optical-to-FIR [O iii] line ratio of the
E-core shows its lower gas-phase metallicity ($\simeq 0.2$
$Z_{\text{$\odot$}}$) than the galaxy W. These results agree with a scenario
where major mergers disturb morphology and induce nuclear dusty starbursts
triggered by less-enriched inflows. B14-65666 shows a picture of complex
stellar buildup processes during major mergers in the epoch of reionization.

Galaxy evolution (594); High-redshift galaxies (734); Galaxy mergers (608);
Near infrared astronomy (1093);

††software: NumPy (Harris et al. 2020), SciPy (Virtanen et al. 2020), IPython
(Perez & Granger 2007), Matplotlib (Hunter 2007), Astropy (Astropy
Collaboration et al. 2013, 2018)

##  1 Introduction

Galaxy mergers are a fundamental process to build up the stellar content
through galaxy evolution. Gravitational perturbations by galaxy interactions
can strongly disturb the morpho-kinematics, which is characterized by tidal
tails and large velocity dispersion. It is theoretically predicted that such
perturbation promotes gas accretion from outskirts into the central parts of
galaxies that can trigger nuclear star formation (Mihos & Hernquist 1994,
1996; Barnes & Hernquist 1996). In observations of local galaxies, the star-
formation rate ($\mathrm{SFR}$) of interacting galaxies increases as projected
pair separation decreases to within several tens kpc (e.g., Barton et al.
2000). This $\mathrm{SFR}$ enhancement is particularly notable in collisions
of galaxies with similar stellar masses (i.e., “major mergers”; e.g., Ellison
et al. 2008). The gas accretion associated with mergers can also induce an
observed reduction in the gas-phase metallicity at a fixed mass or luminosity
(e.g., Kewley et al. 2006; Ellison et al. 2008). Even after the coalescence
phase, the major mergers are expected to cause energetic events in the
subsequent evolution of the merged galaxies, such as luminous infrared
starbursts, rapid black hole growths, and the appearance of quasars (Hopkins
et al. 2008).

At high redshifts, major mergers are predicted to play a more important role
in the evolutionary path because the merger rate theoretically increases with
increasing redshift by $\propto(1+z)^{2\mbox{--}3}$ (e.g., Fakhouri & Ma 2008;
Rodriguez-Gomez et al. 2015). This increase is consistent with observed major
merger fractions at $z\sim 0\mbox{--}5$ (e.g., Man et al. 2016; Ventou et al.
2017) when the redshift-dependent merger timescale is taken into account (Qu
et al. 2017; Snyder et al. 2017; Romano et al. 2021). Recently, rest-frame
ultraviolet (UV) images (Ribeiro et al. 2017; Duncan et al. 2019; Shibuya et
al. 2022) and [C ii] 158 $\mathrm{\mu m}$ data cubes (Romano et al. 2021) were
able to explore the merger fractions up to $z\sim 5\mbox{--}7$. Although the
measured fractions are yet to be converged and depend strongly on galaxy
masses, the high merger fraction, $20\mbox{--}80$% at $z\sim 5\mbox{--}7$,
supports a large number of galaxies that experience major mergers in the early
Universe. In terms of the stellar mass assembly, contributions of major
mergers are suggested to be comparable to in-situ star formation at $z>3$
(Duncan et al. 2019; Romano et al. 2021). While cosmological zoom-in
simulations also support the SFR enhancement by gas-rich mergers at $z>5$
(Ceverino et al. 2018), the effect of mergers on star formation appears to
reduce strongly from the local Universe to $z\sim 1\mbox{--}3$ (e.g., Silva et
al. 2018; Shah et al. 2022).

High angular resolution is key to pushing up the redshift frontier of merger
studies. Merging systems at $z\gtrsim 5$ are reported from observations of
Hubble Space Telescope (HST), James Webb Space Telescope (JWST, Gardner et al.
2023), and Atacama Large Millimeter/sub-millimeter Array (ALMA), which have
both high angular resolutions and highest sensitivities that can resolve high-
redshift galaxies. For example, dusty systems at $z\sim 4\mbox{--}7$
corresponding to hyper/ultra luminous infrared galaxies (ULIRGs) sometimes
show signatures of galaxy mergers (e.g. Oteo et al. 2016; Marrone et al. 2018;
Hygate et al. 2023). Triple mergers have also been observed, such as a dust-
free, luminous $\mathrm{Ly\alpha}$ emitter at $z=6.6$ (“Himiko”, Ouchi et al.
2013) and the [C ii]-bright galaxy at $z=4.6$ (Jones et al. 2020a). JWST has
recently updated the most distant major merger by spectroscopically confirming
a triply-lensed galaxy at $z=10.17$ that has clumps with different dust
attenuation (Hsiao et al. 2023a, b). For a statistical study, JWST has shown
that galaxy interactions significantly contribute to the bursty star-formation
and strong emission lines in low-mass galaxies at $z\sim 5\mbox{--}6$ (Asada
et al. 2024). However, few merging galaxies are observed and detected with
multiple bands in both UV and far-infrared (FIR), limiting our ability to
interpret the complete properties of distant merging systems. Instead, multi-
wavelength, high angular-resolution observations are essential to examine the
complex properties of stars, gas, and dust through UV-to-FIR spectral energy
distributions (SEDs) of resolved galaxy components.

This paper presents new observations of a bright merging galaxy system at
$z=7.15$, B14-65666, also known as the Big Three Dragons (Hashimoto et al.
2019), with the Near Infrared Camera (NIRCam; Rieke et al. 2023) on board
JWST. This galaxy was first identified as a Lyman-break galaxy from the
UltraVISTA survey (Bowler et al. 2012, 2014). Its UV absolute magnitude of
$M_{\text{UV}}=-22.5$ mag is four times brighter than the characteristic value
at $z=7$ ($M^{*}_{\text{UV}}\simeq-21.0$ mag; e.g., Finkelstein et al. 2015;
Bouwens et al. 2021). Multi-wavelength follow-up observations were conducted
to detect the $\mathrm{Ly\alpha}$ emission line (Furusawa et al. 2016), the UV
continuum with high-spatial resolution (Bowler et al. 2017), the [O iii] 88
$\mathrm{\mu m}$ and [C ii] 158 $\mathrm{\mu m}$ emission lines (Hashimoto et
al. 2019), and dust continua at 88, 122, and 158 $\mathrm{\mu m}$ (Bowler et
al. 2018; Hashimoto et al. 2019; Sugahara et al. 2021). From the dust
continua, the infrared luminosity is estimated to be
$\log{L_{\text{IR}}/L_{\sun}}=11.6^{+0.6}_{-0.3}$ at the dust emissivity of
$\beta_{\text{d}}=2.0$ (Sugahara et al. 2021). The galaxy was also observed by
ALMA Bands 3 and 7, but it remained undetected in CO(6-5), CO(7-6), [C i](2-1)
(Hashimoto et al. 2023a) and [N ii] 122 $\mathrm{\mu m}$ (Sugahara et al.
2021). The systemic redshift was determined from the [O iii] and [C ii]
emission lines to be $z=7.1520\pm 0.0003$ (Hashimoto et al. 2019). These
observations do not provide any evidences towards the presence of an active
galactic nucleus (AGN) in the B14-65666 system. These ancillary multi-
wavelength data, particularly its rich ALMA data, make B14-65666 a valuable
follow-up target. We use our new NIRCam data and the ancillary data to
investigate detailed merger properties during the epoch of reionization (EoR).

This paper is organized as follows. Section 2 describes NIRCam and ancillary
data, and reduction processes. Section 3 explains analysis tools including
image analyses, SED fitting, and photoionization models. Section 4 presents
results about morphology, comparison with ALMA data, and SED fitting. Section
5 discusses relations between UV and FIR photometry, a metallicity
distribution, and our interpretations about B14-65666. Finally, Section 6
summarizes our conclusions. Throughout this paper, we use the AB magnitude
system (Oke & Gunn 1983) and the flat $\Lambda$CDM cosmology
($\Omega_{\mathrm{M}}=0.310$ and $h=H_{0}/[100\ \mathrm{km\ s^{-1}\
Mpc^{-1}}]=0.677$; Planck Collaboration et al. 2020). At $z=7.15$, $1$ arcsec
corresponds to $5.27$ kpc.

##  2 Observations and Data Reductions

The NIRCam observations of B14-65666 were conducted as part of the RIOJA
project: “the Reionization and the ISM/Stellar Origins with JWST and ALMA”
(JWST GO1 PID1840; PIs: J. Álvarez-Márquez and T. Hashimoto; Alvarez-Marquez
et al. 2021; Hashimoto et al. 2023b). NIRCam images were taken on Dec. 20,
2022 with a total exposure time of $644$ s in F200W and F277W, $730$ s in
F115W and F356W, and $1675$ s in F150W and F444W. The F115W, F200W, F277W, and
F356W observations adopted the BRIGHT1 readout pattern with eight or nine
groups and one integration; the F150W and F444W observations adopted the
SHALLOW4 with eight groups one integration. All the exposures were four times
dithered using the INTRAMODULEBOX pattern. In the RIOJA project, the NIRCam
filters were selected such as to trace the continuum emission over the entire
$1$–$5$ $\mathrm{\mu m}$ wavelength range with mostly avoiding the
contamination by the main optical emission lines. Given the spectroscopic
redshift of $z=7.1520$, the short wavelength channels of F115W, F150W, and
F200W correspond to the rest-frame ultraviolet ($1400\mbox{--}2500$ Å). The
F277W and F356W bands cover the stellar continua on both sides of the Balmer
break without any contamination of strong optical emission lines. The F444W
band includes the $\mathrm{H\beta}$ and [O iii]$\lambda\lambda$4960,5008 Å
emission lines as well as underlying stellar continuum. The corresponding
rest-frame wavelengths and the angular resolution of the NIRCam filters are
listed in Table 1.

The data were reduced using our custom pipeline based on the JWST calibration
pipeline version 1.9.4 (Bushouse et al. 2023) under CRDS context
jwst_1041.pmap. Our additional reduction steps are: (1) snowballs and wisps
removal as described in Bagley et al. (2023) and (2) background homogenization
(including $1/f$-noise subtraction) applied in Pérez-González et al. (2023).
The pixel scales of the final products are set to $0.03$ arcsec$/$pixel.

The ALMA data used in this work were presented in previous studies. ALMA Bands
6 and 8 data were taken in Cycles 4 (ID 2016.1.00954.S, PI: A. K. Inoue) and 5
(ID 2017.1.00190.S, PI: A. K. Inoue) to target the [C ii] 158 $\mathrm{\mu m}$
and [O iii] 88 $\mathrm{\mu m}$ emission lines (Hashimoto et al. 2019). ALMA
Band 7 data were taken in Cycle 7 (ID: 2019.1.01491.S, PI: A. K. Inoue) to
target the [N ii] 122 $\mathrm{\mu m}$ emission line (Sugahara et al. 2021).
In all the observations, one of the four spectral windows was placed to
observe the targeted lines while the others were for the underlying dust
continua. The raw data were reduced with standard scripts on Common Astronomy
Software Applications (CASA Team et al. 2022). [O iii] and [C ii] emission-
line maps (i.e., moment-0 maps) were created by integrating wavelength ranges
of $\simeq 600$ $\mathrm{km\ s^{-1}}$ around the line centers after the
continuum subtraction (Hashimoto et al. 2019). Dust continuum maps were
created using channels that are unaffected by the emission lines. The
synthesized beam FWHM is $0\farcs 29\times 0\farcs 23$, $1\farcs 15\times
0\farcs 97$, and $0\farcs 39\times 0\farcs 37$ in Bands 6, 7, and 8
observations with the natural weighting. More information on the observations
and data reduction is described in the references above. We did not use ALMA
Band 3 data (ID: 2018.1.01673.S, PI: T. Hashimoto; Hashimoto et al. 2023a) in
this study because neither continuum nor lines were detected.

Astrometry of the NIRCam images were corrected with stars in the field of view
(FoV). In the NIRCam FoV, there exist three stars listed on the Gaia DR3 main
catalog (Gaia Collaboration et al. 2016, 2023). We aligned the NIRCam images
to the three Gaia stars, and then we corrected small offsets between the
NIRCam images using faint stars. The astrometric uncertainties of the NIRCam
images are $<0\farcs 015$. The peak positional accuracy of ALMA data, which
can be estimated from the beam size and peak $\text{S}/\text{N}$111Equation
(10.7) in ALMA Cycle 10 Technical Handbook:
https://almascience.nao.ac.jp/documents-and-tools/cycle10/alma-technical-
handbook, is $\sim 0\farcs 03$ for [C ii] 158 $\mathrm{\mu m}$ and [O iii] 88
$\mathrm{\mu m}$ line maps and $\sim 0\farcs 06$ for dust continuum maps at
$158$ and $88$ $\mathrm{\mu m}$. As these estimates for the ALMA data are
based on an equation for compact sources, the accuracy may be worse for
extended sources like B14-65666.

##  3 Analyses

###  3.1 Image analyses

When comparing different bands of the NIRCam images pixel by pixel, we matched
the point spread functions (PSFs) of the images to that of the F444W image by
kernel convolution. We constructed effective PSFs following Anderson & King
(2000) using photutils (Bradley et al. 2023). Then, we homogenized the PSFs of
the NIRCam images using kernels created by PyPHER (Boucaud et al. 2016).
Finally, the image pixels were resampled using reproject to match the F444W
image pixels. When comparing the NIRCam and ALMA [O iii] images, the PSFs of
the NIRCam images are homogenized to the synthesized clean beam (i.e., two-
dimensional Gaussian) of the ALMA images following the same procedure. The
ALMA image pixels are resampled to match the F444W image pixels.

Photometry was performed with fixed apertures on the PSF-homogenized images
and aperture correction was not applied. B14-65666 consists of several
components and structures (see Section 4 and Figure 1). Therefore, we chose a
circular aperture of $0\farcs 9$ in diameter to measure the total photometry
of B14-65666 and smaller apertures to measure the photometry of individual
components (see Figure 2). The local sky background computed in an annulus
from $0\farcs 95$ to $1\farcs 95$ were subtracted before the photometry
measurements. The photometric uncertainties were estimated from random
aperture photometry. We scattered the apertures over the images around
B14-65666 where bright objects are masked and we computed the uncertainties as
the standard deviation of the aperture-photometry distributions. In addition,
the absolute flux uncertainty was included in the photometric uncertainties as
a secure value of $5$% (Rigby et al. 2023)222See also “Data Calibration
Considerations” in JWST User Documentation: https://jwst-docs.stsci.edu.

###  3.2 SED fitting

We performed SED fitting to the NIRCam photometry using Bayesian Analysis of
Galaxies for Physical Inference and Parameter EStimation (Bagpipes, Carnall et
al. 2018). We tested SED fitting with and without ALMA fluxes to see changes
in obtained best-fit parameters. The ALMA [O iii] 88 $\mathrm{\mu m}$ line and
continuum fluxes were taken from Hashimoto et al. (2019). We note that these
ALMA fluxes were measured with spatially two-component Gaussian fittings,
which differs from the aperture photometry for the NIRCam images. The redshift
is fixed at $z=7.1520$ in the fitting.

The fitting parameters are the star-formation age $t_{\text{SF}}$, the stellar
mass $M_{*}$, the ionization parameter $U$, the V-band dust attenuation
$A_{\text{V}}$, the stellar metallicity $Z_{*}$, and the mass of an old
stellar component $M_{\text{old}}$. The star-formation history of B14-65666 is
simply assumed to be constant for $t_{\text{SF}}$ down to $z=7.1520$. We also
tested the delayed-$\tau$ and exponential star-formation histories, but the
fitting results were insensitive to different star-formation histories,
because B14-65666 is dominated by a young burst of star formation ($\sim 10$
Myr). The gas-phase metallicity is assumed to be identical to the stellar
metallicity. The ionization parameter $U$ controls nebular emission-line
ratios, which is implemented using a photoionization code Cloudy (Ferland et
al. 1998, 2017). The hydrogen density in the Cloudy models is fixed to be
$n_{\text{H}}=100$ $\mathrm{cm^{-3}}$ in Bagpipes. The dust attenuation curve
follows the Calzetti et al. (2000) law. In addition, considering that
B14-65666 is a major merger, stellar components older than recent starbursts
should exist. To constrain the mass of the old stellar components, we added
another constant star formation lasting for $500$ Myr, from $600$ Myr ago to
$100$ Myr ago with respect to $z=7.15$, which does not strongly contribute to
UV continuum and emission-line fluxes. This selected star-formation timescale
is arbitrary, but changes of this timescale do not affect our conclusion. The
stellar metallicity of the old components is fixed to $Z_{\text{star}}=0.5$
$Z_{\text{$\odot$}}$ (assuming that $Z=0.02$ for solar metallicity; Asplund et
al. 2021) to reduce the number of free parameters, but we checked that results
do not vary significantly with the assumed $Z_{\text{star}}$ value. The best-
fit parameters are computed as the median of the posterior distributions. The
$1\sigma$ fitting uncertainties are evaluated from the $16$ and $84$
percentiles of the posterior distributions.

For FIR SED, we added fitting parameters of the minimum stellar intensity
$\mathcal{U}_{\mathrm{min}}$ and the fraction $\gamma$ of the dust mass heated
by $\mathcal{U}>\mathcal{U}_{\mathrm{min}}$. Bagpipes has adopted the Draine &
Li (2007) dust emission models. Instead of the originally adopted ones, we
used the updated version (Draine et al. 2014) implemented in the CIGALE SED
fitting code (Boquien et al. 2019), which allows up to
$\mathcal{U}_{\mathrm{min}}=50$. The model parameters other than
$\mathcal{U}_{\mathrm{min}}$ and $\gamma$ were fixed: the power-law index of
the stellar intensity $\alpha_{\mathcal{U}}=2$, the maximum stellar intensity
$\mathcal{U}_{\text{max}}=10^{7}$, and the PAH mass fraction
$q_{\text{PAH}}=2.5$%. The former two fixed parameters are appropriate to
reproduce high-redshift galaxies (Burgarella et al. 2022) as well as nearby
galaxies (Draine et al. 2014) and the latter $q_{\text{PAH}}$ was fixed
because our observations are insensitive to mid-IR wavelengths. We note that
while we added two fitting parameters for the FIR SED, dust continua are
constrained at only $2\mbox{--}3$ bands. Thus, adding ALMA dust continua in
the SED fitting is not quite effective to constrain the stellar properties
better (as shown in Section 4.3), but it helps estimate the dust temperature
necessary for the energy balance between dust attenuation and emission.

For input stellar models, the Binary Population and Spectral Synthesis code
(BPASS, Eldridge et al. 2017) version 2.3 (Byrne et al. 2022) was used. BPASS
includes effects of binary populations, which are important to reproduce
strong [O iii] emission lines observed at high redshift (Sugahara et al.
2022). Specifically, the version 2.3 implements high $\alpha/\text{Fe}$ cases,
reflecting abundance patterns found at $z\gtrsim 2$ dominated by core-collapse
supernovae (Steidel et al. 2014, 2016). This enhancement of $\alpha$ elements
is also key to reproducing [O iii] emission line strengths and hard UV
spectral slopes at high redshift (e.g., Shapley et al. 2019; Sugahara et al.
2022). We used the case of $\varDelta\log(\alpha/\text{Fe})=+0.6$, where the
oxygen abundance is almost the solar value $12+\log\text{O}/\text{H}=8.69$
(Asplund et al. 2021) at $Z=0.02$ (Byrne et al. 2022). The adopted initial
mass function (IMF) is a Kroupa (2001) IMF in the mass range from $0.1$ to
$300$ $\mathrm{M_{\sun}}$. Their choice of the upper mass limit leads to a 45%
higher ionizing flux and stronger hydrogen recombination lines than a choice
of an upper limit of $100$ $\mathrm{M_{\sun}}$ (Stanway et al. 2016). The
nebular-continuum and emission-line models are generated based on these
stellar models.

Figure 1:  Left: Pseudo-RGB composite image of B14-65666, where F356W (red),
F277W (green), and F150W (blue) bands reflect continuum light. The RGB images
are convolved to match the spatial resolution of the F444W-band image.
B14-65666 shows complex morphology including the elongated galaxy W and the
galaxy E, which has a compact core being surrounded by the red tails. Right:
NIRCam band stamps of B14-65666, which show the same sky area as the left
panel. They are illustrated in the native PSFs, which are shown with the white
circles at the bottom lefts.

###  3.3 Photoionization models with Cloudy

In our emission-line modeling, we used a photoionization code Cloudy version
17.02 (Ferland et al. 1998, 2017). We note that our setup is different from
the one used in the Bagpipes SED fitting. We assumed a plane-parallel geometry
under the constant pressure and elemental and dust-grain abundances in the H
ii regions that are stored in Cloudy. For the helium abundance, we simply
added stellar yields to Big Bang nucleosynthesis as a function of the gas-
phase metallicity (Groves et al. 2004). The input ionizing spectra were
constructed from the BPASS version 2.3 and the stellar metallicity is assumed
to be as identical to the gas-phase metallicity. The star-formation history is
set to be a $10$-Myr constant star formation. Adopted variable parameters are
the ionization parameter $U$ and hydrogen density $n_{\text{H}}$ at the
illuminated surface, and the gas-phase metallicity $Z_{\text{gas}}$. The
calculations stop at the edge of the H ii region where the electron fraction
is $10^{-2}$. The output emission-line fluxes are normalized by the
$\mathrm{H\beta}$ line fluxes.

Table 1: PFS sizes of the NIRCam images.

Band |  $\lambda_{\text{rest}}$ |  PSF FWHM
---|---|---
| [$\AA$] | [arcsec]
F115W |  $1416$ Å | $0.040$
F150W |  $1841$ Å | $0.050$
F200W |  $2441$ Å | $0.065$
F277W |  $3418$ Å | $0.092$
F356W |  $4371$ Å | $0.116$
F444W |  $5424$ Å | $0.145$

Note. — The rest-frame wavelengths, $\lambda_{\text{rest}}$, are converted
from the pivot wavelengths. The empirical PSF FWHM is taken from “NIRCam Point
Spread Functions” in JWST User Documentation.

Figure 2: Left: Apertures used for photometry of different components in
B14-65666, on the F444W image. The total aperture (white) contains flux from
the entire system, and the galaxies E and W are represented with the red and
blue circles, respectively. The cores of the galaxies are depicted with the
dashed circles. The thin white circles show the tails (N-tail and S-tail). The
background contours are at the $2^{n}\sigma$ significance levels
($n=1,2,3,\ldots$). Middle: [O iii]$\lambda\lambda$4960,5008 flux map
estimated from the F356W and F444W images (see text for details). The flux map
is drawn in pixels with $\text{S}/\text{N}>2$ in the F444W image. The gray
contours are the F444W flux map, which is the same as in the left panel.
Right: Rest-frame [O iii] equivalent width map, which is re-binned to improve
$\text{S}/\text{N}$ per pixel. The equivalent widths are estimated by dividing
the flux map by the stellar continuum inferred from F356W. The [O iii] fluxes
reach the peaks at the cores of the galaxies E and W. The rest-equivalent
width values are $\text{EW}_{0}([\mbox{O\,{\sc iii}}])=2050\pm 290$ and
$740\pm 170$ Å for the E-core and W-core, respectively, which implies the
presence of the nuclear starbursts induced by the major merger. The equivalent
widths in the tails are not significant in the $3\sigma$ level (Table 3).
Table 2: NIRCam photometry and ALMA measurements of B14-65666 components.

Aperture | F115W | F150W | F200W | F277W | F356W | F444W | Band8† | Band7†† | Band6† | [OIII]88†
---|---|---|---|---|---|---|---|---|---|---
| [$\mathrm{\mu Jy}$] | [$\mathrm{\mu Jy}$] | [$\mathrm{\mu Jy}$] | [$\mathrm{\mu Jy}$] | [$\mathrm{\mu Jy}$] | [$\mathrm{\mu Jy}$] | [$\mathrm{\mu Jy}$] | [$\mathrm{\mu Jy}$] | [$\mathrm{\mu Jy}$] | [$\mathrm{Jy~{}km~{}s^{-1}}$]
Total | $0.374\pm 0.100$ | $0.447\pm 0.043$ | $0.429\pm 0.112$ | $0.661\pm 0.047$ | $0.705\pm 0.046$ | $1.346\pm 0.070$ | $470\pm 128$ | $218\pm 19$ | $130\pm 25$ | $1.50\pm 0.18$
E | $0.159\pm 0.060$ | $0.162\pm 0.023$ | $0.261\pm 0.051$ | $0.343\pm 0.027$ | $0.355\pm 0.024$ | $0.773\pm 0.040$ | $208\pm 83$ | - | $41\pm 23$ | $0.92\pm 0.14$
W | $0.225\pm 0.045$ | $0.298\pm 0.019$ | $0.187\pm 0.040$ | $0.298\pm 0.022$ | $0.305\pm 0.020$ | $0.496\pm 0.026$ | $246\pm 73$ | - | $87\pm 26$ | $0.57\pm 0.09$
E-core | $0.107\pm 0.022$ | $0.128\pm 0.008$ | $0.163\pm 0.018$ | $0.180\pm 0.011$ | $0.167\pm 0.010$ | $0.440\pm 0.022$ | - | - | - | -
W-core | $0.203\pm 0.023$ | $0.232\pm 0.013$ | $0.186\pm 0.019$ | $0.215\pm 0.013$ | $0.220\pm 0.012$ | $0.349\pm 0.018$ | - | - | - | -
N-tail | $0.042\pm 0.025$ | $0.016\pm 0.007$ | $0.042\pm 0.020$ | $0.057\pm 0.009$ | $0.080\pm 0.009$ | $0.116\pm 0.008$ | - | - | - | -
S-tail | $-0.005\pm 0.038$ | $0.014\pm 0.011$ | $0.050\pm 0.033$ | $0.091\pm 0.014$ | $0.091\pm 0.011$ | $0.161\pm 0.011$ | - | - | - | -

References. — †Hashimoto et al. (2019), ††Sugahara et al. (2021)

##  4 Results

###  4.1 Morphology, Colors, and Optical [O iii] Line Map

B14-65666 consists of two bright components: galaxies E (east) and W (west),
which were referred to as the clumps A and B in Hashimoto et al. (2019),
respectively. Hashimoto et al. (2019) reported that these two galaxies have
similar UV, IR, [O iii] 88 $\mathrm{\mu m}$, and [C ii] 158 $\mathrm{\mu m}$
luminosities within a factor of two and their line-of-sight velocity
separation is $\simeq 200$ $\mathrm{km\ s^{-1}}$. For these reasons, B14-65666
is thought to be a major merger at $z=7.1520$. Figure 1 shows the NIRCam-band
color image and the six band stamps of the two galaxies of B14-65666.

The NIRCam imaging clearly illustrates that the galaxies E and W have
different morphology. The galaxy E has a bright compact core. This compact
core is unresolved even in the F115W image, which has the highest angular
resolution in the images used in this work. Under an assumption that the core
of the galaxy E has an exponential surface brightness profile, the deconvolved
effective radius in F115W is $r_{\text{e}}<0.016$ arcsec ($85$ pc) at the
$2\sigma$ significance limit. These core effective radii are smaller than the
average effective radii ($0.7\mbox{--}1$ kpc) of observed and simulated
galaxies at $M_{*}\sim 10^{9}$ $\mathrm{M_{\sun}}$ at $z=7$ (Ormerod et al.
2024; Costantin et al. 2023).

In contrast, the galaxy W is elongated from southeast to northwest. This
feature is seen in all the NIRCam images and its length is $\simeq 0.3$ arcsec
($1.5$ kpc). Such elongated morphology look like tidal tails and it is often
observed under tidal effect like major mergers in the local Universe (e.g.,
Borne et al. 2000; García-Marín et al. 2009). In addition, the F115W and F150W
images suggest a clumpy internal structure of the galaxy W, which may consist
of three clumps. These compactness and clumpiness of B14-65666 agree with
rest-frame UV clumps in bright high-redshift galaxies revealed by recent JWST
observations (e.g., Treu et al. 2023; Trussler et al. 2023; Marques-Chaves et
al. 2024) and numerical simulations (e.g., Mandelker et al. 2017; Ceverino et
al. 2023; Nakazato et al. 2024).

In addition to the compact and clumpy structure, the NIRCam images have
identified some extended emission around the galaxy E. This emission is
clearly seen in the F444W image and consists of two diffuse sub-components
north and south-east of the galaxy E (Figure 1). Actually, the spatially
extended components were already suggested in the HST F140W image (Bowler et
al. 2017) and the [O iii] 88 $\mathrm{\mu m}$ line map (Hashimoto et al.
2019), but they were not explicitly reported in the literature due to low
$\text{S}/\text{N}$. In the NIRCam imaging, the diffuse components are
confirmed at $\text{S}/\text{N}>6$ in the F277W, F356W, and F444W bands and
they are also suggested at lower significance levels ($\text{S}/\text{N}\simeq
1\mbox{--}2$) in the F150W and F200W images. The F444W bandpass includes
$\mathrm{H\beta}$ and [O iii]$\lambda\lambda$4960,5008 emission lines. The
F277W bandpass samples shorter wavelengths than the Balmer break, at which the
stellar continuum by old stars becomes faint. Therefore, these multiple band
detections imply that the diffuse sub-components are bright in young stellar
continuum or nebular emission (e.g., the [O iii] and [O
ii]$\lambda\lambda$3727,3730 Å lines and nebular continuum). Given that
B14-65666 is a major merging system, we assumed that they are likely to be
tidal tails stripped from the galaxies.

Figure 3:  ALMA line and dust contours overlaid on the background NIRCam
image. The NIRCam image is a stack of F115W to F356W bands to trace the
continuum light, which is homogenized to the F444W PSF. The white contours are
drawn with $2\sigma$ intervals from the $2\sigma$ level, where $\sigma=0.05$
and $0.02$ $\mathrm{Jy\ km\ s^{-1}}$ for the [O iii] and [C ii] line maps and
$\sigma=30,20$, and $9$ $\mathrm{\mu Jy\ beam^{-1}}$ for the dust maps,
respectively. The gray hatched ellipses at the left bottom corners are the
synthesized beam sizes of ALMA observations.  Table 3: Observed and Physical
Properties of Components.

Components | flux([O iii]) | EW${}_{0}([\mbox{O\,{\sc iii}}])$ | SFR0 | sSFR0 | SFR${}_{\text{10Myr}}$ | $\beta_{\text{UV}}$
---|---|---|---|---|---|---
| [$\mathrm{10^{-17}\,erg\,s^{-1}\,cm^{-2}}$] | [$\mathrm{\mathring{A}}$] | [$\mathrm{M_{\odot}\,yr^{-1}}$] | [$\mathrm{Gyr^{-1}}$] | [$\mathrm{M_{\odot}\,yr^{-1}}$] |
Total | $10.1\pm 1.3$ | $1140\pm 220$ | $225^{+71}_{-56}$ | $52^{+56}_{-28}$ | $207^{+65}_{-51}$ | $-1.57^{+0.11}_{-0.13}$
E | $6.6\pm 0.7$ | $1480\pm 260$ | $178^{+52}_{-39}$ | $53^{+26}_{-20}$ | $178^{+50}_{-39}$ | $-1.18^{+0.15}_{-0.16}$
W | $2.6\pm 0.4$ | $750\pm 170$ | $37^{+12}_{-8}$ | $23^{+15}_{-11}$ | $37^{+12}_{-8}$ | $-2.19^{+0.11}_{-0.09}$
E-core | $4.3\pm 0.4$ | $2050\pm 290$ | $180^{+113}_{-58}$ | $506^{+253}_{-182}$ | $34^{+13}_{-9}$ | $-1.78^{+0.11}_{-0.13}$
W-core | $2.0\pm 0.3$ | $740\pm 170$ | $29^{+13}_{-7}$ | $27^{+28}_{-12}$ | $29^{+14}_{-6}$ | $-2.24^{+0.13}_{-0.10}$
N-tail | $0.6\pm 0.2$ | $580\pm 260$ | - | - | - | -
S-tail | $1.1\pm 0.2$ | $960\pm 340$ | - | - | - | -

Note. — SFR0 is the instant star-formation rate
($\mathrm{SFR}_{0}=M_{*}/t_{\text{SF}}$) and SFR${}_{\text{10Myr}}$ is the
star-formation rate averaged in 10 Myr, that is,
$\mathrm{SFR}_{\text{10Myr}}=M_{*}(<10\ \text{Myr})/10\ \text{Myr}$.

Table 3.3 lists measured photometry of the total and the galaxies E and W. The
table also lists photometry of the tails (N-tail and S-tail) and the cores of
the galaxies E and W (E-core and W-core). The applied apertures are shown in
the left panel of Figure 2. Our measurements of the total photometry are
consistent with the values on the COSMOS2020 catalog (J, H, and Ks bands in
UltraVISTA and Ch1 and 2 in Spitzer; Weaver et al. 2022). The F115W magnitudes
of the galaxies E and W are measured to be $25.9^{+0.5}_{-0.3}$ and
$25.5^{+0.2}_{-0.2}$ mag, respectively, corresponding to
$M_{\text{UV}}=-21.1^{+0.5}_{-0.3}$ and $-21.5^{+0.2}_{-0.2}$ mag. These
magnitudes are similar to and brighter than the characteristic magnitude,
$M^{*}_{\text{UV}}\simeq-21.0$ mag, of the UV luminosity function at $z\sim 7$
(e.g., Finkelstein et al. 2015; Bouwens et al. 2021). From UV to optical
wavelengths, the galaxy E has a redder color than the galaxy W. The tails
largely contribute this red color of the galaxy E, as shown in Figure 1. This
color difference suggests that the galaxies E and W have different stellar and
dust properties.

A characteristic color index is the F444W excess, corresponding to the [O
iii]$\lambda\lambda$4960,5008 emission line strength. The F444W band probes
the [O iii] and $\mathrm{H\beta}$ lines, while the F356W band probes the
$\mathrm{H\beta}$ line. The filter transmission of the $\mathrm{H\beta}$ line
is similar in both of the bands, 37% (F356W) and 50% (F444W). Therefore, by
assuming a constant continuum flux density between F356W and F444W, we created
a optical [O iii] line map from the F356W$-$F444W color with little
contamination of $\mathrm{H\beta}$ emission. Other emission lines of
$\mathrm{H\gamma}$, $\mathrm{H\delta}$, and [O iii] 4363 Å may lead to
underestimation of the [O iii] flux by $8\mbox{--}15$%, according to Cloudy
calculations at $0.1<Z_{\text{gas}}/Z_{\text{$\odot$}}<0.4$ and
$-1.5<\log{U}<-2.5$.

The middle panel of Figure 2 shows the optical [O iii] line map, which
indicates strong [O iii] emission at the cores. The E-core and W-core exhibit
integrated [O iii] fluxes of $(4.3\pm 0.4)\times 10^{-17}$ and $(2.0\pm
0.3)\times 10^{-17}$ $\mathrm{erg\ s^{-1}\ cm^{-2}}$, respectively. Being
regarded as a proxy of star formation, this strong [O iii] emission indicates
nuclear starbursts likely induced by the galaxy interaction, as seen in major
mergers in the local Universe (e.g., Ellison et al. 2008). The right panel of
Figure 2 shows a map of the rest-frame [O iii] equivalent width
($\text{EW}_{0}$). The E-core and W-core exhibit $\text{EW}_{0}([\mbox{O\,{\sc
iii}}])\sim 2050\pm 290$ and $740\pm 170$ Å, respectively, which is comparable
to [O iii] emitters at $z=5\mbox{--}7$ (Matthee et al. 2023; Boyett et al.
2024). In the tails, the [O iii] fluxes are detected at the
$3\mbox{--}4\sigma$ significance levels, although this value may be
overestimated because their red color may reflect their red optical continuum
slopes, which is different from our assumption. The obtained [O iii] fluxes
and equivalent widths are listed in Table 3.

Figure 4: SEDs of B14-65666 and the results of the SED fitting. (a) The SED
fitting results for the NIRCam and ALMA measurements. Each color shows the
data of each component: total (black) and the galaxies E (red) and W (blue).
The circles indicate the measured photometry and the error bars indicate the
measurement errors including the absolute flux uncertainties. The crosses show
the obtained best-fit photometry of the SED fitting. The solid lines and
shades show the best-fit curves and uncertainties taking $16$ to $84$
percentiles of the posterior distributions. The bottom gray shades of the left
panel show the transmission curves of the NIRCam filters. (b) The SED fitting
results only for the NIRCam photometry, without ALMA measurements. The best-
fit parameters are comparable to the ones obtained for the NIRCam and ALMA
measurements (Table 4). (c) The SED of the E-core (red) and W-core (blue) and
the obtained best-fit results for them. The dotted line is shown for
comparison, which is the same as the black line in the panel (b).

###  4.2 NIRCam and ALMA imaging: Stellar versus dust distribution

The combination of the new NIRCam images with archival high angular-resolution
($0\farcs 5$) ALMA data provides us an unique opportunity to investigate the
detailed stellar, ionized, and dust structure in an interacting system at
$z>7$. Figure 3 compares the NIRCam image, and ALMA line and dust maps. The
NIRCam image is a composite of all the bands, except for F444W, to exhibit the
stellar continuum. All the ALMA detections overlap the positions of the
galaxies E and W, showing no spatial offsets between the UV-to-optical and FIR
dust and line emission. The [C ii] 158 $\mathrm{\mu m}$ and [O iii] 88
$\mathrm{\mu m}$ emission lines in the galaxy E are $1\mbox{--}1.5$ times as
bright as those in the galaxy W (Hashimoto et al. 2019). This trend is
qualitatively the same as the one seen in our [O iii]$\lambda\lambda$4960,5008
line map (Table 3). In contrast to the emission lines, the dust emission at
$88$ and $158$ $\mathrm{\mu m}$ is brighter in the galaxy W than in the galaxy
E (Hashimoto et al. 2019). This dust spatial distribution might indicate that
the galaxy W experiences dustier star-formation than the galaxy E. As denoted
in Section 4.1, however, the galaxy W presents a bluer UV color than the
galaxy E, which is inconsistent with a previous finding that dust emission
tend to be accompanied with red components in $z>6$ galaxies (Bowler et al.
2022). We will come back to this point in Section 5.1.

The tails are detected only in the [O iii] 88 $\mathrm{\mu m}$ line map,
whereas they appear neither in the [C ii] 158 $\mathrm{\mu m}$ line map nor in
the dust maps. NIRCam detects the tails in the stacked continuum image,
generated using the bands from F115W to F356W, and the [O
iii]$\lambda\lambda$4960,5008 line map. However, the tails are faint in the
rest-frame UV filters (e.g., F150W), indicating red continuum slope from
optical to UV. The presence of continuum and line emission implies that the
tails are actively forming stars, but the red continuum slope and the no
emission of dust continua suggest a complex nature of the tails.

###  4.3 SED fitting of total and resolved components

We performed SED fitting combining the NIRCam and ALMA datasets for the entire
system and the resolved galaxy components of B14-65666. The apertures for each
component is illustrated in the left panel of Figure 2 and the NIRCam and ALMA
fluxes used in the SED fitting are summarized in Table 3.3. The best-fit SEDs
and measured physical parameters are presented in Figure 4 and Table 4,
respectively.

Table 4: Results of the SED fitting.

Components | Age | Formed Stellar Mass | Av | $\log U$ | $Z_{*}$ | Old Stellar Mass | $U_{\mathrm{min}}$ | $\gamma$
---|---|---|---|---|---|---|---|---
| [$\mathrm{Myr}$] | [$\mathrm{\log M_{\odot}}$] | [$\mathrm{mag}$$\mathrm{\left(\mathrm{AB}\right)}$] |  | [$\mathrm{Z_{\odot}}$] | [$\mathrm{\log M_{\odot}}$] |  |
NIRCam + ALMA data
Total | $35^{+20}_{-12}$ | $9.78^{+0.16}_{-0.18}$ | $0.78^{+0.07}_{-0.06}$ | $-2.09^{+0.58}_{-0.34}$ | $0.21^{+0.05}_{-0.04}$ | $<9.59$ | $>10.45$ | $0.88^{+0.07}_{-0.07}$
E | $<93$ | $9.55^{+0.19}_{-0.20}$ | $0.98^{+0.08}_{-0.07}$ | $-2.06^{+0.61}_{-0.34}$ | $0.19^{+0.04}_{-0.04}$ | $<9.55$ | $>16.33$ | $>0.83$
W | $22^{+13}_{-8}$ | $9.09^{+0.16}_{-0.18}$ | $0.46^{+0.08}_{-0.06}$ | $-1.95^{+0.54}_{-0.39}$ | $0.32^{+0.06}_{-0.05}$ | $<8.95$ | $>7.61$ | $<0.80$
Only NIRCam data
Total | $18^{+21}_{-9}$ | $9.61^{+0.25}_{-0.27}$ | $0.86^{+0.11}_{-0.11}$ | $-2.14^{+0.47}_{-0.34}$ | $0.26^{+0.09}_{-0.08}$ | $<9.68$ | - | -
E | $<53$ | $9.51^{+0.17}_{-0.14}$ | $1.21^{+0.13}_{-0.14}$ | $-2.19^{+0.42}_{-0.30}$ | $0.28^{+0.09}_{-0.08}$ | $<9.28$ | - | -
W | $42^{+39}_{-17}$ | $9.22^{+0.17}_{-0.17}$ | $0.34^{+0.12}_{-0.09}$ | $-2.08^{+0.54}_{-0.37}$ | $0.24^{+0.09}_{-0.08}$ | $<8.95$ | - | -
E-core | $<5$ | $8.53^{+0.15}_{-0.13}$ | $0.72^{+0.15}_{-0.14}$ | $-1.82^{+0.49}_{-0.29}$ | $0.33^{+0.11}_{-0.10}$ | $<8.87$ | - | -
W-core | $35^{+29}_{-18}$ | $9.02^{+0.19}_{-0.20}$ | $0.30^{+0.14}_{-0.10}$ | $-1.99^{+0.67}_{-0.38}$ | $0.26^{+0.10}_{-0.10}$ | $<8.89$ | - | -

Note. — The upper and lower limits are the $2\sigma$ limits, that is, 98 and 2
percentiles of the post distributions, respectively.

####  4.3.1 Entire Merging System

The SED fit to the integrated flux supports that the B14-65666 system is a
young starburst, regardless of whether including the ALMA data. In the case
without the ALMA data, the stellar age and the stellar mass are
$t_{\text{SF}}=20^{+20}_{-10}$ Myr and $\log
M_{*}/M_{\sun}=9.6^{+0.3}_{-0.3}$, and then, instant star-formation rate is
$\mathrm{SFR}_{0}=M_{*}/t_{\text{SF}}=230^{+70}_{-60}$ $\mathrm{M_{\sun}\
yr^{-1}}$. The ionization parameter is $\log U=-2.1^{+0.5}_{-0.3}$. This young
starburst with the high ionization parameter gives rise to strong [O iii]
emission lines that can explain the F444W excess. The relatively red stellar
continuum even with the young starburst is attributed to the dust attenuation
of $A_{V}=0.86^{+0.1}_{-0.1}$ mag. The old stellar component is constrained to
be $\log M_{*}/M_{\sun}<9.7$ at the $2\sigma$ upper limit, which is comparable
to the stellar mass formed by the present star-formation. Thus, the starburst
induced by the major merger formed a stellar mass comparable to or higher than
the old stellar mass built before the on-going merger.

The SED fitting including the dust continua shows higher
$\mathcal{U}_{\mathrm{min}}$ and $\gamma$ values than those of nearby galaxies
($\mathcal{U}_{\mathrm{min}}\lesssim 10$ and $\gamma\lesssim 0.1$ in Draine &
Li 2007). Such high $\mathcal{U}_{\mathrm{min}}$ and $\gamma$ were also
obtained for UV-selected galaxies at $z>4.5$ (Burgarella et al. 2022),
indicating the necessity of high dust temperature to explain the FIR dust
emission in high-redshift galaxies. The characteristic dust
temperature333Within the given fitting parameter ranges of
$\mathcal{U}_{\mathrm{min}}$ and $\gamma$, the maximum characteristic
temperature is $52.4$ K., derived from Equation (33) in Draine & Li (2007) and
Equation (13) in Draine et al. (2014), is $T_{\text{d, char}}\simeq
18\langle\mathcal{U}\rangle^{1/6}\,\text{K}=48^{+2}_{-3}$ K, which is
consistent with the dust temperature obtained from the modified blackbody
fitting (Sugahara et al. 2021). Although including the ALMA data slightly
decreases the uncertainties of the best-fit dust attenuation and stellar
metallicity, the changes in the best-fit stellar properties are insignificant.

Previous analyses of the entire properties of B14-65666 yielded stellar masses
that are less than one-third of the one presented in this work (Bowler et al.
2018; Hashimoto et al. 2019). On the other hand, the instant
$\mathrm{SFR}_{0}$ that they obtained are consistent with our estimate
(Hashimoto et al. 2019). The stellar mass discrepancy is because of the
differences on the IRAC and NIRCam photometry, where IRAC flux densities were
$1.5\mbox{--}2$ times as low as the NIRCam flux densities at $3.6$ and $4.4$
$\mathrm{\mu m}$. We note again that our new photometry agrees with that in
the COSMOS2020 catalog, which is also higher than the previous photometry at
$3.6$ $\mathrm{\mu m}$ and $4.5$ $\mathrm{\mu m}$.

####  4.3.2 Resolved components

The high angular resolution of the NIRCam imaging enables us to perform the
SED fitting to each of the galaxy components. Figure 4 makes it clear that the
galaxies E and W have different colors in the rest-frame UV-to-optical
wavelengths. In the cases without the ALMA data, they are inferred as the
young starbursts with ages of $<53$ Myr and $40^{+40}_{-20}$ Myr,
respectively, similar to the total. The $\mathrm{SFR}_{10\text{Myr}}\
(=M_{*}(<10\,\text{Myr})/10\,\text{Myr})$ of the galaxy E ($180^{+50}_{-40}$
$\mathrm{M_{\sun}\ yr^{-1}}$) is $\simeq 4$ times as high as that of the
galaxy W ($40^{+10}_{-10}$ $\mathrm{M_{\sun}\ yr^{-1}}$). Moreover, the galaxy
E is expected to have $\gtrsim 0.5$ mag stronger dust attenuation than the
galaxy W. These results reflect a dusty, bursty phase of the galaxy E, which
is also supported by the strong F444W excess, i.e., the highest [O iii] line
strength at the core of the galaxy E (Section 4.1). In the SED fitting
including the ALMA data, the [O iii] 88 $\mathrm{\mu m}$ emission line fluxes
help more strictly constrain the metallicity of the galaxies; the galaxy E
exhibits a lower stellar metallicity than the galaxy W while they have a
similar metallicity in the SED fitting without the ALMA data. These
metallicity constraints originate from better constraints on the gas-phase
metallicity by the optical-to-FIR [O iii] line ratios. Both of the galaxies
show small Balmer jumps (or inverse Balmer break), a flux excess at
wavelengths shorter than $4000$ Å caused by the nebular continuum, which is
especially strong for the galaxy E. The contribution of this nebular continuum
and [O ii]$\lambda\lambda$3727,3730 emission lines help to boost the F277W
flux at a level similar to that of the F356W filter. Finally, the stellar-mass
ratio of the galaxies E and W is estimated to be 3:1 to 2:1. This mass ratio
is comparable to the dynamical mass ratio of the galaxy E ($5.7\pm 1.6\times
10^{10}$ $\mathrm{M_{\sun}}$) to the galaxy W ($3.1\pm 1.1\times 10^{10}$
$\mathrm{M_{\sun}}$) that are inferred from the [O iii] 88 $\mathrm{\mu m}$
line widths on the assumption of the virial theorem (Hashimoto et al. 2019).
Thus, this study has revealed that B14-65666 is surely a major merger in terms
of the stellar mass ratio of the galaxies, while previous studies identified
it from the luminosity ratio.

In the FIR SED, it is notable that the galaxies E and W have different peaks
of the dust continua. The NIRCam and ALMA observations have shown that the
galaxy E has redder UV spectral slope while the galaxy W exhibits higher dust-
continuum flux densities. The SED fitting explains these observations with a
difference of the dust temperatures. In particular, the galaxy E requires so
high $\mathcal{U}_{\mathrm{min}}$ and $\gamma$ that their probability
distribution functions reach the maximum values of their permitted parameter
ranges. The characteristic dust temperatures are $T_{\text{d,
char}}=50^{+1}_{-2}$ and $40^{+4}_{-3}$ K for the galaxies E and W,
respectively. We will further discuss relations between the UV and FIR light
in Section 5.1.

For finer resolved analyses, we divided the system into the cores of the
galaxies and the tails in the galaxy E. We did not include the ALMA data in
these SED fittings, which have relatively lower angular resolution. The panel
(c) of Figure 4 depicts the observed SED and best-fit spectra for the E-core
and W-core. The W-core has a similar SED shape to the galaxy W and is
$10\mbox{--}30$% fainter than it; this similarity means that the W-core is
dominant in the SED of the galaxy W. On the other hand, the E-core is twice as
faint as the galaxy E at the long wavelength channels. The obtained best-fit
age is very young, $<5$ Myr, at the $2\sigma$ upper limit and the stellar mass
is $\simeq 1.0$ dex lower than that of the galaxy E. As shown in the panels
(b) and (c), there are a $>0.4$ dex difference in the best-fit stellar
continua at wavelengths longer than the Balmer jump. These differences
indicate that the red tails strongly affect the SED of the galaxy E at the
rest-frame optical wavelengths. One possible scenario to interpret this result
might be that the galaxy E is composed of the dusty compact starbursts at the
nucleus and the older stars in the tails which dominate the stellar mass.
However, the non-detections in the F115W, F150W, and F200W bands of the tails
make it challenging to understand the origin of their red color by solving the
degeneracy between the old stellar populations, dust attenuation, and strong
emission lines. Detailed properties of the tails will be estimated by
including emission lines taken in approved JWST NIRSpec and MIRI observations.

##  5 Discussion

###  5.1 IRX–$\beta_{\text{UV}}$ relation of the clumps

The NIRCam and ALMA images have indicated complex morphology of B14-65666 in
the rest-frame UV, optical, and FIR views. While the galaxy W is bluer than
the galaxy E in the rest-frame UV (Figure 4), it exhibits brighter dust
emission (Figure 3). This originates from different properties of star
formation and dust attenuation between the galaxies E and W, as suggested in
the SED fitting.

We investigate the relation between the infrared excess
($\text{IRX}=L_{\text{IR}}/L_{\text{UV}}$) and the UV spectral slope
$\beta_{\text{UV}}$ to study whether the incongruity between the UV colors and
dust emission is due to the dust attenuation curve. We computed the
monochromatic rest-frame UV luminosity $L_{\text{UV}}$ and $\beta_{\text{UV}}$
from the best-fit SED models derived without the ALMA data. The UV luminosity
was computed from the UV continuum as $L_{\text{UV}}=\lambda L_{\lambda}$ at
$\lambda=1600$ Å. The UV slope was obtained by linear fitting to the UV
continuum within the wavelength windows proposed by Calzetti et al. (1994).
The measurement uncertainties were propagated from the fitting uncertainties;
we measured values from the 500 SED models derived from the parameter
posterior distributions and estimated uncertainties from the measurement
distribution. The IR luminosity $L_{\text{IR}}$ was taken from Hashimoto et
al. (2019), who estimated it by the modified blackbody fitting under the
assumption of the dust temperature of $T_{\text{d}}=48$ K and the dust
emissivity of $\beta_{\text{d}}=2.0$.

The obtained IRX values of the entire system (total) and both of the galaxies
are similar, $\log\text{IRX}\simeq 0.7$. However, the $\beta_{\text{UV}}$
values range from $\beta_{\text{UV}}=-2.2$ (galaxy W) to $-1.2$ (galaxy E).
The UV slope $\beta_{\text{UV}}$ depends on stellar and dust properties
including age, dust composition, and the stellar-and-dust geometry. Therefore,
the variation in $\beta_{\text{UV}}$ at a given IRX supports a spatial
variation of stellar and dust properties between the components.

Figure 5: IRX–$\beta_{\text{UV}}$ diagram. The black, red, and blue circles
show the measurements for the total, galaxy E, and galaxy W of B14-65666. The
E-core and W-core are depicted with the small circles with the same colors of
the galaxies E and W, respectively. The light-color triangles show the
measurements for B14-65666 obtained by Bowler et al. (2022). We note that the
galaxy-E symbol of Bowler et al. is completely overlapped with our E-core
symbol. The light gray symbols show $z\sim 7$ galaxies taken from Bowler et
al. (2022, diamond) and Hashimoto et al. (2019, square). All the IRX values of
the data points from the literature are corrected for $T_{\text{d}}=48$ K and
$\beta_{\text{d}}=2.0$. The lines show the parameterized
IRX–$\beta_{\text{UV}}$ relations: the black solid line is the Meurer et al.
(1999) relation modified with $\beta_{\text{UV,0}}=-2.62$ (Reddy et al. 2018),
the thin solid line is the original Meurer relation
($\beta_{\text{UV,0}}=-2.23$) for UV-selected local starbursts, the dotted
line is an aperture-corrected Meurer relation (Takeuchi et al. 2012), and the
dashed lines are relations obtained for UV-selected galaxies at $z\sim 2$
(Reddy et al. 2018) and $z\sim 4.5$ and $5.5$ (Fudamoto et al. 2020). The
bottom right arrows show how the IRX values change if the assumed dust
temperature increases or decreases from $48$ K.

Previously Bowler et al. (2022) measured the IRX and $\beta_{\text{UV}}$
values for B14-65666. In Figure 5, we plotted their data points after
correcting the IRX values for $T_{\text{d}}=48$ K and $\beta_{\text{d}}=2.0$.
Compared with their values, our measurements are systematically redder in
$\beta_{\text{UV}}$. Although the origin of the systematics is unclear, the
$\beta_{\text{UV}}$ difference would be caused by different spatial
resolution. Bowler et al. (2022) used UltraVISTA images with the spatial
resolutions of $\sim 0\farcs 7$, much coarser than those of our NIRCam images,
and they resolved objects with TPHOT (Merlin et al. 2015) based on an HST
image. Thus, unexpected systematics may be included in their rest-frame UV
photometry. This comparison underlines the power of high resolution for
spatially resolved studies.

The IRX–$\beta_{\text{UV}}$ relations are connected with attenuation curves
under the energy balance between absorption and re-radiation by dust (Meurer
et al. 1999). Canonical relations are parameterized with attenuation curves,
for example, for local starbursts (Meurer et al. 1999; Overzier et al. 2011;
Takeuchi et al. 2012) and for the Small Magellanic Cloud (SMC; Prevot et al.
1984; Gordon et al. 2003). For high-redshift galaxies, the intrinsic UV slope,
$\beta_{\text{UV,0}}$, is bluer than that for the local starbursts (Reddy et
al. 2018) due to the young ages, low stellar metallicities, and binary
interactions (Steidel et al. 2014, 2016), although the nebular continuum
emission could make it redder. We checked $\beta_{\text{UV,0}}$ of our best-
fit intrinsic SED curves before dust reddening and found that it ranges from
$\beta_{\text{UV,0}}=-2.6$ to $-2.55$, which is bluer than $-2.23$ used in
Meurer et al. (1999). For simplicity we use $\beta_{\text{UV,0}}=-2.62$ (Reddy
et al. 2018) as a fiducial value at high redshift.

Figure 5 compares the measurements for B14-65666 and various
IRX–$\beta_{\text{UV}}$ relations. The total is in a good agreement with the
widely-used canonical relation, the Meurer et al. (1999) relation, modified to
$\beta_{\text{UV,0}}=-2.62$. From this result, (1) the agreement with the
Meurer relation supports our assumption of using the Calzetti et al. (2000)
attenuation curve in the SED fitting and (2) $\beta_{\text{UV,0}}=-2.62$
supports lower stellar metallicity and younger ages of B14-65666 than those of
local starbursts (Meurer et al. 1999; Overzier et al. 2011; Takeuchi et al.
2012). The galaxies E and W are, however, not consistent with the
$\beta_{\text{UV,0}}$-modified Meurer relation. The galaxies E and W exhibit
lower and higher IRX than the relation at their $\beta_{\text{UV}}$,
respectively. Deviations from the canonical IRX-$\beta_{\text{UV}}$ relations
arise from several physical processes (e.g., Popping et al. 2017). For the
galaxy E, the low IRX value would be explained by hypotheses that the galaxy E
has higher dust temperatures of $T_{\text{d}}\geq 63\mbox{--}68$ K, which is
suggested by the SED fitting (Section 4.3.2); or that the galaxy E exhibits a
steeper SMC-like dust attenuation curve than the Calzetti et al. (2000) one.
Alternatively, the measurement of $\beta_{\text{UV}}$ based on the SED fitting
may be affected by the red tails. If we assume that the measured dust emission
only arises from the E-core, the E-core is located on the
$\beta_{\text{UV,0}}$-modified Meurer relation on the IRX–$\beta_{\text{UV}}$
diagram. For the galaxy W, hypotheses to explain the IRX values are a low dust
temperature of $T_{\text{d}}\leq 27\mbox{--}33$ K, which qualitatively agrees
with the SED-fitting result; or the patchy stellar-and-dust geometry
unresolved even with NIRCam and ALMA ($\lesssim 0.2\mbox{--}1.5$ kpc), which
produces flatter dust-attenuation curves resulting in high IRX at blue
$\beta_{\text{UV}}$. The patchy stellar-and-dust geometry has been already
found in a Lyman-break galaxy at $z=8.31$ (Tamura et al. 2023) and supported
by the FirstLight simulations, where complex stellar-and-dust geometry
produces high IRX (Mushtaq et al. 2023). In any scenarios, or in a combination
of several scenarios, our rest-frame UV and FIR observations have highlighted
the different dust properties between the galaxies E and W.

B14-65666 and most other $z\gtrsim 7$ galaxies (Hashimoto et al. 2019; Bowler
et al. 2022) are distributed between the Meurer relations with
$\beta_{\text{UV,0}}=-2.23$ (black thin line in Figure 5) and $-2.62$ (black
thick line). From this distribution, Hashimoto et al. (2019) and Bowler et al.
(2022) concluded that dust-detected galaxies at $z\gtrsim 7$ follow the
Calzetti et al. (2000) attenuation curve. The apparent IRX dispersion would be
explained by a combination of a $\beta_{\text{UV,0}}$ variation (i.e., a
variation of stellar age and metallicity) and patchy stellar-and-dust geometry
(Reddy et al. 2018). We note that their IRX is corrected for $T_{\text{d}}$
and $\beta_{\text{d}}$ used in this study and that their $\beta_{\text{UV}}$
was derived from photometry, which is different from our method. These results
are consistent with stacking analyses of Lyman-break galaxies at $z\sim 3$
(e.g., Álvarez-Márquez et al. 2016, 2019; Bowler et al. 2024), but seem to be
inconsistent with steeper attenuation curves favored in stacking analyses of
UV-selected galaxies at $z\sim 2$ (Reddy et al. 2018) and $z\sim 4.5$ and
$5.5$ (Fudamoto et al. 2020). However, some literature galaxies in Figure 5
have stringent IRX upper limits below the canonical relations at high
$\beta_{\text{UV}}$ and the IRX estimates significantly depend on the
assumptions of $T_{\text{d}}$ and $\beta_{\text{d}}$ and how to measure
$\beta_{\text{UV}}$. The “standard” IRX–$\beta_{\text{UV}}$ relation at high
redshift is still under debate.

Figure 6: Top: Map of the optical-to-FIR [O iii] emission line ratio, color
coded by $\log[\mbox{O\,{\sc iii}}]\ 4960+5008\ \mbox{\AA}/[\mbox{O\,{\sc
iii}}]\ 88\ \mbox{$\mathrm{\mu m}$}$. The white contours show the [O iii] 88
$\mathrm{\mu m}$ fluxes, which is the same as in Figure 3. The red and blue
crosses depict the peak positions of the E-core and W-core, respectively.
Bottom: the optical-to-FIR [O iii] emission line ratio as a function of the
electron density. The horizontal lines show the line ratios of the entire
system (black) and the value at the peak positions of the E-core (red) and
W-core (blue). The gray curves illustrate theoretical predictions computed
from Cloudy. The line thicknesses depend on the gas-phase metallicity and the
solid and dashed lines are at $\log U=-1.5$ and $-2.5$, respectively. At a
typical electron density ($\sim 200$ $\mathrm{cm^{-3}}$), the gas-phase
metallicity of the B14-65666 is $\simeq 0.2$ $Z_{\text{$\odot$}}$.

###  5.2 [O iii] Line Ratio: Insights into Gas-phase Metallicity

NIRCam images have identified a flux excess in the F444W band interpreted as a
strong [O iii]$\lambda\lambda$4960,5008 emission. The ALMA Band 8 has detected
the [O iii] 88 $\mathrm{\mu m}$ emission. Due to the large difference in their
excitation energy, the ratio of the optical to FIR [O iii] emission lines is
sensitive to the physical conditions (electron density and temperature) of the
ionized emitting gas (Osterbrock & Ferland 2006). This line ratio also depends
on the gas-phase metallicity under the assumptions of photoionization models
(e.g., Nakazato et al. 2023). Thus, a combination of the NIRCam images and
ALMA moment-0 map provides us with insights into spatial distributions of gas-
phase metallicity.

We created a line-ratio map by taking a ratio of the optical [O
iii]$\lambda\lambda$4960,5008 emission-line map to the FIR [O iii] 88
$\mathrm{\mu m}$ image. The optical [O iii] line map was created as in Section
4.1, but here the F356W and F444W images were corrected for the dust
attenuation before creating the line map. We corrected flux densities within
the apertures E ($A_{\text{V}}=1.2$ mag), W ($0.34$ mag), and E-core ($0.72$
mag) so that the flux density recovered in each aperture agreed with the best-
fit $A_{\text{V}}$ value obtained in the SED fitting. The created attenuation-
corrected optical line map was convolved to match the angular resolution of
the ALMA synthesized beam ($0\farcs 38$). The map area is limited to the
region where both lines are detected at more than 1$\sigma$ level and it
corresponds to $3\mbox{--}4$ beam solid angle.

The top panel of Figure 6 illustrates the optical-to-FIR [O iii] line-ratio
map. The line ratio at the peak of the [O iii] 88 $\mathrm{\mu m}$ surface
brightness is $\simeq 7.8$, and $10.6\pm 1.9$ for the integrated flux, where
the error does not include the uncertainty of $A_{\text{V}}$. The figure shows
the spatial variation of the line ratio over the galaxies, which suggests a
variation of the electron temperature and/or electron density in the galaxies
E and W. At the peak positions of the E-core and W-core, the [O iii] line
ratios are $9.2$ and $5.6$, respectively. The highest line ratios are seen in
the S-tail; however, these high values are based on the assumed strong dust
correction and include large uncertainties.

To be more quantitative, we used Cloudy to explore nebular parameters that
reproduce the obtained [O iii] line ratios. In our photoionization models, the
[O iii] line ratio depends on nebular parameters: the gas-phase metallicity
$Z_{\text{gas}}$, ionization parameter $U$, and hydrogen density
$n_{\text{H}}$. As the F356W band includes weak emission lines, which are also
subtracted from F444W flux, we corrected the model [O
iii]$\lambda\lambda$4960,5008 line fluxes by subtracting expected
$\mathrm{H\gamma}$ and $\mathrm{H\delta}$ line fluxes. The bottom panel of
Figure 6 shows optical-to-FIR [O iii] line ratios as a function of the nebular
parameters. The illustrated model lines depend on $Z_{\text{gas}}$ and $U$;
low gas-phase metallicity and high ionization parameter increase the electron
temperature of the ionized gas, leading to strong [O
iii]$\lambda\lambda$4960,5008 lines with respect to [O iii] 88 $\mathrm{\mu
m}$ line. Instead of $n_{\text{H}}$ at the illuminated surface, we used the
averaged electron density $n_{\text{e}}$ for the $x$-axis that was computed
from the modeled [O iii] 52 and 88 $\mathrm{\mu m}$ line ratios at the
electron temperature of $1.2\times 10^{4}$ K using PyNeb (Luridiana et al.
2015).

The estimated gas-phase metallicity is $Z_{\text{gas}}\sim 0.2$
$Z_{\text{$\odot$}}$ for the total, at an assumed electron density of
$n_{\text{e}}\sim 200$ $\mathrm{cm^{-3}}$, which is a typical estimate from
optical and FIR emission lines at high redshift (Sanders et al. 2016; Fujimoto
et al. 2022; Stiavelli et al. 2023). Using the line ratios at the peak
positions of the E-core and W-core, the E-core exhibits a lower gas-phase
metallicity of $Z\sim 0.3$ $Z_{\text{$\odot$}}$ than the W-core ($>0.4$
$Z_{\text{$\odot$}}$). These results also confirm conclusions of the previous
studies using the [O iii] 88 $\mathrm{\mu m}$ emission line that B14-65666 has
sub-solar metallicity but not extremely low at $<0.1$ $Z_{\text{$\odot$}}$
(Hashimoto et al. 2019; Jones et al. 2020b; Yang & Lidz 2020). In terms of the
electron density, the optical-to-FIR [O iii] line ratios are not consistent
with the presence of a high electron density gas ($n_{\text{e}}>10^{4}$
$\mathrm{cm^{-3}}$), which have been traced by high-ionization UV emission
lines in some high-redshift galaxies (e.g., Senchyna et al. 2023; Marques-
Chaves et al. 2024), under the assumption that galaxies in the EoR would have
sub-solar metallicity.

###  5.3 Overall scenario for B14-65666 system

The NIRCam and ALMA observations have shed light on the morphology and
physical properties of B14-65666 in the EoR. We analyzed the rest-frame
optical and FIR data to show the complex morphology of the merging galaxies:
the galaxy E is composed of the compact core with faint dust emission and the
diffuse red tails; the galaxy W presents an elongated and clumpy morphology
with bright dust emission. These clumpy, elongated structures and diffuse
emission like tidal tails are reproduced in the simulations of high redshift
galaxies ($z=6\mbox{--}9$, Nakazato et al. 2024). Furthermore, the galaxy W
may exhibit patchy stellar-and-dust geometry that is not resolved in our
observations. These results indicate that B14-65666 is experiencing strong
tidal forces accompanied with a merger, disturbing the entire morphology of
the system and each of the galaxies involved in the process.

The projected separation of the galaxies E and W is $d\simeq 4$ kpc and their
line-of-sight velocity difference is $v\simeq 200$ $\mathrm{km\ s^{-1}}$
(Hashimoto et al. 2019). Although this is the first-order estimate, the
timescale elapsed from the previous passage is estimated to be $d/v\simeq 20$
Myr, which is comparable to the star-formation age inferred from the SED
fitting. This merging timescale is feasible because it is shorter than the
expected merger observability timescale of $180$ Myr at $z=7.15$ (Snyder et
al. 2017). This consistency supports that the gravitational interactions have
driven the star formation in the B14-65666 after the first (or second)
passage.

Major mergers would promote gas accretion by removing the angular momentum
from disk gas, triggering starbursts at the cores of the galaxies (e.g., Mihos
& Hernquist 1994). B14-65666 shows signatures of nuclear dusty starbursts,
including strong optical [O iii] emission, strong dust attenuation, and high
$\mathrm{SFR}$. These nuclear dusty starbursts have indeed enhanced the star-
forming activity of B14-65666. The instant $\mathrm{SFR}_{0}$ of $230$
$\mathrm{M_{\sun}\ yr^{-1}}$ is $\simeq 1$ dex higher than the star-formation
main sequence at $\log M_{*}/\mathrm{M_{\sun}}=9.7$ and $z\sim 7$ (Popesso et
al. 2023). The merger-induced star-formation at high redshift is consistent
with a finding in the FirstLight simulations that bursty star-formation is
induced by gas-rich mergers at $z>5$ (Ceverino et al. 2018). While many
observations demonstrate that major mergers enhance $\mathrm{SFR}$ in the
local Universe (e.g., Ellison et al. 2008), the $\mathrm{SFR}$ enhancement is
reported to be inefficient at $z>1$ (e.g., Silva et al. 2018; Shah et al.
2022). B14-65666 is a clear example experiencing enhanced star formation at
higher redshift.

In particular, the E-core experiences a compact, centrally concentrated,
strong starburst. Given $r_{\text{e}}<85$ pc of the E-core in the F115W image
and $\mathrm{SFR}_{0}$ values in Table 3, the $\mathrm{SFR}_{0}$ surface
density of the E-core is $\Sigma_{\mathrm{SFR},0}\gtrsim 4\times 10^{3}$
$\mathrm{M_{\sun}\ yr^{-1}\ kpc^{-2}}$, which is comparable to IR
$\mathrm{SFR}$ surface densities of local ULIRGs or obscured AGN (Pereira-
Santaella et al. 2021). We expect that this compact starburst is caused by a
significant gas accretion to the E-core. The high $\Sigma_{\mathrm{SFR},0}$ is
also consistent with the expected high dust temperature of the E-core (see
Liang et al. 2019). We note that star-formation timescale of the E-core ($<5$
Myr) is much shorter than those of ULIRGs, as $\mathrm{SFR}$ of ULIRGs are
derived from the FIR luminosity (i.e., $\sim 100$ Myr). Although the molecular
gas mass is highly uncertain for B14-65666 ($10^{8.7\mbox{--}11}$
$\mathrm{M_{\sun}}$, Hashimoto et al. 2023a), the high
$\Sigma_{\mathrm{SFR},0}$ would lead to short gas depletion time of $<10$ Myr
in the range of molecular gas mass surface densities of local starbursts
($<10^{4}$ $\mathrm{M_{\sun}\ pc^{-2}}$; Hodge et al. 2015; Pereira-Santaella
et al. 2021). In addition, the high $\Sigma_{\mathrm{SFR},0}$ and
s$\mathrm{SFR}_{0}$ of the E-core suggest that the radiation presssure is high
enough to launch galactic outflows (e.g., Pereira-Santaella et al. 2021;
Ferrara 2023), which lead to negative feedback, therefore reducing or even
quenching the star formation (Hopkins et al. 2012). Thus, the compact dusty
starburst at the E-core may cease within several tens Myr.

The gas-phase metallicity for the entire system is estimated to be $\simeq
0.2$ $Z_{\text{$\odot$}}$ from the SED fitting, and these estimates are
supported by the optical-to-FIR [O iii] emission line ratio. B14-65666 is on
the average mass-metallicity relation at $z=4\mbox{--}9$ (Nakajima et al.
2023; Curti et al. 2023), but given its high $\mathrm{SFR}$, it would not be
on the fundamental metallicity plane (Andrews & Martini 2013). Our analyses
also show that the galaxy E has lower metallicity than the galaxy W, which is
consistent with a scenario where a strong accretion of less-enriched gas
dilutes the pre-existing gas at the galaxy center (e.g., Kewley et al. 2006).
At $Z_{\text{gas}}\simeq 0.2$ $Z_{\text{$\odot$}}$, an ionization parameter
estimated by Sugahara et al. (2021) using the [O iii] 88 $\mathrm{\mu m}$ and
[C ii] 158 $\mathrm{\mu m}$ lines is $\log U\simeq-2.0$, which is consistent
with the SED fitting results. In other words, high [O iii]$/$[C ii] ratio of
B14-65666 can be naturally explained by the high ionization parameter (see
Hashimoto et al. 2019; Harikane et al. 2020). The [N ii] 122 $\mathrm{\mu m}$
non-detection is also easily explained by the high ionization parameter where
most of the nitrogen are in the doubly ionized phase (Sugahara et al. 2021).

Finally, we mention some limitations of our observations. Although we have
revealed many properties of the cores of the galaxies, less properties are
constrained for the tails. This is mainly because it is difficult to solve the
degeneracy in physics that can explain their red SED, dust non-detections, and
detection of [O iii] 88 $\mathrm{\mu m}$ emission, with different angular
resolutions of the NIRCam and ALMA. We expected that the extended diffuse
emission originates from the tidal tails produced by the major merger, but we
did not reject other possibilities like ionized gas outflows (e.g., Yuma et
al. 2019) and offset old stellar populations (e.g., Colina et al. 2023).
Approved MIRI and NIRSpec IFU spectroscopy and ALMA high angular-resolution
observations will give a hint to properties of the tails. To reach statistical
conclusions, a large sample of major mergers at high redshift is necessary. It
is unclear whether B14-65666 is representative of major mergers at this
redshift. Nevertheless, our results suggest that the major merger, which is an
abundant population at high redshift (Rodriguez-Gomez et al. 2015; Duncan et
al. 2019; Romano et al. 2021), makes significant contributions to galaxy mass
assembly in the early Universe. B14-65666 is still an important target that
provides us detailed distributions of stellar, gaseous, and dust populations
during major mergers.

##  6 Conclusions

In this paper, we present JWST NIRCam observations of the UV-bright
($M_{\text{UV}}=-22.5$ mag) Lyman-break galaxy system B14-65666 at $z=7.1520$,
also known as Big Three Dragons. The NIRCam filters are selected to sample the
rest-frame UV wavelengths, the Balmer break clean of strong optical emission
lines, and the rest-frame optical wavelengths ($\sim 5000$ Å) including
$\mathrm{H\beta}$ and [O iii] emission lines. Additionally, rich ancillary
ALMA data trace the [O iii] 88 $\mathrm{\mu m}$ and [C ii] 158 $\mathrm{\mu
m}$ emission lines and underlying dust continua. The high angular resolution
of these NIRCam and ALMA observations enables spatially resolved analyses at
the rest-frame UV, optical, and FIR wavelengths of B14-65666, a bright major
merger in the EoR.

B14-65666 consists of two galaxy components, E and W, and diffuse emission
surrounding the galaxy E. The galaxy E is red (UV spectral slope
$\beta_{\text{UV}}=-1.2$) and has a compact core (E-core) with an effective
radius of $r_{\text{e}}<0.016$ arcsec ($85$ pc) at the rest-frame $1400$ Å.
The galaxy W is blue ($\beta_{\text{UV}}=-2.2$) and elongated to $\simeq 0.3$
arcsec ($1.5$ kpc) with clumpy morphology in the rest-frame UV wavelengths.
Surprisingly, in contrast to their colors, the blue galaxy W is brighter than
the red galaxy E in the dust continua at $88$ and $158$ $\mathrm{\mu m}$. The
core of the galaxy E (E-core) is surrounded by diffuse extended emission,
which is only detected in the long NIRCam wavelength channels and is also
bright in [O iii] 88 $\mathrm{\mu m}$. We assume that this extended emission
is likely to be the tidal tails created by the gravitational interactions.
Both of the galaxies E and W show the F444W flux excesses over the F356W
fluxes ($0.84^{+0.09}_{-0.10}$ and $0.53^{+0.09}_{-0.10}$ mag, respectively),
indicating the strong [O iii]$\lambda\lambda$4960,5008 emission lines. The [O
iii] line map estimated from the F356W$-$F444W color illustrates strong [O
iii] emission at the E-core and W-core, with the rest-frame equivalent widths
of $\text{EW}_{0}\simeq 2000$ and $750$ Å, respectively. These results show
that the major merger has disturbed the morphology of the galaxies and induced
the nuclear dusty starbursts.

The SED fitting shows that B14-65666 has a total stellar mass of $\log
M_{*}/\mathrm{M_{\sun}}=9.8\pm 0.2$ and $\mathrm{SFR}_{\text{10Myr}}$ of
$\simeq 230^{+70}_{-60}$ $\mathrm{M_{\sun}\ yr^{-1}}$. These merger-induced
starbursts would build young stellar masses comparable to or heavier than the
underlying old stellar components. The stellar-mass ratio of the galaxies E
and W spans from 3:1 to 2:1, which confirms the fact that B14-65666 is a major
merger in terms of the stellar-mass ratio. The galaxy E experiences young
starbursts with $\mathrm{SFR}_{\text{10Myr}}=180^{+50}_{-40}$
$\mathrm{M_{\sun}\ yr^{-1}}$ at ages of $<53$ Myr, supported by the high
optical [O iii] emission. The red color of the galaxy E is explained by the
strong dust attenuation of $A_{\text{V}}=1.2\pm 0.1$ mag, in opposite to the
galaxy W, which has blue color and $A_{\text{V}}=0.3\pm 0.1$ mag. Including
ALMA data, the SED fitting explains the red color and faint dust continua of
the galaxy E by the scenario that the galaxy E exhibits higher dust
temperature than the galaxy W.

As combinations of optical and FIR data, we investigate the
IRX-$\beta_{\text{UV}}$ diagram and the optical-to-FIR [O iii] line ratios.
The entire system is consistent with the $\beta_{\text{UV,0}}$-modified Meurer
relation, supporting the Calzetti attenuation law. The galaxies E and W,
however, deviate from the relation; this deviation can be explained by
different stellar and dust properties: variations of the dust attenuation law,
patchy stellar-and-dust geometry, different dust temperature, and possible
effects of the tails. From the estimated [O iii]$\lambda\lambda$4960,5008 map
and the [O iii] 88 $\mathrm{\mu m}$ moment-0 maps, we draw the optical-to-FIR
[O iii] line ratio map. This map illustrates gradual spatial variation in the
[O iii] line ratio, implying spatial variation in the electron temperature and
electron density. From the [O iii] line ratios, the gas-phase metallicity of
the entire galaxy is estimated to be $Z_{\text{gas}}\sim 0.2$
$Z_{\text{$\odot$}}$ by photoionization models under the assumption of the
electron density of $\sim 200$ $\mathrm{cm^{-3}}$. The E-core shows higher [O
iii] line ratio than the W-core, indicating lower gas-phase metallicity in the
E-core.

In summary, our results support a scenario where B14-65666 is a major merger
that gravitationally disturbs the morphology and experiences nuclear dusty
starbursts likely triggered by less-enriched gas inflows, after the first (or
second) passage. Although the galaxies E and W have similar UV magnitude, they
exhibit different stellar, gaseous, and dust properties and morphology. The
starbursts have indeed enhanced the star-formation activity compared with
main-sequence star-forming galaxies at the same redshift. B14-65666 is an
important object that provides us opportunities to investigate complex stellar
buildup processes during major mergers, which are abundant at high redshifts,
during the EoR.

We thank Rebecca Bowler for helpful discussions on B14-65666. This research is
supported by NAOJ ALMA Scientific Research Grant number 2020-16B. JAM, ACG,
and LC acknowledge support by grant PIB2021-127718NB-100 from the Spanish
Ministry of Science and Innovation/State Agency of Research
MCIN/AEI/10.13039/501100011033 and by “ERDF A way of making Europe”. TH was
supported by Leading Initiative for Excellent Young Researchers, MEXT, Japan
(HJH02007) and by JSPS KAKENHI Grant Number 22H01258. KM, YN, and YT
acknowledge financial support from JSPS through KAKENHI grant Number 20K14516,
23KJ0728, and 22H04939, respectively. CBP acknowledges support by grant
CM21_CAB_M2_01 from the Program “Garantía Juveníl” from the “Comunidad de
Madrid” 2021. MPS acknowledges support from grant RYC2021-033094-I funded by
MCIN/AEI/10.13039/501100011033 and the European Union NextGenerationEU/PRTR.
This paper makes use of the following ALMA data: ADS/JAO.ALMA#2016.1.00954.S,
ADS/JAO.ALMA#2017.1.00190.S, and ADS/JAO.ALMA#2019.1.01491.S. ALMA is a
partnership of ESO (representing its member states), NSF (USA), and NINS
(Japan), together with NRC (Canada), MOST and ASIAA (Taiwan), and KASI
(Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA
Observatory is operated by ESO, AUI/NRAO, and NAOJ. This work has made use of
data from the European Space Agency (ESA) mission Gaia
(https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and
Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement. The Gaia data are retrieved
from the JVO portal (http://jvo.nao.ac.jp/portal) operated by the NAOJ. This
research has made use of NASA’s Astrophysics Data System.

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