# The Detection of Higher-Order Millimeter Hydrogen Recombination Lines in the
Large Magellanic Cloud

Marta Sewiło Exoplanets and Stellar Astrophysics Laboratory, NASA Goddard
Space Flight Center, Greenbelt, MD 20771, USA  Department of Astronomy,
University of Maryland, College Park, MD 20742, USA  Center for Research and
Exploration in Space Science and Technology, NASA Goddard Space Flight Center,
Greenbelt, MD 20771  Kazuki Tokuda Department of Earth and Planetary Sciences,
Faculty of Sciences, Kyushu University, Nishi-ku, Fukuoka 819-0395, Japan
National Astronomical Observatory of Japan, National Institutes of Natural
Sciences, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan  Stan E. Kurtz Instituto
de Radioastronomía y Astrofísica, Universidad Nacional Autónoma de México,
Apdo. Postal 3-72, 58090 Morelia, Michoacán, Mexico  Steven B. Charnley
Astrochemistry Laboratory, NASA Goddard Space Flight Center, Greenbelt, MD
20771, USA  Thomas Möller I. Physikalisches Institut der Universität zu Köln,
Zülpicher Str. 77, 50937, Köln, Germany  Jennifer Wiseman Exoplanets and
Stellar Astrophysics Laboratory, NASA Goddard Space Flight Center, Greenbelt,
MD 20771, USA  C.-H. Rosie Chen Max-Planck-Institut für Radioastronomie, Auf
dem Hügel 69 D-53121 Bonn, Germany  Remy Indebetouw Department of Astronomy,
University of Virginia, PO Box 400325, Charlottesville, VA 22904, USA
National Radio Astronomy Observatory, 520 Edgemont Rd, Charlottesville, VA
22903, USA  Álvaro Sánchez-Monge Institut de Ciències de l’Espai (ICE, CSIC),
Can Magrans s/n, E-08193, Bellaterra, Barcelona, Spain  Institut d’Estudis
Espacials de Catalunya (IEEC), Barcelona, Spain  Kei E. I. Tanaka Tokyo
Institute of Technology, 2 Chome-12-1 Ookayama, Meguro City, Tokyo 152-8550,
Japan  Peter Schilke I. Physikalisches Institut der Universität zu Köln,
Zülpicher Str. 77, 50937, Köln, Germany  Toshikazu Onishi Department of
Physics, Graduate School of Science, Osaka Metropolitan University, 1-1
Gakuen-cho, Naka-ku, Sakai, Osaka 599-8531, Japan  Naoto Harada Department of
Earth and Planetary Sciences, Faculty of Sciences, Kyushu University, Nishi-
ku, Fukuoka 819-0395, Japan

###### Abstract

We report the first extragalactic detection of the higher-order millimeter
hydrogen recombination lines ($\Delta n>2$). The $\gamma$-, $\epsilon$-, and
$\eta$-transitions have been detected toward the millimeter continuum source N
105–1 A in the star-forming region N 105 in the Large Magellanic Cloud (LMC)
with the Atacama Large Millimeter/submillimeter Array (ALMA). We use the
H40$\alpha$ line, the brightest of the detected recombination lines
(H40$\alpha$, H36$\beta$, H50$\beta$, H41$\gamma$, H57$\gamma$, H49$\epsilon$,
H53$\eta$, and H54$\eta$), and/or the 3 mm free-free continuum emission to
determine the physical parameters of N 105–1 A (the electron temperature,
emission measure, electron density, and size) and study ionized gas
kinematics. We compare the physical properties of N 105–1 A to a large sample
of Galactic compact and ultracompact (UC) H ii regions and conclude that N
105–1 A is similar to the most luminous ($L>10^{5}$ $L_{\odot}$) UC H ii
regions in the Galaxy. N 105–1 A is ionized by an O5.5 V star, it is deeply
embedded in its natal molecular clump, and likely associated with a
(proto)cluster. We incorporate high-resolution molecular line data including
CS, SO, SO2, and CH3OH ($\sim$0.12 pc), and HCO+ and CO ($\sim$0.087 pc) to
explore the molecular environment of N 105–1 A. Based on the CO data, we find
evidence for a cloud-cloud collision that likely triggered star formation in
the region. We find no clear outflow signatures, but the presence of filaments
and streamers indicates on-going accretion onto the clump hosting the UC H ii
region. Sulfur chemistry in N 105–1 A is consistent with the accretion shock
model predictions.

##  1 Introduction

The hydrogen radio and millimeter/submillimeter recombination lines (RRL and
mm-RLs, respectively) are excellent tools to measure the temperature and
column density of ionized gas in star-forming regions, and provide information
on the ionized gas kinematics (e.g., Dupree & Goldberg 1970; Brown et al.
1978; Gordon & Sorochenko 2002; Tanaka et al. 2016). They allow us to peer
into the dense and dusty regions of the molecular clouds inaccessible through
observations at other wavelengths.

Recombination lines are formed by a radiative recombination in which a free
electron recombines with an ion. In this process, the electron ends up in any
of the bound electronic states and the atom emits a photon that carries away
the energy lost by the electron. Following the recombination, the electron
cascades down the energy levels from the level it recombined to (with the
principal quantum number $n$$>$1) to the ground state ($n$=1) by a sequence of
spontaneous emissions, producing a series of emission lines (the recombination
lines) at wavelengths that depend on the difference in energy between the
levels. For higher $n$, this difference in energy is smaller, and thus the
atom emits at longer wavelengths. If the electron recombines to a highly-
excited state (large $n$), the process will produce spectral lines from radio
to ultraviolet (UV) wavelengths. The recombination to the ground state of the
hydrogen atom emits the Lyman $\alpha$ (Ly$\alpha$) spectral line in the UV,
the strongest emission line observed toward astrophysical objects. Atoms other
than hydrogen produce recombination lines as well (such as helium, carbon, and
oxygen); however, they have much lower abundances compared to hydrogen and
thus the recombination lines are weaker (e.g., Gordon & Sorochenko 2002).
Recombination lines can also be emitted by ions; the detection of the He ii, C
ii, and O ii recombination lines has been reported in the literature (e.g.,
Chaisson & Malkan 1976; Liu et al. 2023).

The RRL/mm-RL transitions are identified by the name of the element, the
principal quantum number of the final level, and the change in the principal
quantum number ($\Delta n$) indicated with successive letters in the Greek
alphabet. The recombination lines of H from level $n+\Delta n$ to $n$ are
denoted (H$n\alpha$, H$n\beta$, H$n\gamma$, H$n\delta$, H$n\epsilon$,
H$n\zeta$, H$n\eta$…) for $\Delta n$ = (1, 2, 3, 4, 5, 6, 7…). The
$\alpha$-transitions ($\Delta n$=1) have the largest Einstein coefficient for
spontaneous emission ($A$) and thus are the most likely transitions, producing
the brightest recombination lines that can be detected at large distances. The
hydrogen RL $\alpha$\- and $\beta$-transitions were first detected outside the
Galaxy in the 1970s. The higher-order ($\Delta n$$>$2) mm/radio transitions
had not been detected until our observations reported in this paper.

The first extragalactic detection of radio recombination lines was reported by
Mezger et al. (1970). The H109$\alpha$, He109$\alpha$, and H137$\beta$
transitions at 6 cm (5 GHz) were detected toward the star-forming region 30
Doradus in the Large Magellanic Cloud (LMC) with the Parkes 64 m radio
telescope (half-power beam width, HPBW$\approx$4′). The H109$\alpha$ line was
later detected toward seven other regions in the LMC in addition to 30 Dor
(out of 14 surveyed) by McGee et al. (1974). In the higher resolution
observations of 30 Dor with the Australia Telescope Compact Array (ATCA;
HPBW$\sim$15${}^{{}^{\prime\prime}}$), Peck et al. (1997) detected the
H90$\alpha$, H113$\beta$, and He90$\alpha$ ($\sim$8.9 GHz), H92$\alpha$ (8.3
GHz), and again the H109$\alpha$ (5 GHz) recombination lines. 30 Dor has been
the LMC star-forming region best-studied in RRLs. The H30$\alpha$ and
H40$\alpha$ lines are now routinely detected toward the LMC star-forming
regions with the Atacama Large Millimeter/submillimeter Array (ALMA), in
programs targeting cold molecular gas, such as CO (e.g., Indebetouw et al.
2013: 30 Doradus; Saigo et al. 2017: N 159–East; Nayak et al. 2019: N 79). The
only detection of a carbon RL in the LMC (C30$\alpha$ in N 79) is uncertain
since it may in fact be a helium RL (He30$\alpha$).

Figure 1: The 1.2 mm (left) and 3 mm (right) continuum images of N 105–1 A. In
the left/right panel, white contours correspond to (3, 10, 30, 250)/(3, 20,
100) times $6.9\times 10^{-5}$/$2.5\times 10^{-4}$ Jy beam-1, the 1.2/3 mm
continuum image rms noise ($\sigma$). Red contours in the right panel are the
same as white contours in the left panel. The size of the ALMA synthesized
beam is indicated in the bottom left corner in each image: $0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}51\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}47$ for 1.2 mm and $2\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}49\times 1\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}83$ for the 3 mm continuum image,
respectively.

The first extragalactic detections of RRLs beyond the Magellanic Clouds were
made toward M 82 (H166$\alpha$, Shaver et al. 1977; H102$\alpha$, Bell &
Seaquist 1977; H92$\alpha$, Chaisson & Rodriguez 1977) and NGC 253
(H102$\alpha$, Seaquist & Bell 1977). After these first observations of
extragalactic RRLs, no new detections were reported until the H53$\alpha$ mm-
RL was detected toward NGC 2146 over 10 years later (Puxley et al. 1991). RLs
in the millimeter and submillimeter range have proven to be excellent tools in
studying the star formation activity in dusty regions in the centers of nearby
galaxies (e.g., Scoville & Murchikova 2013). ALMA observations have provided
new detections of the H$n\alpha$ mm-RLs toward nearby galaxies (e.g., Bendo et
al. 2015; Bendo et al. 2016; Bendo et al. 2017; Michiyama et al. 2020) and the
first detection of the fainter H$n\beta$ and He$n\alpha$ lines toward NGC 253
(Meier et al. 2015; Martín et al. 2021). At cosmological distances, Emig et
al. (2019) claimed the first detection of the RRL after stacking 13
$\alpha$-transitions with principal quantum numbers $n=266-301$, detected in
the spectrum of the radio quasar 3C 190 ($z=1.1946$) with LOw Frequency ARray
(LOFAR; van Haarlem et al. 2013).

In this paper, we report the detection of the H40$\alpha$, H36$\beta$,
H50$\beta$, H41$\gamma$, H57$\gamma$, H49$\epsilon$, H53$\eta$, and H54$\eta$
mm-RLs and a tentative detection of the H55$\theta$ line with ALMA toward the
1.2 mm continuum source N 105–1 A in the N 105 star-forming region in the LMC
(LHA 120–N 105, Henize 1956; DEM L86, Davies et al. 1976). The LMC is the
Milky Way’s satellite galaxy, the nearest star-forming galaxy, located at the
distance of $49.59\pm 0.09$ (statistical) $\pm 0.54$ (systematic) kpc
(Pietrzyński et al. 2019). The (sub)mm-RL $\gamma$-, $\epsilon$-, $\eta$-, and
$\theta$-transitions are detected outside the Galaxy for the first time.

The environment of the LMC is distinct from that found in our Galaxy - it is
characterized by a low metallicity (lower abundances of gaseous atoms heavier
than He; $Z_{\rm LMC}$ $\sim$ 0.3–0.5 $Z_{\odot}$, e.g., Russell & Dopita
1992; Westerlund 1997; Rolleston et al. 2002), lower dust-to-gas ratio (e.g.,
Dufour 1975, 1984; Koornneef 1984; Roman-Duval et al. 2014), higher intensity
of the interstellar UV radiation (e.g., Browning et al. 2003; Welty et al.
2006), and lower cosmic-ray density (e.g., Abdo et al. 2010; Knödlseder 2013).
Metal (such as C, O, and N) abundances are expected to set the electron
temperature of H ii regions since the collisionally excited lines from metals
are the primary cooling mechanism in the ionized gas (e.g., Shaver et al.
1983; Balser et al. 2011 and references therein). Thus, low abundances of
metals in the LMC are expected to produce high temperatures. All the
properties of the LMC’s environment can affect the chemistry of the sources
emitting mm-RLs/RRLs. The relatively small distance to the LMC enables studies
of individual stars and protostars in this unique environment that is similar
to galaxies at the peak of star formation in the Universe ($z$$\sim$1.5–2;
e.g., Pei et al. 1999; Mehlert et al. 2002; Madau & Dickinson 2014).

The paper is organized as follows. In Section 2, we describe our ALMA
observations and the archival data used in our analysis. We summarize previous
studies on N 105–1 A in Section 3. The data analysis results and discussion
are presented in Section 4 and Section 5, respectively. In Section 6, we
provide the summary and conclusions of our study.

##  2 ALMA Observations

The N 105–1 field hosting the source with the detection of mm-RLs (N 105–1 A)
was observed with ALMA 12 m Array in Band 6 as part of the Cycle 7 project
2019.1.01720.S (PI M. Sewiło). The data were calibrated with version 5.6.1-8
of the ALMA pipeline in CASA (Common Astronomy Software Applications; CASA
Team et al. 2022). The observations were executed twice on October 21, 2019
with 43 antennas and baselines from 15 m to 783 m. The (bandpass, flux, phase)
calibrators were (J0519$-$4546, J0519$-$4546, J0440$-$6952) and (J0538$-$4405,
J0538$-$4405, J0511$-$6806) for the first and second run, respectively. N
105–1 was observed again on October 23, 2019 with 43 antennas, baselines from
15 m to 782 m, and the same calibrators. The total on-source integration time
was $\sim$13.1 minutes including both executions.

The spectral setup included four 1875 MHz spectral windows centered on
frequencies of 242.4 GHz, 244.8 GHz, 257.9 GHz, and 259.7 GHz. Continuum was
subtracted in the uv domain from the line spectral window and imaged. The CASA
task tclean was used for imaging using the Hogbom deconvolver, standard
gridder, Briggs weighting with a robust parameter of 0.5, and auto-multithresh
masking. The rms sensitivity of 0.05 mJy per 0$\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}51\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}47$ beam (4.4 mK) was achieved in the
continuum. Sensitivity of 1.97 mJy (0.15 K) per 0$\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}54\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}50$ beam was achieved in the 242.4 GHz
spectral cube; 1.88 mJy (0.15 K) per 0$\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}53\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}49$ beam in the 244.8 GHz cube; 2.05 mJy
(0.16 K) per 0$\hbox to0.0pt{.\hss}{{}^{\prime\prime}}51\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}47$ beam in the 257.9 GHz cube; 2.28 mJy
(0.18 K) per 0$\hbox to0.0pt{.\hss}{{}^{\prime\prime}}51\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}47$ beam in the 259.7 GHz cube; all four with
488.3 kHz (0.6 km s-1) channels. The images were corrected for primary beam
attenuation. The Band 6 molecular line data for the ALMA N 105–1 field were
analyzed and discussed in Sewiło et al. (2022).

In our analysis, we also utilize the ALMA Band 3 12m- and 7m-Array
observations of two young stellar objects (YSOs) located in N 105–1 (one
coinciding with 1 A; see Fig. C.1) from our Cycle 5 project 2017.1.00093.S (PI
T. Onishi). We used the data from observations of both YSOs for imaging since
the fields of view overlap considerably. In this paper, we analyze spectral
windows centered on the H40$\alpha$ line (a rest frequency of 99.02295 GHz;
12m data only to use with our Band 6 12m data) and 13CO (1–0) (110.20135 GHz,
an upper state energy, $E_{\rm U}=5.3$ K; combined 12m and 7m data).

The 12m observations toward N 105–1 were executed on April 2 and April 3, 2018
with 41 and 43 antennas and baselines between 15 m and 543 m and 15 m and 500
m, respectively. The (bandpass, flux, phase) calibrators were (J0635$-$7516,
J0635$-$7516, J0529$-$7245) for both runs. The total on-source integration
time was $\sim$4 minutes. The 1875 MHz spectral window centered on the
H40$\alpha$ line was divided into 1920 channels of $\sim$976.6 kHz ($\sim$2.96
km s-1) each. The 59 MHz spectral window centered on the 13CO (1–0) line was
divided into 1920 channels of $\sim$30.7 kHz ($\sim$0.084 km s-1) each. The
data were calibrated with CASA 5.1.1-5.

Continuum was subtracted in the uv domain from the H40$\alpha$ line spectral
window. Sensitivity of 4.4 mJy per $2\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}56\times 1\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}95$ beam (0.11 K) was achieved in the
H40$\alpha$ data cube. The H40$\alpha$ spectral cube has a cell size of
$0\hbox to0.0pt{.\hss}{{}^{\prime\prime}}3\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}3\times 2.96$ km s-1. Sensitivity of 0.25 mJy
per $2\hbox to0.0pt{.\hss}{{}^{\prime\prime}}49\times 1\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}83$ beam (6.8 mK) was achieved in the
continuum (99.023 GHz or 3.03 mm) constructed from the line-free channels. The
images were corrected for primary beam attenuation.

Figure 2: Left: Three-color mosaic of the northern part of the N 105 star-
forming region combining the H$\alpha$ image from the MCELS2 survey with the
MOSAIC II camera on the Blanco 4m Telescope at CTIO (red; a pixel size of
$0\hbox to0.0pt{.\hss}{{}^{\prime\prime}}27\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}27$, PI You-Hua Chu), the CS (5–4) image of
the ALMA field N 105–1 from Sewiło et al. (2022, green; an angular resolution
$\sim$0$\hbox to0.0pt{.\hss}{{}^{\prime\prime}}$5), and the ATCA 4.8 GHz / 6
cm image (blue; $\sim$2′′, Indebetouw et al. 2004). The white box indicates
the area shown in the image in the right panel. Right: Three-color mosaic of N
105–1 A combining the H36$\beta$ recombination line (red, this paper) and the
CS (5–4) (green; the same as in the left panel) and SO ${}^{3}\Sigma$ 66–55
(blue) images from Sewiło et al. (2022). All the images have the same angular
resolution of $\sim$0$\hbox to0.0pt{.\hss}{{}^{\prime\prime}}$5 ($\sim$0.12
pc). Green/blue contours correspond to the CS/SO emission with contour levels
of (20, 50, 90)% of the CS/SO integrated intensity peak of 230.8/262.8 mJy
beam-1 km s-1. White contours correspond to the 1.2 mm continuum emission with
contour levels of (3, 10, 100) times the image rms noise level of $6.9\times
10^{-5}$ Jy beam-1.

The 7m observations were executed eight times between December 29, 2017 and
January 10, 2018 with 10 (1 execution) and 11 (7) antennas, with baselines
between 8.9 m and 48.9 m. The total on-source integration time was $\sim$19.7
minutes for each target YSO. Both bandpass and flux calibrators were
J0522$-$3627 (7 executions) and J0006$-$0623 (1). J0450$-$8101 (4 executions)
and J0529$-$7245 (4) were used as phase calibrators. The 62 MHz spectral
window centered on the 13CO (1–0) line was divided into 2048 channels of
$\sim$30.3 kHz ($\sim$0.082 km s-1) each. We combined the 12m and 7m data to
image the 13CO (1–0) line. The 13CO (1–0) spectral cube has a cell size of
$0\hbox to0.0pt{.\hss}{{}^{\prime\prime}}3\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}3\times 0.2$ km s-1 and it is corrected for
primary beam attenuation. The sensitivity of 16.9 mJy per $2\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}29\times 1\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}67$ beam (0.44 K) was achieved in the 13CO
(1–0) data cube.

The 1.2 mm and 3 mm continuum images are shown in Fig. 1. The Band 3 spectrum
from the spectral window covering H recombination lines and all the Band 6
spectra (first presented and analyzed in Sewiło et al. 2022) are shown in
Appendix A. The mm-RL spectra for 1 A (both Band 6 and Band 3) were extracted
for the analysis as the mean within the contour corresponding to 50% of the
1.2 mm continuum emission peak.

###  2.1 ALMA Archival Data

We retrieved the Band 7 12CO (3–2) and HCO+ (4–3) data from the Cycle 7
project 2019.1.01770.S from the ALMA archive to explore, respectively, the
diffuse and dense gas surrounding N 105–1 A. Target Lh06 corresponding to our
ALMA field N 105–1, was observed as a part of the “MAGellanic Outflow and
chemistry Survey” (MAGOS; PI K. Tanaka). A detailed description of the
observations can be found in Tanaka et al. (2023, in prep.). The CO and HCO+
data were calibrated and imaged using CASA 6.1.2.7. We used the CASA task
tclean for imaging with the multi-scale deconvolver and the Briggs weighting
with a robust parameter of 0.5.

The CO (3–2) line with the rest frequency of 345,795.99 MHz and $E_{\rm U}$ of
11.5 K was included in the 937.5 MHz-wide spectral window centered at 345.3
GHz and divided into 1920 channels of 488.28 kHz each (or 0.42 km s-1). The
resulting CO (3–2) spectral cube has a $0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}41\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}33$ beam and sensitivity of 4.5 mJy beam-1
(or 0.35 K) at a channel width of 1 km s-1.

The HCO+ (4–3) line with the rest frequency of 356,734.22 MHz and $E_{\rm U}$
of 42.8 K was included in the 1.875 GHz-wide spectral window centered at 357.1
GHz and divided into 3840 channels of 488.32 kHz each (or 0.41 km s-1). The
resulting HCO+ (4–3) spectral cube has a $0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}40\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}32$ beam and sensitivity of 7.0 mJy beam-1
(or 0.52 K) at a channel width of 0.5 km s-1.

The 345.798 GHz ($\sim$870 $\mu$m) continuum image was constructed from line-
free channels in the Band 7 data; the sensitivity of 0.32 mJy per $0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}47\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}$39 beam was achieved in the continuum.

##  3 Previous Studies on N 105–1 A

The N 105 star-forming region is located at the western edge of the LMC bar
(e.g., Ambrocio-Cruz et al. 1998) and is associated with the OB association LH
31 (e.g., Lucke & Hodge 1970; 18 OB stars and two Wolf-Rayet stars) and the
sparse cluster NGC 1858 (e.g., Bica et al. 1996; age estimates in a range 8–17
Myr; Vallenari et al. 1994; Alcaino & Liller 1986). The ALMA observations
discussed in this paper are coincident with the brightest part of the optical
nebula (N 105A) and the peak of the molecular cloud in N 105 (traced by 12CO
1–0, e.g., Wong et al. 2011), dense gas emission peaks (traced by HCO+ and HCN
1–0; Seale et al. 2012), and the thermal radio continuum source MC 23 or
B0510–6857 (e.g., McGee et al. 1972, Ellingsen et al. 1994, Filipovic et al.
1998). The source with the detection of higher-order mm-RLs, N 105–1 A, is the
brightest 8.6 GHz (3 cm) and 4.8 GHz (6 cm) source in N 105 (B0510$-$6857 W)
detected by Indebetouw et al. (2004) with the Australia Telescope Compact
Array (ATCA; a resolution of $\sim$1$\hbox to0.0pt{.\hss}{{}^{\prime\prime}}$5
and $\sim$2′′ at 3 cm and 6 cm, respectively) and it was classified as an
ultracompact (UC) H ii region.

The infrared source coincident with N 105–1 A / B0510$-$6857 W was reported in
the literature as a candidate protostar by Epchtein et al. (1984) based on
five-band near-infrared (near-IR) photometry. The YSO classification was later
supported by Oliveira et al. (2006, their source N 105A IRS1) with high
spatial resolution near-IR spectroscopic observations. The 3–4 $\mu$m spectrum
from the Infrared Spectrometer And Array Camera (ISAAC) on the ESO-VLT
displays a very red continuum, very strong Br$\alpha$ and Pf$\gamma$
recombination line emission, and a non-detection of the Pf$\delta$ line which
was postulated to be the result of a high dust column density in front of the
source ($A_{\rm V}$$\sim$40 mag). Further evidence supporting the
interpretation of the source as an embedded protostar includes broad
Br$\alpha$ line wings likely indicating the presence of an outflow, and the
fact that the source is very bright in $L^{\prime}$-band and it is extremely
red ($K_{\rm S}$-$L^{\prime}$ = 3.9 mag). Strong hydrogen RLs indicate that
the massive YSO has started ionizing its immediate surroundings. The recent
VLT/KMOS near-IR spectroscopic observations reported in Sewiło et al. (2022)
also revealed strong hydrogen RLs - a full Brackett series emission in the
H$+$K bands.

N 105–1 A was also classified as a YSO based on the Spitzer Space Telescope
mid-IR photometric (Whitney et al. 2008, source #318 or SSTISAGE1C
J050950.53$-$685305.4; Gruendl & Chu 2009, source 050950.53$-$685305.5;
Carlson et al. 2012) and spectroscopic (Spitzer’s The InfraRed Spectrograph,
IRS 5–38 $\mu$m, Seale et al. 2012; Jones et al. 2017) observations. The
Spitzer/IRS spectrum exhibits relatively strong fine-structure lines such as
[S IV] 10.5 $\mu$m, [Ne II] 12.8 $\mu$m, [Ne III] 15.5 $\mu$m, [S III] 18.7
$\mu$m and 33.5 $\mu$m, and [S III] 34.8 $\mu$m, and weak emission lines from
the Polycyclic Aromatic Hydrocarbons (PAHs). The spectral energy distribution
(SED, from near- to mid-IR) of N 105–1 A was well-fit with the Robitaille et
al. (2006) YSO radiative transfer models by Carlson et al. (2012) with the
best-fit stellar mass and luminosity of $31.3\pm 2.6$ M⊙ and ($1.4\pm
0.2)\times 10^{5}$ L⊙, respectively.

N 105–1 A was included in the molecular spectral line analysis of the 1.2 mm
continuum sources in the ALMA field N 105–1 performed by Sewiło et al. (2022).
The CH3OH, SO2, SO, CS, H13CO+, H13CN, and (tentatively) H2CS lines are
detected toward 1 A (see Appendix A for more details). The detection of
multiple CH3OH and SO2 lines allowed for an independent rotational temperature
determination for each of these species.

The distributions of the H$\alpha$, CS, SO, H36$\beta$, 1.2 mm, and 6 cm
emission are compared in Fig. 2. The high-resolution H$\alpha$ image shows
that 1 A is associated with the extended low-intensity H$\alpha$ emission
around the mm/radio continuum peak, while the mm continuum emission extending
to the north from the main peak overlaps with what appears to be the
H$\alpha$-dark region. The CS emission has a filamentary structure and extends
even further to the north ($\sim$8′′ or 1.9 pc), filling in the H$\alpha$-dark
region. The H36$\beta$ mm-RL emission (and the emission from all other mm-RLs
except H57$\gamma$; see below) coincides with the mm/radio continuum peak,
while the CS and SO peaks are offset to the north. No molecular line emission
for species detected by Sewiło et al. (2022) coincides with the mm/radio
continuum peak.

N 105–1 A is associated with cold CH3OH ($12\pm 1$ K) and hot SO2 ($96\pm 20$
K). The presence of the hot SO2 component led Sewiło et al. (2022) to suggest
that 1 A may host a hot molecular core. Hot cores are compact ($D\lesssim 0.1$
pc), hot ($T_{\rm kin}\gtrsim 100$ K), and dense ($n_{\rm H}\gtrsim 10^{6-7}$
cm-3) regions around forming massive stars (e.g., Garay & Lizano 1999; Kurtz
et al. 2000; Cesaroni 2005; Palau et al. 2011). They produce rich molecular
line spectra with many lines from complex organic molecules (COMs; containing
6 or more atoms including C, Herbst & van Dishoeck 2009). COMs are the
products of interstellar grain-surface chemistry (released to the gas phase by
thermal evaporation or shock sputtering), or post-desorption gas chemistry
(e.g., Herbst & van Dishoeck 2009; Oberg 2016; Öberg & Bergin 2021). There are
only a handful of known bona fide hot cores outside the Galaxy, all are
located in the LMC (Shimonishi et al. 2016; Sewiło et al. 2018; Sewiło et al.
2019; Shimonishi et al. 2020; Sewiło et al. 2022).

Sewiło et al. (2022) argue that cold CH3OH, but hot SO2 detected toward the
1.2 mm continuum source in 1 A may indicate that the SO2 emission originates
from the area offset from the continuum source and/or CH3OH is sub-thermally
excited. Sewiło et al. (2022) conclude that it is more likely that the hot
core in 1 A is coincident with the SO2 peak offset from the continuum peak by
$\sim$0$\hbox to0.0pt{.\hss}{{}^{\prime\prime}}$6 where CH3OH is 10 K warmer
and that it is externally illuminated since no IR source is detected in the
existing observations; it was suggested that the hot core might be associated
with outflow shocks.

Figure 3: A comparison between the predicted and observed peak intensity
ratios of hydrogen mm-RLs (the peak intensity of a given mm-RL with respect to
the peak intensity of the H36$\beta$ transition) detected in Band 6 toward N
105–1 A (see Section 4.1 for details). The uncertainties of the observed peak
intensity ratios shown in the plot were calculated by error propagation using
the 1$\sigma$ uncertainties of the mm-RL peak intensities provided in Table 1.

##  4 Results and Analysis

###  4.1 Hydrogen Recombination Line Emission Toward N 105–1 A

We detected five hydrogen mm-RLs toward N 105–1 A in our Band 6 observations:
H36$\beta$ (260.03278 GHz), H41$\gamma$ (257.63549 GHz), H49$\epsilon$
(241.86116 GHz), H53$\eta$ (257.19399 GHz), and H54$\eta$ (243.94239 GHz). An
additional transition, H55$\theta$ (258.52592 GHz), is tentatively detected.
These are all hydrogen RLs with $\Delta n\leq 8$ (up to $\theta$-transitions)
falling within the frequency range of our observations (e.g., Gordon &
Sorochenko 2002). Three more mm-RLs (H40$\alpha$, H50$\beta$, and H57$\gamma$
at 99.02295 GHz, 99.22521 GHz, and 98.67189 GHz, respectively) were detected
toward N 105–1 A in our unrelated, lower-resolution ALMA Band 3 project
targeting high-mass YSOs in the LMC (see Section 2).

Figure 4: From left to right, top to bottom: Band 6 H36$\beta$, H41$\gamma$,
H49$\epsilon$, H53$\eta$, and H53$\eta$ recombination line integrated
intensity images of N 105–1. The white contours in each image correspond to
the 1.2 mm continuum emission with contour levels of (3, 10, 100)$\sigma$.

Figure 5: From left to right: Band 3 H40$\alpha$, H50$\beta$, and H57$\gamma$
recombination line integrated intensity images of N 105–1 A. The white
contours are the same as in Fig. 4.

We verified the identification of the faintest reliably-detected mm-RLs,
H53$\eta$ and H54$\eta$, by comparing their observed peak intensities relative
to the peak intensity of the H36$\beta$ line (the brightest RL in our Band 6
observations), to those predicted by theory, assuming local thermodynamic
equilibrium (LTE) and optically thin conditions (the assumptions supported by
our subsequent analysis; see below and Sections 4.2–4.3). We used theoretical
intensities from Table B2 in Gordon & Sorochenko (2002), which reproduces
entries in Table 1 of Towle et al. (1996). The relative mm-RL intensities are
plotted in Fig. 3. There is a good agreement for H41$\gamma$ and H49$\epsilon$
lines. For the H53$\eta$ and H54$\eta$ lines, the observed relative
intensities deviate from those predicted by the theory, but this deviation is
small enough to be explained by the uncertainties in their analysis
($\sim$6$\sigma$-detections). The spatial distribution of the H53$\eta$ and
H54$\eta$ line emission is in good agreement with other mm-RLs (see Figs. 4
and 5) and the lines are too broad for molecular lines (Fig. 6).

To our knowledge, only the $\alpha$\- and $\beta$-transitions have been
observed outside the Galaxy to date, making the $\gamma$-, $\epsilon$-, and
$\eta$-transitions observed toward N 105–1 A the first extragalactic
detections. The $\eta$-transitions are the first to be detected at millimeter
wavelengths.

The H36$\beta$, H41$\gamma$, H49$\epsilon$, H53$\eta$, and H54$\eta$
integrated intensity images are shown in Fig. 4 ($\sim$0$\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}$5 resolution), and the H40$\alpha$,
H50$\beta$, and H57$\gamma$ integrated intensity images are shown in Fig. 5
($\sim$2$\hbox to0.0pt{.\hss}{{}^{\prime\prime}}$2). The mm-RL emission peaks
for all observed transitions except H57$\gamma$ (detected in the 3 mm band)
are coincident with the millimeter and radio continuum peaks, but all are
offset from the CS, SO, and SO2 molecular line emission peaks (by $\sim$1–1.4
times the beam size at 1.2 mm; see Fig. 2 and Figs. 4–5). The H57$\gamma$
emission peak appears to be offset to the east from the continuum peak by
$\sim$$0\hbox to0.0pt{.\hss}{{}^{\prime\prime}}6$; since this offset
corresponds to a relatively small fraction of the 3 mm beam ($\sim$30%) and
the signal-to-noise ratio for this line is low (lower than for any other mm-RL
detected toward 1 A), we do not attempt to assign any physical meaning to it.

We fitted Gauss functions to all mm-RLs using the CASSIS interactive spectrum
analyzer (Vastel et al. 2015). Table 1 shows the Gaussian fitting results: the
peak velocity ($v_{LSR}$), linewidth ($\Delta v_{\rm FWHM}$, the full-width at
half maximum), peak intensity ($T_{b}$), and the integrated intensity ($\int
T_{b}dv$) for each transition. No correction for beam dilution was applied.
The observed mm-RLs with the fitted Gauss profiles are shown in Fig. 6.

Figure 6: Left: Hydrogen recombination lines observed toward N 105–1 A in Band
6 (from top to bottom in order of decreasing frequency): H36$\beta$,
H41$\gamma$, H53$\eta$, H54$\eta$, and H49$\epsilon$. Right: Hydrogen
recombination lines observed toward N 105–1 A in Band 3 (from top to bottom):
H40$\alpha$, H50$\beta$, and H57$\gamma$. The observed spectra are shown in
light gray; the spectra shown in color were smoothed using the Hanning window
with the window size of 11 (Band 6) and 5 (Band 3; using
PyAstronomy.pyasl.smooth). The dashed and dotted lines in the left panel
indicate the central velocity of the H36$\beta$ (231.9 km s-1) and H13CO+
(237.2 km s-1) lines, respectively, while the dashed line in the right panel
indicates the central velocity of the H40$\alpha$ line (230.1 km s-1). The
fitted Gauss functions are overlaid. The results of the fitting are reported
in Table 1.

A comparison of the observed intensity ratios of H$n\alpha$ and higher-order
lines observed at approximately the same frequency to theoretical predictions
provides an empirical test for departures from LTE. For a pair of lines with
similar frequencies the size of the beam will be similar and thus
complications due to the inhomogeneity in the nebula will be significantly
reduced or eliminated. The H40$\alpha$ and H50$\beta$ lines detected toward 1
A are an ideal line pair to investigate the conditions under which the
observed recombination lines are formed. The theoretical value of the
H50$\beta$/H40$\alpha$ line ratio is 0.275 (Menzel 1969; Dupree & Goldberg
1970) under the assumption that the impact broadening can be neglected and
both lines are optically thin and formed in LTE. The observed
H50$\beta$/H40$\alpha$ line ratio toward 1 A of 0.29 $\pm$ 0.06 is in good
agreement with the theoretical value, indicating that the region is likely to
be in LTE.

Table 1: Gauss Fitting Results for Hydrogen Recombination Lines Detected Toward N 105–1 Aa,ba,bfootnotemark:  Transition | Frequency | $v_{LSR}$ |  $\Delta v$ccThe full width at half-maximum (FWHM) corrected for instrumental broadening: $\Delta v=\sqrt{\Delta v_{\rm obs}^{2}-\Delta v_{\rm instr}^{2}}$, where $\Delta v_{\rm obs}$ is the observed linewidth and $\Delta v_{\rm instr}$ is the channel width after Hanning smoothing.  | $T_{b}$ | $\int T_{b}dv$
---|---|---|---|---|---
| (GHz) | (km s-1) | (km s-1) | (K) | (K km s-1)
Band 6
H36$\beta$ | 260.03278 | 231.9(0.6) | 32.3(1.3) | 0.84(0.03) | 29.0(1.5)
H41$\gamma$ | 257.63549 | 231.1(0.8) | 33.5(1.8) | 0.39(0.02) | 14.2(1.0)
H49$\epsilon$ | 241.86116 | 227.7(2.7) | 35.7(6.4) | 0.14(0.02) | 5.5(1.3)
H53$\eta$ | 257.19399 | 239.3(1.9) | 25.7(4.5) | 0.12(0.02) | 3.3(0.7)
H54$\eta$ | 243.94239 | 233.9(1.4) | 24.4(3.4) | 0.12(0.02) | 3.3(0.6)
Band 3
H40$\alpha$ | 99.02295 | 230.1(0.6) | 31.1(1.5) | 0.91(0.04) | 30.6(1.9)
H50$\beta$ | 99.22521 | 229.9(2.9) | 29.3(6.9) | 0.26(0.05) | 8.4(2.6)
H57$\gamma$ | 98.67189 | 228.6(5.7) | 25.5(13.5) | 0.13(0.06) | 3.6(2.4)

###  4.2 XCLASS Modeling of Recombination Lines

We model the observed mm-RLs using the eXtended CASA Line Analysis Software
Suite (XCLASS111https://xclass.astro.uni-koeln.de/; Möller et al. 2017) with
additional extensions (T. Möller, priv. comm.). In XCLASS, a contribution of
each mm-RL to a model spectrum is described by a certain number of components,
with each component described by the source size ($\theta_{\rm source}$),
electron temperature ($T_{\rm e}$), emission measure (EM), line width ($\Delta
v$), velocity offset ($v_{\rm offset}$), and the position along the line of
sight. The physics behind the model is summarized in Appendix B.

All the model parameters are fitted to the observational data using the
Levenberg-Marquardt algorithm (Marquardt 1963) provided by the optimization
package MAGIX222https://magix.astro.uni-koeln.de/ (Möller et al. 2013). The
errors of the mm-RL parameters are estimated using the
emcee333https://emcee.readthedocs.io/en/stable/ package (Foreman-Mackey et al.
2013), which implements the affine-invariant ensemble sampler of Goodman &
Weare (2010) to perform a Markov chain Monte Carlo (MCMC) algorithm,
approximating the posterior distribution of the model parameters by random
sampling in a probabilistic space. Details of the error estimation are
described in Möller et al. (2021). To get a more reliable error estimation,
the errors for the emission measures and the electron densities are calculated
on a logarithmic scale, i.e., these parameters are converted to their log10
values before applying the MCMC algorithm and converted back to linear scale
after finishing the error estimation procedure.

In our XCLASS model, we include all mm-RL transitions up to
$\theta$-transitions ($\Delta n=8$) covered by our ALMA observations in both
Band 3 and Band 6: H40$\alpha$, H36$\beta$, H50$\beta$, H41$\gamma$,
H57$\gamma$, H49$\epsilon$, H67$\epsilon$, H53$\eta$, H54$\eta$, H74$\eta$,
H55$\theta$, and H77$\theta$; spectra of N 105–1 A centered on the rest
frequencies of these mm-RLs are shown in Fig. B.1. We also include several
molecular species detected toward N 105–1 A (CH3OH, H13CO+, H13CN, CS, H2CS,
SO2, and SO) to identify any contributions from the molecular lines to the mm-
RL profiles; the LTE-parameters for these molecules are taken from Sewiło et
al. (2022). Recombination lines of atoms other than H (such as He, C, N, O,
and S) are not reliably detected toward N 105–1 A.

We perform the XCLASS analysis for three sets of lines: all 12 mm-RL
transitions from Band 3 and Band 6, Band 3 transitions only, and Band 6
transitions only. For each set, we examine four different scenarios. In the
first scenario, we assume LTE conditions and describe all mm-RLs by a single
component, i.e., we use a single set of $T_{\rm e}$ and EM to model all mm-RL
transitions. In the second scenario, we again assume that all mm-RLs are in
LTE, but now we use an additional (second) component. In the remaining two
scenarios, we model all mm-RL transitions in non-LTE, using one or two
components.

We are unable to obtain a good model fit (i.e., reliably determine $T_{\rm e}$
and EM) for any combination of the line sets and scenarios. For line sets with
single band mm-RL transitions, this is likely because many lines have a very
low signal-to-noise ratio. For the line set combining all Band 3 and Band 6
mm-RL transitions, the difference in the angular resolution between the Band 3
and Band 6 observations (and the resulting beam dilution in Band 3) prevents
us from obtaining a satisfactory fit to the data. We did not find any
significant differences between the LTE and non-LTE model fits, indicating
that LTE is a reasonable assumption.

In Fig. B.1 in Appendix B, we present an example XCLASS synthetic spectrum
fitted to Band 6 mm-RL transitions, representing the LTE model with one
component. The corner plot of the MCMC error estimate for the same model is
shown in Fig. B.2; it indicates that it is not possible to constrain $T_{\rm
e}$ and EM based on this model fit. Figure B.1 still provides us with some
useful information. Firstly, it demonstrates that we detected (or marginally
detected in the case of the H55$\theta$ line) all hydrogen mm-RLs with $\Delta
n\leq 8$ covered by our Band 6 observations. Secondly, Fig. B.1 shows that the
observed mm-RL intensities in Band 6 are consistent with the LTE model.
Lastly, extending the model to lower frequencies reveals that the intensities
of mm-RLs detected in Band 3 are significantly lower than predicted by the
model, indicating that beam dilution effects are considerable in the Band 3
observations having about five times larger beam than in Band 6.

###  4.3 Physical Properties of the Ionized Gas in N 105–1 A Based on the
H40$\alpha$ and 3 mm Continuum Data

We determine the physical properties of the ionized gas in N 105–1 A using the
standard analytical methods incorporating a continuum free-free and/or an
H$n\alpha$ recombination line emission. The H40$\alpha$ line is the only
$\alpha$-transition currently available for 1 A.

####  4.3.1 The Electron Temperature

Assuming the recombination lines are emitted under pure LTE conditions and
pressure broadening of the lines by electron impacts is negligible, the
observed recombination line to continuum ratio can be utilized to directly
estimate the LTE electron temperature ($T_{\rm e}^{*}$) of the ionized gas
(e.g., Dupree & Goldberg 1970 and references therein; Roelfsema & Goss 1992;
Gordon & Sorochenko 2002) from the following formula for H$n\alpha$ lines
(Wilson et al. 2009):

$\displaystyle\frac{I_{\rm L}}{I_{\rm C}}\left(\frac{\Delta v}{{\rm
km\,s^{-1}}}\right)\approx\frac{6985}{a(T_{\rm e},\nu)}\left(\frac{\nu_{\rm
L}}{{\rm GHz}}\right)^{1.1}$ (1) $\displaystyle\cdot\left(\frac{T_{\rm
e}^{*}}{{\rm K}}\right)^{-1.15}\left(\frac{1}{1+y^{+}}\right)$

where $I_{\rm L}$ is the line intensity, $I_{\rm C}$ is the free-free
continuum intensity, $\Delta v$ is the FWHM line width, $\nu_{\rm L}$ is the
H$n\alpha$ line rest frequency, $y^{\rm+}$ is a singly ionized helium-to-
hydrogen abundance ratio ($N_{\rm He^{+}}/N_{\rm H^{+}}$), $a(T_{e},\nu)$ is a
dimensionless factor of the order of 1 slowly varying with $T_{e}$ and $\nu$,
tabulated in Mezger & Henderson (1967). At the frequency of $\sim$100 GHz,
$a$$\sim$0.9 for $T_{e}$ from 8,000 to 12,000 K, typical for H ii regions.
$I_{\rm L}/I_{\rm C}$ ratio allows $I_{\rm L}$ and $I_{\rm C}$ to be given in
any convenient units.

We utilize the H40$\alpha$ line ($\nu_{\rm L}$ = 99.02295 GHz) to estimate
$T_{\rm e}^{*}$ in N 105–1 A. The $y^{\rm+}$ parameter cannot be determined
for 1 A in this work since no He recombination line has been detected toward
this source in our observations. A typical value of $y^{\rm+}$ in the Galaxy
is $\sim$8–10%, although significant variations have been found (e.g.,
Churchwell et al. 1974; Lockman & Brown 1975; Balser et al. 2001). Previous
studies suggest that $y^{\rm+}$ in the LMC is similar to that measured in the
Galaxy (e.g., Dufour 1975; Rosa & Mathis 1987; Peck et al. 1997; Tsamis et al.
2003). Dufour (1975) determined $y^{\rm+}$ of 0.1 specifically for N 105 using
optical spectroscopy.

We obtain $I_{\rm L}\,\Delta v$ from the H40$\alpha$ integrated intensity,
$\int I_{\rm L}dv$; for Gaussian line profiles, $\int I_{\rm L}dv\approx
1.064\,I_{\rm L}\,\Delta v$ (e.g., Brown et al. 1978). We determine $\int
I_{\rm L}dv$ of $1110\pm 10$ mJy km s-1 from the H40$\alpha$ integrated
intensity (moment 0) map by measuring the emission within the 3$\sigma$
contour. We measured the 99 GHz continuum flux density of $43.2\pm 0.4$ mJy
within the same area from the 99 GHz continuum map. The 3 mm emission is
expected to be mostly free-free, but some contribution from the dust thermal
emission is likely. We utilized the archival 870 $\mu$m continuum image to
estimate the dust emission at 3 mm under the assumption that all the emission
at 870 $\mu$m originates from dust. We measured the 870 $\mu$m continuum flux
density (Band 7, B7) on the image smoothed to the beam of the 3 mm image and
scaled it to the expected 3 mm continuum dust-only flux density (B3) using the
formula: $S_{\rm B3}=S_{\rm B7}\,(\nu_{\rm B3}/\nu_{\rm B7})^{2+\beta}$, where
$\beta$ is a dust emissivity index. We adopt $\beta$ of 1.7 for N 105 (Gordon
et al. 2014; see also Sewiło et al. 2022). The comparison of the predicted
dust-only and total flux density measured from the same region at 3 mm
indicates that only 3.6% of the 3 mm emission is the dust thermal emission. To
calculate $T_{\rm e}^{*}$, we removed the contribution from the dust emission
from the 3 mm (99 GHz) flux density.

From Eq. 1, we determine $T_{\rm e}^{*}$ of $10\,940\pm 1\,350$ K toward 1 A
which is consistent with $T_{\rm e}$ measurements toward Galactic H ii regions
(see Section 5.1); the uncertainties were calculated by error propagation
after incorporating the 10% flux calibration error into the flux density
uncertainties. Since the observed H50$\beta$/H40$\alpha$ line ratio is close
to its theoretical LTE value in 1 A, $T_{\rm e}^{*}$ is a good representation
of the true electron temperature ($T_{\rm e}$).

####  4.3.2 Emission Measure and Electron Density

We calculate the emission measure (EM) toward N 105–1 A using the following
formula (e.g., Wilson et al. 2009):

$\displaystyle{\rm EM}\,({\rm pc\,cm^{-6}})=1.74\times
10^{-3}\left(\frac{T_{\rm e}}{{\rm K}}\right)^{3/2}\left(\frac{\nu_{\rm
L}}{{\rm GHz}}\right)$ (2) $\displaystyle\cdot\,\eta^{-1}\left(\frac{T_{\rm
L}}{{\rm K}}\right)\left(\frac{\Delta v}{{\rm km}\,{\rm s}^{-1}}\right),$

where $T_{\rm L}$ is the H$n\alpha$ line intensity in K, $\eta$ is the beam
filling factor, and other parameters are the same as in Eq. 1. We use $T_{\rm
e}=10\,940$ K (Section 4.3.1). We calculate $T_{\rm L}\,\Delta v$ from the
integrated intensity of the H40$\alpha$ line ($\int I_{\rm L}dv$ converted to
K km s-1) as described above. The beam filling factor
$\eta=\theta_{s}^{2}/(\theta_{s}^{2}+\theta_{\rm beam}^{2})$, where
$\theta_{s}^{2}$ and $\theta_{\rm beam}^{2}$ are the FWHM size of the source
and the synthesized beam, respectively. For 1 A, we obtain $\theta_{s}$ of
$0\hbox to0.0pt{.\hss}{{}^{\prime\prime}}55\pm 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}10$ based on the 2D Gauss fitting to the 3 mm
image with the CASA task imfit; it is the geometric mean of the major and
minor source FWHMs deconvolved from the beam. We obtain EM of $(8.9\pm
1.7)\times 10^{7}$ pc cm-6 toward N 105–1 A.

The emission measure can also be estimated based on the free-free continuum
emission (EMcont); it is expressed by the formula (e.g., Mezger & Henderson
1967; Roelfsema & Goss 1992):

$\displaystyle{\rm EM_{cont}}\,({\rm pc\,cm^{-6}})=\frac{\tau_{\rm
C}}{8.235\times 10^{-2}\,a(T_{\rm e},\nu)\,T_{\rm e}^{-1.35}\,\nu^{-2.1}},$
(3)

where $\tau_{\rm C}$ is the continuum optical depth, $\nu$ is the continuum
frequency in GHz, and $a(T_{\rm e},\nu)$ is the same as in Eq. 1. We calculate
the peak $\tau_{\rm C}$ and EMcont based on the 3 mm continuum data.

For the 99.023 GHz / 3 mm free-free continuum peak intensity of $43.64\pm
2.10$ mJy per $2\hbox to0.0pt{.\hss}{{}^{\prime\prime}}49\times 1\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}83$ beam (96.4% of the observed continuum
peak of 45.27 mJy beam-1), we obtain the brightness temperature ($T_{\rm b}$)
of $1.19\pm 0.06$ K, assuming the Rayleigh-Jeans approximation. The
uncertainties include the continuum image rms and a 10% flux density
calibration error. We estimate the peak continuum optical depth ($\tau_{\rm
C}$) from the formula $T_{\rm b}=T_{\rm e}(1-e^{-\tau_{\rm C}})$ after
correcting $T_{\rm b}$ for the beam dilution effect: $\tau_{\rm C}=-{\rm
ln}(1-\eta^{-1}\,T_{\rm b}/T_{\rm e})=(1.7\pm 0.2)\times 10^{-3}$. From Eq. 3,
we obtain the peak EMcont of $(1.0\pm 0.2)\times 10^{8}$ pc cm-6 toward N
105–1 A, in a very good agreement with the EM derived from the mm-RL data.

To estimate the rms electron density, $n_{\rm e}$, we use the expression
$n_{\rm e}=\sqrt{{\rm EM}/\Delta s}$, where $\Delta s$ (in pc) is the size of
the source along the line of sight. Assuming that the extent of the source
along the line of sight is the same as in the plane of the sky, we adopt the
source size obtained from the 2D Gauss fitting ($\theta_{s}$) as $\Delta s$
($0.13\pm 0.02$ pc). The resulting $n_{\rm e}$ is $(2.6\pm 0.3)\times 10^{4}$
cm-3 for N 105–1 A.

The values of (FWHM size, EM, $n_{\rm e}$) of (0.13 pc, $\sim$$9\times 10^{7}$
pc cm-6, $\sim$$3\times 10^{4}$ cm-3) derived for N 105–1 A indicate that this
LMC source is likely an UC H ii region. UC H ii regions are one of the
earliest phases of massive star formation, defined observationally as regions
with sizes $\lesssim$0.1 pc, $n_{\rm e}\gtrsim 10^{4}$ cm-3, and ${\rm
EM}\gtrsim 10^{7}$ pc cm-6 (e.g., Wood & Churchwell 1989; Kurtz et al. 2000;
Hoare et al. 2007). The UC H ii region phase follows the hypercompact (HC) H
ii region phase (regions ten times smaller and 100 times denser than UC H iis,
characterized by positive spectral indices and often showing broad RLs; e.g.,
Kurtz 2005; Hoare et al. 2007) and precedes the compact H ii region phase. The
formation of the H ii region signals the arrival of a massive protostar onto
the main sequence and thus is the earliest manifestation of an OB star. The
massive stars reach the main sequence while still deeply embedded in their
natal molecular cloud and with ongoing accretion (e.g., Churchwell 2002;
Zinnecker & Yorke 2007; Tanaka et al. 2016). The immediate surroundings of the
ionizing star(s) of HC/UC H ii regions are therefore expected to be very
dynamic due to possible infall, outflows, stellar winds, accretion disk
rotation, turbulence, and shocks.

####  4.3.3 Ionizing Flux of the Exciting Star

For a spherical, homogeneous, optically thin, and dust-free model, the number
of Lyman continuum photons ($N_{\rm Ly}$) required to ionize the source can be
expressed by the formula (e.g., Goudis 1977 and references therein):

$\displaystyle N_{\rm Ly}(s^{-1})=4.76\times
10^{48}\left(\frac{S_{\rm\nu}^{\rm ff}}{{\rm Jy}}\right)\left(\frac{D}{{\rm
kpc}}\right)^{2}$ (4) $\displaystyle\cdot\left(\frac{\nu}{{\rm
GHz}}\right)^{0.1}\left(\frac{T_{\rm e}}{{\rm K}}\right)^{-0.45},$

where $S_{\rm\nu}^{\rm ff}$ is the free-free continuum integrated flux density
at a frequency $\nu$, and $D$ is a distance. For $S_{\rm\nu}^{\rm ff}$, we
adopt the 3 mm integrated flux density obtained using the 2D Gauss fitting
with the CASA task `imfit`: $48.5\pm 0.8$ mJy (in an excellent agreement with
the flux density integrated over the region enclosed by the 3$\sigma$ contour,
see Table C), after removing the contribution from the thermal dust emission
(3.6%). We adopt $T_{\rm e}$ determined in Section 4.3.1 and the LMC distance
of 49.6 kpc (Pietrzyński et al. 2019)

For 1 A, Eq. 2 results in $N_{\rm Ly}$ of $1.32\times 10^{49}$ photons per
second (log $N_{\rm Ly}=49.12$) which corresponds to an O5.5 V star (log
$N_{\rm Ly}$=49.11) according to models of Galactic O stars by Martins et al.
(2005). The model luminosity and spectral mass for an O5.5 V star are
$2.6\times 10^{5}$ $L_{\odot}$ and 34.4 $M_{\odot}$, respectively.

In an earlier work by Smith et al. (2002), the authors modeled ionizing flux
densities for the metallicity range from 0.05 to 2 $Z_{\odot}$. In their
models, the ionizing flux of log $N_{\rm Ly}$=49.12 corresponds to the O6 V
spectral type and luminosity of $\sim$$3\times 10^{5}$ $L_{\odot}$ for the
metallicity within the range observed in the LMC ($Z=0.4\,Z_{\odot}$). Smith
et al. (2002) predict the same ionizing fluxes for the solar metallicity for
the relevant spectral type range for dwarfs.

###  4.4 SED Fitting: Physical Properties of the YSO Associated with N 105–1 A

Figure 7: The spectral energy distribution (SED) for N 105–1 A, covering a
wavelength range from 1.25 $\mu$m to 2.7 mm. Filled circles and triangles are
valid flux values and flux upper limits, respectively. The flux error bars are
smaller than the data points. The photometry and references are provided in
Table C.

The IR source coinciding with N 105–1 A was identified as a YSO candidate in
several previous studies (see Section 3). We fit the spectral energy
distribution (SED) of YSO 050950.53$-$685305.5 with a set of radiative
transfer model SEDs for YSOs developed by Robitaille (2017) using the
Robitaille et al. (2007) SED fitting tool. The SED fitting was previously
performed for the YSO associated with 1 A (Carlson et al. 2012), but without
the long-wavelength photometry ($>$20 $\mu$m) that is crucial to constrain the
evolutionary stage of the YSOs, and using the older version of the YSO models
(Robitaille et al. 2006).

To construct the multiwavelength SED of 1 A, we compiled the photometric data
covering a wavelength range from $\sim$1.3 $\mu$m to 3 mm. The SED is shown in
Fig. 7, while flux densities with references and technical details are
provided in Appendix C. Figure C.1 shows the images of N 105–1, from optical
(H$\alpha$) to radio (6 cm) wavelengths.

We do not include the Herschel photometry except the 100 $\mu$m flux density
in the fitting because at the longer Herschel wavelengths (at the lower
spatial resolution), the emission from 1 A is unresolved from the neighboring
source to the east (see Fig. C.1). In addition, we set the two longest-
wavelength data points out of 11 extracted from the Spitzer/IRS spectrum (see
Appendix C) as upper limits for the fitting due to the reduced Spitzer’s
spatial resolution at these wavelengths ($>$20 $\mu$m; see also Fig. C.1). We
also use the 870 $\mu$m and 1.2 mm (dust-only) flux densities as upper limits
because they include the emission from both the compact source and the
extended component (they were measured within the corresponding 3$\sigma$
contour, see Table C). The dust-only 3 mm flux density (3.6% of the total 3 mm
flux density; see Section 4.3.1) is used with the 20% uncertainty to account
for the flux calibration errors and the uncertainty in estimating the dust
contribution to the free-free emission at this wavelength.

Figure 8: The SED of 1 A showing the data points used for the fitting and with
the best-fit YSO SED model from the Robitaille (2017) grid overplotted. Filled
circles and triangles are valid flux values and flux upper limits,
respectively. The values of a reduced $\chi^{2}$, interstellar visual
extinction ($A_{\rm V}$), and a distance scale for the best-fit model are
indicated in the plots. The icon in the upper right corner of the plot
represents the basic layout of the best-fit YSO model (Robitaille 2017).

The Robitaille (2017) YSO SED models consist of a combination of several
components: a star, a disk, an infalling envelope, bipolar cavities, and an
ambient medium. They were computed as 18 sets of models with increasing
complexity, described by two up to 12 variables. To find the best-fitting YSO
model, Robitaille (2017) recommends to identify the best model set first using
a Bayesian approach and adopt the model with the lowest $\chi^{2}$ from the
most likely set (see also Sewiło et al. 2019). The most likely model set
contains the highest fraction of models that provide a “good” fit defined as
models with $\chi^{2}-\chi^{2}_{\rm best}<F\,n_{\rm data}$ where
$\chi^{2}_{\rm best}$ is $\chi^{2}$ of the best-fit model across all the model
sets and all models, $n_{\rm data}$ is the number of data points used for the
fitting, and $F$ is is a threshold parameter that is determined empirically.
In case of 1 A, $F=6$ is a reasonable choice; however, the well-fit model
statistics is very small. Ultimately, the best-fit model is the model with the
lowest $\chi^{2}$ across all the model sets and all models. The SED with the
best-fit model overlaid is shown in Fig. 8. The best-fit model comes from the
model set `spubsmi` and includes a central source, passive disk, Ulrich
(rotationally flattened) envelope, bipolar cavities, an ambient interstellar
medium (see an inset cartoon in Fig. 8, from Robitaille 2017). The presence of
both the envelope and the disk in the best-fit model indicates that
050950.53$-$685305.5 is likely a Class I YSO.

We determine the luminosity of the central source of $1.3\times 10^{5}$
$L_{\odot}$ from the stellar radius and effective temperature returned by the
fitter for the best-fit YSO model using the Stefan-Boltzmann equation. To
obtain mass, we compare the position of the best-fit model in the Hertzsprung-
Russell (H-R) diagram with the PARSEC evolutionary tracks that include the
pre-main sequence (PMS) stage; the evolutionary tracks were calculated for the
initial stellar masses from 0.1 to 350 $M_{\odot}$, and a range of
metallicities (Bressan et al. 2012; Chen et al. 2015). We adopt the mass of
the closest PARSEC track in the H-R diagram as the mass of YSO
050950.53$-$685305.5. The age of the PARSEC photospheric model is interpolated
from the closest point on the track to the model on the H-R diagram to
constrain the age of the source. The resulting mass and age of the YSO are 28
$M_{\odot}$ and $1.1\times 10^{5}$ years, respectively. The luminosity and
mass of 1 A obtained in our analysis are in good agreement with the SED
fitting results of Carlson et al. (2012, see Section 3).

The luminosity based on the SED fitting is two times lower and the mass is
$\sim$20% lower than the values calculated from the mm-RLs and the 3 mm
continuum. Despite these lower values, the SED results are generally
consistent with the mm-RL and 3 mm values. We consider the recombination line
values for luminosity and mass to be more reliable, as they are based on
direct observations of the ionizing photons.

Figure 9: The H40$\alpha$ (upper panel) and H36$\beta$ (lower panel)
integrated line intensity (moment 0, left), intensity weighted velocity
(moment 1, center), and observed line width (right) images with the 1.2 mm
continuum contours overlaid for reference, with contour levels the same as in
Fig. 1. To make moment maps, we only used the emission above 3$\sigma$. The
velocity gradient of 17 km s-1 over 3$\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}$5 or 20 km s-1 pc-1 is detected in
H40$\alpha$ in a roughly SE–NW direction. The size of the ALMA synthesized
beam ($2\hbox to0.0pt{.\hss}{{}^{\prime\prime}}56\times 1\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}95$ for H40$\alpha$ and $0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}51\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}47$ for H36$\beta$) is shown in the lower
left corner in each image.

###  4.5 Ionized Gas Kinematics

There are three groups of the line broadening mechanisms that can affect the
recombination line widths: the natural, Doppler, and pressure (collisional)
broadening (e.g., Gordon & Sorochenko 2002). The effects of the natural and
collisional broadening are negligible for mm/submm RLs (e.g., Brocklehurst &
Seaton 1972; Roelfsema & Goss 1992; Gordon & Sorochenko 2002), thus we only
consider the Doppler thermal and non-thermal contributions to the line widths
of mm-RLs detected toward N 105–1 A.

The thermal line width can be determined from the formula:

$\displaystyle\Delta v_{\rm th}=\sqrt{8\,{\rm ln}2\,\frac{k_{\rm B}T_{\rm
e}}{\mu m_{\rm H}}}$ (5)

where $k_{B}$ is the Boltzmann constant, $T_{\rm e}$ is the gas electron
temperature, $\mu$ is the molecular weight in atomic units (for atomic
hydrogen, $\mu_{\rm H}$ = 1.008 amu), and $m_{\rm H}$ is the mass of the
hydrogen atom. For the temperature of the hydrogen gas of $10\,940\pm 1\,350$
K, $\Delta v_{\rm th}$ is $22.3\pm 1.4$ km s-1.

The non-thermal (turbulence and large-scale motions) line width can be
estimated using the equation:

$\displaystyle\Delta v_{\rm nth}=\sqrt{\Delta v^{2}-\Delta v_{\rm th}^{2}}$
(6)

where $\Delta v$ is the measured line width corrected for instrumental
broadening (see Table 1). The average line width for the three brightest
recombination lines is 32.3 km s-1, resulting in $\Delta v_{\rm nth}$ of
$\sim$23.4 km s-1, similar to $\Delta v_{\rm th}$.

The relatively large non-thermal component toward 1 A likely indicates the
presence of bulk motions in the region. To demonstrate this, we constructed
the velocity and line width (FWHM) images for the two brightest recombination
lines detected toward 1 A: H40$\alpha$ (Band 3) and H36$\beta$ (Band 6;
$\sim$5$\times$ higher angular resolution than the Band 3 observations); all
the images are shown in Fig. 9. The H40$\alpha$ velocity map reveals the
velocity gradient of 17 km s-1 over 3$\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}$5 (or 20 km s-1 pc-1) along the line going
through the peak of the 1.2 mm continuum emission in a roughly SE–NW
direction. It is striking though that the velocity changes the most at the
outskirts of the source while it remains roughly the same across its central
part. The H40$\alpha$ lines are the broadest in the area southeast from the
continuum peak, possibly hinting on some additional line broadening caused by
external bulk motions. The results on the kinematic structure of the ionized
gas toward 1 A based on the current H40$\alpha$ data have to be treated with
caution since the signal-to-noise ratio of the H40$\alpha$ line is low in the
outermost regions. The H36$\beta$ emission traces much smaller scales close to
the central star; the overall trend in velocity seems to be similar to that
seen in H40$\alpha$. The H36$\beta$ lines are the broadest at the continuum
peak.

The velocity gradients observed in the ionized gas may trace infall, outflow,
rotation, or their combination, or highlight overlapping velocity components.
The observed velocity patterns are influenced by the viewing angle, often
making the interpretation difficult. Based on the near-IR spectroscopic data,
Oliveira et al. (2006) found evidence for the outflow in 1 A (their source N
105A IRS1), therefore the velocity gradient detected in mm-RLs in the same
direction as observed in the near-IR data may trace the outflow. However, we
did not find any clear outflow signatures in the 12CO (3–2) data (see Section
5.3.4). The H40$\alpha$ and H36$\beta$ velocity gradients are not homogenous,
therefore the rotation is not the most likely interpretation, but cannot be
ruled out.

##  5 Discussion

In this section, we compare the physical properties obtained for the LMC
source N 105–1 A to those of Galactic H ii regions to search for any
differences that could be explained by different environments in these
galaxies. We also investigate the distribution and kinematics of the ionized
gas in N 105–1 A in the context of its molecular environment to further
explore the nature of the source.

###  5.1 Electron Temperature of N 105–1 A vs. Galactic H ii Regions

The recombination line observations of Galactic star-forming regions revealed
the $T_{\rm e}$ Galactocentric radial gradient which was interpreted as the
metallicity gradient (e.g., Churchwell & Walmsley 1975; Churchwell et al.
1978; Shaver et al. 1983). The electron temperature increases while the
metallicity decreases with increasing Galactocentric distance ($R_{\rm GC}$,
e.g., Shaver et al. 1983; Balser et al. 2011; Fernández-Martín et al. 2017;
Maciel & Andrievsky 2019). The lower abundance of metal coolants (such as C
and O) affects the balance between the heating and cooling within H ii
regions, resulting in warmer temperatures. Shaver et al. (1983) performed the
RL observations (to obtain reliable $T_{\rm e}$) and optical spectra (to
derive metal abundances using $T_{\rm e}$ from RLs) of a large sample of
Galactic H ii regions to investigate $T_{\rm e}$ and metallicity variations
with $R_{\rm GC}$; they found that $T_{\rm e}=(433\pm 40)\,R_{\rm GC}+(3150\pm
110)$, where $T_{\rm e}$ is in K and $R_{\rm GC}$ in kpc.

Based on the Shaver et al. (1983)’s results, we find that Galactic H ii
regions with $T_{\rm e}$ of 10 940 K, the value we determined for N 105–1 A,
are expected to be at $R_{\rm GC}$ of $\sim$18 kpc and have the oxygen
abundance (${\rm 12+[O/H]}=-0.067\,R_{\rm GC}+9.38$) of $\sim$8.2 (or
metallicity $Z$ of $\sim$0.32 Z⊙). This is consistent with the oxygen
abundance ($[12+{\rm log(O/H)}]_{\rm LMC}=8.4$) and metallicity ($Z_{\rm LMC}$
= 0.3–0.5 $Z_{\odot}$) measured in the LMC (e.g., Russell & Dopita 1992). The
oxygen abundance variations determined in other studies provide similar
results, e.g., $R_{\rm GC}$ at which the oxygen abundance is the same as in
the LMC is $\sim$16 kpc based on abundance studies involving hundreds of
Cepheid variables (Maciel & Andrievsky 2019; Luck & Lambert 2011).

###  5.2 N 105–1 A vs. Galactic H ii Regions with the Detection of mm-RLs

We compare the properties of N 105–1 A to those of a large sample of Galactic
compact and UC H ii regions, and several HC H ii region candidates, with the
detection of mm-RLs from Kim et al. (2017). Kim et al. (2017) detected mm-RLs
with $n$ from 39 to 65, and $\Delta n$ = 1, 2, 3, and 4 ($\alpha$-, $\beta$-,
$\gamma$-, and $\delta$-transitions) toward 178 sources using the IRAM 30 m
(HPBW$\sim$29′′) and Mopra 22 m (HPBW$\sim$36′′) telescopes. Based on the
known distances to 120 sources, the physical scales probed by the Kim et al.
(2017)’s single-dish observations range from 0.1 pc to $\sim$1 pc. These
physical scales are similar to those we trace with ALMA in the LMC: 0.51 pc (3
mm / Band 3) and 0.12 pc (1.2 mm / Band 6), therefore the Kim et al. (2017)
catalog of compact/UC H ii regions is an ideal dataset to utilize to compare
the properties of the LMC and Galactic objects with mm-RL detections. We have
started building a sample of compact/UC H ii regions in the LMC based on the
archival ALMA data and such comparisons will be more conclusive in the future.
Here, we focus on a single object with the first detection of the higher order
hydrogen mm-RLs outside the Galaxy (up to $\Delta n$=7) to understand its
nature.

Kim et al. (2017) targeted 976 compact dust clumps selected from the APEX
Telescope Large Area Survey of the Galaxy (ATLASGAL; Schuller et al. 2009).
About 10,000 dense clumps have been identified based on the ATLASGAL 870
$\mu$m continuum data, covering the pre-stellar, protostellar, and H ii region
massive star formation stages (Contreras et al. 2013, Urquhart et al. 2014;
Csengeri et al. 2014). The Kim et al. (2017)’s IRAM/Mopra sample includes
roughly the same number of the mid-IR bright and mid-IR quiet clumps that are
among the brightest 870 $\mu$m clumps in their categories. Kim et al. (2017)
detected mm-RLs toward 18% of the targeted ATLASGAL massive clumps. Out of 178
clumps with the detection of an $\alpha$-transition, (65, 23, 22) sources were
also detected in (H$n\beta$, H$n\gamma$, H$n\delta$).

Figure 10: Integrated H$n\alpha$ (stacked $\alpha$-transitions; upper panel)
and H40$\alpha$ (lower panel) line intensities as a function of bolometric
luminosity ($L_{\rm bol}$) for the Galactic H ii regions with the detection of
(sub)mm-RLs (Kim et al. 2017) with the position of N 105–1 A indicated with
the red diamond. The integrated H40$\alpha$ line intensity (not corrected for
beam dilution) is plotted for 1 A in both plots. Galactic sources with
asymmetric submm-RLs discussed in Kim et al. (2018) are also overlaid as
indicated in the legend. The integrated H40$\alpha$ line intensity observed
toward 1 A is similar to that observed toward the most luminous Galactic
sources from the Kim et al. (2017, 2018) sample.

About 75% of the clumps with the detection of mm-RLs from Kim et al. (2017)
are associated with both the radio continuum and mid-IR emission (in the 22
$\mu$m band of the Wide-field Infrared Survey Explorer, WISE, survey), and
were previously identified in the literature as compact or UC H ii regions.
Similarly to these Galactic sources, N 105–1 A is detected at radio and mid-IR
wavelengths (including WISE at 22 $\mu$m, although due to a relatively low
angular resolution, the emission is contaminated by the emission from source 1
B). N 105–1 A’s 6 cm radio luminosity ($S_{\rm 6\,cm}\times D_{\rm
LMC,kpc}^{2}$) of $\sim$$6.4\times 10^{4}$ mJy kpc2 lies well within the
distribution of the 6 cm radio luminosities of the Kim et al. (2017)’s H ii
regions with mm-RLs, with the median of $2.4\times 10^{5}$ mJy kpc2 (see Fig.
10 in Kim et al. 2017).

In the top panel of Fig. 10, we show the integrated H$n\alpha$ line
intensities of Galactic H ii regions with the detection of mm-RLs from Kim et
al. (2017, stacked H39$\alpha$, H40$\alpha$, H41$\alpha$, and H42$\alpha$
transitions; not corrected for beam dilution) as a function of bolometric
luminosity ($L_{\rm bol}$), with the position of N 105–1 A indicated. Since
the H40$\alpha$ transition is the only $\alpha$-transition observed toward 1 A
to date, we also show the plot only including Galactic H ii regions with the
detection of the H40$\alpha$ transition. The plots demonstrate that 1 A lies
within the parameter space covered by the Galactic sources. The integrated
H40$\alpha$ line intensity of $\sim$31 K km s-1 observed toward 1 A (Table 1)
is among the highest values reported by Kim et al. (2017) for Galactic H ii
regions. In their sample, the integrated H40$\alpha$ line intensity ranges
from 0.5 to 63.24 K km s-1, with only four sources (out of 60) above 30 K km
s-1: AGAL010.624$-$00.384 (34.34 K km s-1), AGAL034.258$+$00.154 (38.31),
AGAL043.166$+$00.011 (55.96), AGAL012.804$-$00.199 (63.24). In case of the
stacked $\alpha$-transitions (H$n\alpha$), the integrated H$n\alpha$ line
intensity ranges from 0.37 to 64.28 K km s-1 with five sources (out of 178)
above 30 K km s-1, including AGAL333.604$-$00.212 (64.28 K km s-1) that was
not detected in H40$\alpha$. The sources with the highest values of the
integrated H40$\alpha$ and/or H$n\alpha$ line intensities are among the most
luminous in the Kim et al. (2017) sample ($L_{\rm bol}>10^{5}$ $L_{\odot}$).

For example, clump AGAL043.166$+$00.011 is the central part of one of the most
active massive star-forming region in the Galaxy ($M=4.3\times 10^{4}$ M⊙,
Giannetti et al. 2017), W 49, and it is a massive protocluster candidate
(Urquhart et al. 2018). AGAL034.258$+$00.154 and AGAL010.624$-$00.384 are
well-studied massive star-forming regions (G34.26+0.15 and G10.62-0.38)
hosting UC H ii regions and hot cores; both regions are unresolved in Kim et
al. (2017)’s single-dish observations.

N 105–1 A lies within the physical parameter space (FWHM size, EM, and $n_{\rm
e}$; see Section 4.3) covered by the Kim et al. (2017) sample of Galactic H ii
regions with mm-RLs (see their Fig. 16) and consistent with being at the UC ii
region stage of the massive star evolution.

N 105–1 A corresponds to source IRAS 05101$-$6855A in the main IRAS catalog
(Helou & Walker 1988) and IRAS F05101-6856 in the IRAS Faint Source Catalog
(with a better positional match and the same photometry within the
uncertainties; Moshir et al. 1990). The photometry of both IRAS 05101$-$6855A
and IRAS F05101-6856 fulfills the criteria for UC H ii regions proposed by
Wood & Churchwell (1989).

Figure 11: Top panel: Three-color image combining the HCO+ integrated
intensity (red), the 1.2 mm continuum (green), and the H40$\alpha$ integrated
intensity (blue) images with the spectral extraction regions indicated for
spectra shown in the three lower panels. Lower panels: The H40$\alpha$
(unsmoothed), HCO+ (4-3), H13CO+ (3-2) (Sewiło et al. 2022), and 12CO (3-2)
line profile comparison for several locations in N 105–1 A.

####  5.2.1 Galactic H ii Regions with Asymmetric Submm-RLs

A subsample of 104 sources from the Kim et al. (2017) catalog of H ii regions
with mm-RLs/ATLASGAL clumps was observed by Kim et al. (2018) at submm-RLs
with the APEX 12 m telescope (HPBW$\sim$16′′–27′′, tracing linear scales of
$\sim$0.1–1 pc with the 16′′ beam and $\sim$1.2–1.7 pc with the 27′′ beam).
Submm-RLs H25$\alpha$, H28$\alpha$, and H35$\beta$ were observed toward the
majority of the sample, and the H26$\alpha$, H27$\alpha$, and H29$\alpha$,
H30$\beta$ lines toward a small subsample. Submm-RLs were detected toward 93
clumps. Kim et al. (2018) found that clumps associated with submm/mm-RLs are
the most massive and luminous ($L_{\rm bol}>10^{4}$ $L_{\odot}$) clumps in the
Galaxy. The highest-frequency transitions observed toward N 105–1 A and within
the frequency range of transitions observed by Kim et al. (2018) are
(H36$\beta$, H41$\gamma$, H53$\eta$) at (260.03, 257.635, 257.194) GHz. The
properties of 1 A are consistent with the submm/mm-RL sample of clumps and
associated H ii regions from Kim et al. (2017, 2018),

Interestingly, Kim et al. (2018) detected six sources with RL profiles that
are a combination of a narrow and a broad Gaussian components. Three out of
six of these sources are those with similar physical properties to 1 A,
including the integrated H$n\alpha$ line intensity above 30 K km s-1:
AGAL012.804$-$00.199, AGAL034.258$+$00.154, and AGAL043.166$+$00.011. Kim et
al. (2018) argue that the high-velocity components (revealed by either blue-
or red-shifted wings in the RL profiles) toward four H ii regions can be
explained by the presence of high-velocity ionized flows. Submm-RL profiles of
the other two clumps are likely the result of unresolved clusters of compact H
ii regions.

The inspection of the unsmoothed H40$\alpha$ lines detected toward N 105–1 A,
shown in Fig. 11, also reveals asymmetries in line profiles, providing an
additional evidence for the presence of bulk motions in the region (see
Section 4.5).

###  5.3 Molecular Environment of N 105–1 A

We investigate the distribution and kinematics of the molecular gas in the
ALMA field N 105–1 to understand the origin and nature of the source N 105–1
A, and identify any physical processes in addition to the photoionization by
the central star that may be contributing to the measured ionizing flux (e.g.,
a shock ionization), producing bright enough higher-order mm-RLs to be
detected outside the Galaxy. A detailed modeling of the kinematic structure of
the source is out of scope of the present paper.

For our analysis, we use the molecular gas tracers covering a wide range of
gas densities. From our Band 6 observations presented in Sewiło et al. (2022),
we utilize the CS (5–4), SO ${}^{3}\Sigma$ 66–55, CH3OH $5_{0,5}$–$4_{0,4}$ A,
and H13CO+ (3–2) data. These observations trace the physical scales of
$\sim$0.12 pc or $\sim$25,000 au, assuming the LMC distance of 49.59 kpc. We
also use the 13CO (1–0) data from our Band 3 observations that probe the
$\sim$0.47 pc/$\sim$97,000 au scales. The archival Band 7 CO (3–2) and HCO+
(4–3) data allow us to explore the spatial distribution and the velocity
structure of the diffuse and dense gas, respectively, with the highest
available spatial resolution (0.087 pc, $\sim$18,000 au).

We present the integrated intensity (moment 0), intensity weighted velocity
(moment 1), and line full-width at half-maximum (FWHM; calculated from the
intensity weighted velocity dispersion, moment 2) maps in Appendix D: Fig.
(D.1, D.2, D.4, D.6, D.7, D.9) for (13CO, CS, SO, CH3OH, CO, HCO+). The
integrated intensity and velocity maps are available for all the species,
while the FWHM maps are only available if they provide reliable results (13CO,
CS, and SO). For CS, SO, CO, and HCO+, we also show channel maps (Figs. D.3,
D.5, D.8, and D.10, respectively) that provide a more detailed view of the
velocity structure in N 105–1.

####  5.3.1 Evidence for a Cloud-Cloud Collision in the Region Leading to the
Formation of N 105–1 A

Figure 12: Upper panel: Two-color image highlighting the presence of two CO
(3–2) velocity components toward the 1.2 mm continuum source N 105–1 A. The CO
(3–2) redshifted (241–243 km s-1) and blueshifted (233–235 km s-1) velocity
components are shown in red and cyan, respectively. Both the blue and red
single contours correspond to 0.17 Jy beam-1 km s-1 and outline the lower
emission boundaries of the two velocity components. The 1.2 mm continuum
contours are shown in yellow with contour levels of (10, 30, 250)$\sigma$ (see
also Fig. 1). The size of the synthesized beam ($0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}41\times 0\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}33$) is indicated in the lower left corner of
the image. Lower panel: The CO (3–2) position-velocity diagram along the gray
rectangle shown in the image in the upper panel (X–X′). The x-axis shows an
offset from the 1.2 mm continuum peak (source 1 A) indicated with the white
dotted line; the offset directions are labeled X and X′ to ease a comparison
to the image in the upper panel. The horizontal red and cyan dotted lines show
the CO (3–2) integrated velocity ranges used to plot the redshifted and
blueshifted velocity components in the image in the upper panel.  Figure 13:
Three-color mosaic of the star-forming region N 105 combining the H$\alpha$
(red), [S ii] (green), and [O iii] (blue) images from the MCELS survey (Smith
& MCELS Team 1998; Paredes et al. 2015; an angular resolution of
$\lesssim$5′′) is shown in the left panel. The white box indicates the area
shown in the middle panel: a zoom in on the region hosting source N 105–1 A
with the detection of higher-order mm-RLs. The [S ii] image of the same area
is presented in the right panel. White contours in the middle and right panels
correspond to the 3 mm continuum emission with contour levels the same as in
Fig. 1. Green contours in the right panel correspond to (50, 70, 95)% of the
[S ii] emission peak.

The most striking feature in the velocity maps of all the investigated species
is a velocity gradient stretching across the region with the velocity range of
up to 10 km s-1. The 13CO (1–0) observations provide the best view of the
larger scale velocity structure toward N 105–1 due to the largest field-of-
view and the availability of the ALMA 7m data. The 13CO velocity map shows
that the molecular gas velocity gradient extends over $\sim$28′′ ($\sim$6.7
pc) in right ascension and $\sim$18′′ ($\sim$4.3 pc) in declination (see Fig.
D.1). The velocity increases with increasing RA from the area west of 1 A to
the continuum source 1 B in the east. At the higher-velocity end of the
gradient, there is an additional U-shaped filamentary structure extending
$\sim$16′′ ($\sim$3.8 pc) north from 1 B and then bending toward southwest in
the direction of 1 A.

The velocity gradient across N 105–1 is evident in all other velocity maps,
including those for species with a much more compact distribution than 13CO,
i.e., CS (Fig. D.2), SO (Fig. D.4), and even CH3OH (Fig. D.6).

The sharpest view of the diffuse gas kinematics toward N 105–1 (although least
extensive; see Fig. D.1 and D.7) is provided by the 12CO (3–2) observations
due to the highest spatial resolution and the high dynamic range. The
inspection of the CO velocity channels reveals two distinct velocity
components in N 105–1 that seem to intersect at the position of N 105–1 A
(Fig. D.8). These velocity components are highlighted in the two-color image
shown in the upper panel of Fig. 12, combining the redshifted (241–243 km s-1)
and blueshifted (233–235 km s-1) CO (3–2) emission in N 105–1. The redshifted
CO emission has an S-shaped filamentary structure extending to the 1.2 mm
continuum source 1 B in the northeast. The blueshifted CO emission has a ring-
like structure with some extended emission (“a tail”) pointing roughly toward
northwest. Source N 105–1 A is located at the overlap region of the two
velocity components, hinting on a possibility that the formation of the
massive protostar is a consequence of the collision between the S-shaped
filamentary and the ring-line clouds.

The cloud-cloud collision has been shown to be an important mechanism in the
formation of massive clusters and isolated high-mass stars (see Fukui et al.
2021 for a review). To date, evidence for cloud-cloud collision has been found
in over 50 Galactic star-forming regions (e.g., W 49 N, e.g., Serabyn et al.
1993; Sgr B2, e.g., Ginsburg et al. 2018; G35.20$-$0.74, e.g., Dewangan et al.
2017) and in several extragalactic star-forming regions (in the LMC, M33, and
the Antennae Galaxies; e.g., Fukui et al. 2015, 2017; Tokuda et al. 2019,
2020; see Table 1 in Fukui et al. 2021).

We investigate the collisional signatures in N 105–1 to explore a possible
origin of the O-star exciting the UC H ii region N 105–1 A. We extracted the
CO (3–2) position-velocity diagram (PV-diagram) along the rectangular region
extending from northeast to southwest over 20′′ ($\sim$4.8 pc) and centered on
the proposed collision site (see Fig. 12). At the position of the continuum
source 1 A, the PV-diagram shows the velocity width of the CO gas of $>$10 km
s-1, larger than predicted by the standard scaling relation (a correlation
between the spherical radius $R$ in pc and line width $\sigma$ in km s-1 of
the form $\sigma\propto R^{\,\alpha}$ with $\alpha\approx 0.5$ for the Galaxy,
e.g., Solomon et al. 1987; see e.g., Wong et al. 2019 for the LMC clouds) for
a cloud of a few pc in size under typical Galactic and LMC conditions.

The S-shaped redshifted component is distributed over $\sim$4.5 pc in the CO
PV-diagram with the central velocity of $\sim$242.5 km s-1. The spatially less
extended component with the broad velocity range suddenly appears close to the
position of 1 A (between the offset of $-$0.5 to 1 pc). Such a feature in the
PV-diagram has been observed in other LMC (e.g., N 159 W; Tokuda et al. 2019,
2022) and Galactic regions, and is thought to be a signature of the cloud-
cloud collision. Theoretical simulations of a cloud-cloud collision predict
that two colliding clouds appear as a single broad and continuous cloud in the
PV-diagram (sometimes V-shaped) with the intermediate-velocity gas produced by
the collisional interaction (e.g., Haworth et al. 2015b, a; see Fukui et al.
2021 for a review).

A hub-filamentary structure of the molecular line emission observed toward N
105–1 A (e.g., CS, CO, HCO+, and CH3OH; see Appendix D; see also Section
5.3.2) provides indirect evidence for the cloud-cloud collision event in the
region. The hub-filamentary structure of the molecular gas in the compressed
layer between the colliding clouds is predicted by the cloud-cloud collision
models in the presence of the magnetic field (see Inoue et al. 2018).

The [S ii] $\lambda\lambda$6716, 6731 Å data from the Magellanic Cloud
Emission Line Survey (MCELS, $\sim$2′′ resolution; Smith & MCELS Team 1998)
provide additional indirect evidence for the cloud-cloud collision. The [S ii]
emission traces shock-excited gas and reveals ionization fronts. Figure 13
shows that in N 105–1, the peak of the [S ii] emission is located toward
southeast from 1 A, near the proposed site of the cloud-cloud collision. In
addition to the [S ii] image, Fig. 13 also shows the three-color mosaic
combining the MCELS H$\alpha$, [O iii] $\lambda$5007 Å, [S ii]
$\lambda\lambda$6716, 6731 Å images that allows us to further explore the
stellar radiation feedback in the region. H$\alpha$ tracers all the ionized
gas, while [O iii] reveals high-excitation regions ionized by O stars. Since N
105 is associated with the OB association (LH 31), the previous generation of
OB stars, it is not surprising that the [O iii] emission is present throughout
the entire region.

The CO (3–2) peak integrated intensity around 1 A is $\sim$400 K km s-1, which
corresponds to the H2 peak column density ($N_{\rm H_{2},peak}$) of $2\times
10^{23}$ cm-2, assuming a CO-to-H2 conversion factor in the LMC of $5\times
10^{20}$ cm-2 (K km s-1)-1 (e.g., Hughes et al. 2010) and the CO (3–2)/CO
(1–0) ratio of 1. Observations show that molecular clouds traced by CO with
$N_{\rm H_{2},peak}>10^{23}$ cm-2 and a collision velocity of $\sim$10 km s-1
can form up to 10 OB-type stars (Enokiya et al. 2021; Abe et al. 2022). The
empirical relation between $N_{\rm H_{2},peak}$ and the number of formed OB
stars predicts the formation of 7 OB stars as a result of the collision of
clouds with $N_{\rm H_{2},peak}$ similar to that observed toward N 105–1
(Enokiya et al. 2021). Only one O-type star is observed toward N 105–1 which
is inconsistent with that prediction. The compact nature of colliding clouds
(a few pc in size) may explain the discrepancy, or the clouds may still be in
an early stage of collision. There are several star-forming regions in the
Galaxy with $N_{\rm H_{2},peak}$ as large as $1\times 10^{23}$ cm-2 where only
one OB star was formed as a result of the cloud-cloud collision (e.g.,
G24.85$+$0.09, G24.71$-$0.13, G24.68$-$0.16, Dewangan et al. 2018; BD $+$40
4124, Looney et al. 2006); however, the relative velocities between the
colliding CO clouds is lower than in N 105–1. While not entirely consistent
with the Galactic observations, it is certainly possible that an intense,
localized cloud-cloud collision event could have been responsible for forming
the massive O star in N 105–1 A, with sufficiently high luminosity to produce
higher-order mm-RLs detectable at a distance of $\sim$50 kpc.

Other cloud-cloud collision signatures include a complementary distribution of
gas from two clouds, as well as a U-shape of the bigger cloud participating in
the collision at the final phase of the collision (see Fukui et al. 2021 for a
review). Identification of the cloud-cloud collision signatures is not always
straightforward. Irregular shapes of the colliding clouds and/or the
projection effects may affect the observed gas morphology and kinematics
significantly, making them difficult to interpret. In general, not all cloud-
cloud collision signatures are observed toward each region where such an event
likely occurred (Fukui et al. 2021 and references therein).

To further investigate the cloud-cloud collision event in N 105–1, it will be
necessary to obtain the ALMA Full Array observations (12m, 7m, and Total
Power) with lower-$J$ CO transitions (tracing lower-density gas) to ensure
that the molecular clouds are traced in their entirety.

####  5.3.2 Zooming-in on N 105–1 A: Dense Molecular Gas Distribution and
Kinematics

Filaments are the prevalent feature in the images of N 105–1 A; their hub-like
morphology provides indirect evidence that the formation of the source
ionizing the UC H ii N 105–1 A, or a (proto)cluster it is likely a part of,
was triggered by the cloud-cloud collision (Section 5.3.1). Figure 14 shows
the HCO+ integrated intensity and velocity images, incorporating the emission
in the 235.5–238 km s-1 velocity range where the filamentary structure of the
dense gas traced by HCO+ is most evident. The filaments that are coincident
with the H$\alpha$-dark region (see Section 3 and Fig. 2) converge in the
central hub at the location of 1 A. The emission from other molecular species
also trace filaments, including CO and CS (see Appendix D). The CH3OH emission
clearly traces filaments as well, even though it is much less extended than
the emission from other species (Fig. D.6). Dust in the filaments is traced by
the 1.2 mm continuum emission.

In addition to large-scale filaments extending to the north from the 1 A
continuum peak, the high-resolution HCO+ images reveal several gas streamers
pointing toward the HCO+ and the continuum emission peaks. Approximate
positions of the HCO+ filaments and streamers are outlined in the single
channel images presented in Fig. 14. The presence of filaments and streamers
indicates that the mass accretion toward the molecular clump hosting the UC H
ii region may still be ongoing.

There is a strong indication in the HCO+ data of the presence of multiple
velocity components in the region that could (at least partially) explain
multi-peaked and asymmetric line profiles. The HCO+ integrated intensity
(moment 0) image in Fig. D.9 reveals three emission peaks, forming an
incomplete ring roughly around the 1.2 mm continuum peak and extending over
$\sim$0.37 pc (or $\sim$76,100 au), with the brightest peak to the west of the
continuum peak. The velocity (moment 1) image in the same figure, shows a
sharp boundary between the higher-velocity gas in the west associated with the
brightest HCO+ emission peak and the lower-velocity gas in other directions,
including the remaining two HCO+ emission peaks. The HCO+ channels maps in
Fig. D.10 reveal an additional, significantly fainter lower-velocity component
with the velocity peak at $\sim$236 km s-1 located slightly to the southwest
from the continuum peak. The 1.2 mm continuum peak is offset from the center
of the apparent HCO+ emission ring by $\sim$0.15 pc ($\sim$76,100 au) toward
the higher-velocity component, but it is not coincident with it.

We highlight the lower- and higher-velocity components observed toward 1 A in
Fig. 15 in two HCO+ integrated intensity images, incorporating the emission
from the velocity ranges 233–235 km s-1 (blueshifted) and 241–243 km s-1
(redshifted). The lower- and higher-velocity components are well-separated
spatially and kinematically in these images. On the larger scale, the velocity
ranges 233–235 km s-1 and 241–243 km s-1 separate two CO velocity components -
the ring-like and S-shaped clouds, respectively (see Fig. 12). The CO emission
peaks corresponding to the HCO+ peaks in Fig. 15 can easily be identified in
Fig. 12, indicating that the lower-velocity/higher-velocity gas traced by HCO+
is associated with the ring-like/S-shaped cloud traced by CO.

In the lower panel in Fig. 15, we compare the blueshifted and redshifted HCO+
velocity components to the distribution of the ionized gas emission traced by
H36$\beta$, and to the SO, SO2, and CS emission in the region. The images
clearly show the offset between the ionized gas/1.2 mm continuum and the HCO+
peaks; the latter are shifted to the north and are associated with the SO,
SO2, and CS emission peaks.

Figure 14: Upper panel: The HCO+ (4–3) integrated intensity (left) and
velocity (right) images of N 105–1 A, incorporating the emission with
velocities in the 235.5–238 km s-1 range. Pixels corresponding to regions with
the integrated intensity below 3$\sigma$ are masked in both images. Lower
panel: Single channel HCO+ (4–3) intensity images for the emission with the
velocity of 236.5 km s-1 (left) and 237.5 km s-1 (right). In the lower left
panel, dashed lines outline the filamentary structure seen in HCO+ and other
tracers north of the mm/radio continuum peak. In the lower right panel, dashed
lines indicate possible accretion streamers. The 1.2 mm continuum contours are
overlaid on all images for reference, with contour levels of (3, 30)$\sigma$.
The ALMA beam size is indicated in red in the bottom left corners of images in
the upper panel.

Figure 15: Upper panel: The HCO+ (4–3) integrated intensity images of N 105–1
A incorporating the emission with velocities in the 233–235 km s-1 (left) and
241–243 km s-1 (right) ranges, corresponding to the blueshifted and redshifted
CO velocity components, respectively, shown in Fig. 12. Lower left panel:
Three-color mosaic combining the HCO+ integrated intensity images shown in the
upper panel (for the velocity range 241–243 km s-1 in red and for 233–235 km
s-1 in blue), and the H36$\beta$ recombination line (green) images. Solid
magenta/cyan/yellow contours correspond to the CS/SO/SO2 emission with contour
levels of (20, 50, 90)%/(20, 50, 90)%/50% of the CS/SO/SO2 emission peak
(Sewiło et al. 2022). Dashed white contours correspond to the H36$\beta$
emission with contour levels of (20, 90)% of the H36$\beta$ emission peak.
Lower right panel: The same image as in the lower left panel, zoomed in on the
central region. Blue/red contours correspond to the HCO+
blueshifted/redshifted emission with contour levels of (20, 90)% times the
corresponding emission peak. Dashed white contours are the same as in the
lower left panel. The black cross symbol in both images in the lower panel
indicate the position of the 1.2 mm continuum peak.

####  5.3.3 Accretion Shocks and Sulfur Chemistry

The spatial distribution and abundances of the S-bearing species (SO2 and SO
in particular) support the hypothesis that the accretion from filaments is
still ongoing in N 105–1 A. The emission originating from warm ($T_{\rm
ex}>50$ K) SO and SO2 has been suggested as a good tracer of accretion shocks
at the interface between the disk and the envelope in protostellar systems
(e.g., Sakai et al. 2017; Oya et al. 2019; Artur de la Villarmois et al.
2019). Under comparable physical conditions, the same chemistry may originate
at the interface between the filament/streamer and the clump/core where the
infalling material may cause a shock, raising the temperature of gas and dust
and drive endoergic chemical processes that would enhance the abundance of SO
and SO2. The distribution of the SO and SO2 emission in N 105–1 A, i.e., the
offset to the north from the continuum peak toward the region where the
filaments seem to converge, and the high SO2 temperature derived by Sewiło et
al. (2022) at this location ($\sim$100 K, see Section 3) indicate that such
chemistry may be at work toward this source.

To investigate this possibility, we compared our data to the predictions of
models of nonmagnetized, irradiated $J$-type protoplanetary disk accretion
shocks presented by van Gelder et al. (2021). These models were computed using
the Paris-Durham shock code (Flower & Pineau des Forêts 2003) and considered
preshock densities, $n_{0}$ ($=n_{\rm H}=2\,n({\rm H_{2}})+n({\rm H})+n({\rm
H^{+}})$) of $10^{5}-10^{8}$ cm-3, shock velocities $V_{\rm s}$ in the range
1–10 km s-1, and irradiation by multiples of the (Mathis) UV radiation field,
$G_{0}$. van Gelder et al. (2021) discussed the viable formation routes for
warm gas-phase SO and SO2 (and the resulting abundances) under variation of
$n_{0}$, $V_{\rm s}$, and $G_{0}$.

For the peak H2 column density of $2.0\times 10^{23}$ cm-2 determined based on
the CO observations (see Section 5.3.1) and the source size of 0.13 pc
(Section 5.2), we estimate $n({\rm H_{2}})$ toward 1 A of $\sim$$5\times
10^{5}$ cm-3, assuming that the extent of the source along the line of sight
is the same as in the plane of the sky. In molecular clouds, hydrogen is
predominantly in a molecular form, thus $n_{\rm H}\approx 2\,n({\rm H_{2}})$,
resulting in $n_{\rm H}$ toward 1 A of $\sim$106 cm-3, within the range of
$n_{0}$ explored by van Gelder et al. (2021).

To compare our results to the shock models of van Gelder et al. (2021), we
calculate the ratio of the SO2 and SO column densities ($N_{\rm SO_{2}}$ and
$N_{\rm SO}$, respectively) toward 1 A. $N_{\rm SO_{2}}$ and $N_{\rm SO}$ are
$(1.6^{+0.3}_{-0.2})\times 10^{14}$ cm-2 and $(3.5^{+0.4}_{-0.3})\times
10^{14}$ cm-2, respectively, resulting in the $N_{\rm SO_{2}}$/$N_{\rm SO}$
ratio of 0.46${}^{+0.4}_{-0.3}$ (Sewiło et al. 2022). For a UV radiation field
strength of $G_{0}=1$ (in units of the Mathis interstellar radiation field,
ISRF), this ratio implies low preshock densities of $n_{0}<10^{6}$ cm-3 and
shock speeds of either $V_{\rm s}<2$ km s-1 or $V_{\rm s}>6.5$ km s-1 (Fig. 10
of van Gelder et al. 2021). At $G_{0}=100$, the $N_{\rm SO_{2}}$/$N_{\rm SO}$
ratio observed toward 1 A can again be similarly reproduced with either
$n_{0}<10^{6}$ cm-3 and $V_{\rm s}<2$ km s-1 or $n_{0}>10^{7}$ cm-3 and
$V_{\rm s}>6$ km s-1.

We only consider the $N_{\rm SO_{2}}$/$N_{\rm SO}$ ratio to compare our
observations to the shock model predictions. We do not compare the absolute
values of $N_{\rm SO_{2}}$ and $N_{\rm SO}$ measured toward 1 A to the maximum
values predicted by the van Gelder et al. (2021)’s shock models for different
physical conditions because of the differences in elemental abundances between
the Galaxy and the LMC. The abundance of atomic S and O in the LMC is lower
than in the Galaxy by a factor of 2.6 and 2.2, respectively (Acharyya & Herbst
2015 and references therein). However, as discussed in van Gelder et al.
(2021), for $n_{0}\lesssim 10^{6}$ cm-3, the abundances of both SO and SO2
drop in a similar way in the environment with the reduced abundance of atomic
S and O. The abundance of atomic C (2.5 lower in the LMC) is also relevant
since it reacts with SO, forming CS, and thus is a main destruction pathway
for this molecule (e.g., Hartquist et al. 1980). Decreasing the abundance of
atomic C results in higher abundances of both SO and SO2 in shocks.

Overall, our measured $N_{\rm SO_{2}}$/$N_{\rm SO}$ ratio in 1 A is consistent
with shocks propagating with $V_{\rm s}>6.5$ km s-1 in gas with low preshock
densities ($n_{0}<10^{6}$ cm-3). Sewiło et al. (2022) derived the SO2
temperature toward the 1.2 mm continuum peak / SO2 peak in N 105–1 A of
$(96\pm 20)$ K / $(98\pm 20)$ K, indicating that the observed composition of 1
A could be explained by simple models of postevaporation chemistry. For shocks
propagating with velocities $\gtrsim$4 km s-1, both SO and SO2 are efficiently
formed in the gas phase. The OH radical is crucial for the formation of SO and
SO2 through the route:

$\rm S\buildrel OH\over{\longrightarrow}SO\buildrel
OH\over{\longrightarrow}SO_{2}$.

It is formed through photodissociation of H2O that was formed before the shock
cooled down to below $\sim$300 K. It can also be efficiently produced in
shocks through the endothermic reaction between H2 and atomic O. Atomic S, as
well as SH, can be produced in hydrogen atom abstraction reactions with H2S in
hot gas following the thermal desorption of H2S ice from grain mantles into
the gas (Charnley 1997):

$\rm H_{2}S\buildrel H\over{\longrightarrow}SH\buildrel
H\over{\longrightarrow}S$.

The reactions between SH and atomic O further increase the abundance of SO and
SO2:

$\rm SH\buildrel O\over{\longrightarrow}SO\buildrel
OH\over{\longrightarrow}SO_{2}$.

The abundance of radicals such as OH and SH depends not only on the
temperature in the shock, but also on the strength of the local UV field. The
UV radiation is stronger in the LMC than in the Galaxy and locally in 1 A,
there is an additional contribution from the O-type star.However, although the
source is deeply embedded and the effects of the strong UV radiation are
somewhat diminished, the UV radiation may still be strong enough to promote
the efficient formation of SO and SO2.

The nondetection of the SiO (6–5) transition toward 1 A by Sewiło et al.
(2022) is consistent with the van Gelder et al. (2021)’s shock models since
they do not include high-velocity shocks powerful enough to destroy dust
grains and release the Si atoms to the gas, making them available for chemical
reactions and leading to the formation of SiO (e.g., Schilke et al. 1997;
Gusdorf et al. 2008; Sánchez-Monge et al. 2013). SiO may form (less
efficiently) by grain-grain interactions in higher-density regions (e.g.,
Guillet et al. 2007); however, the higher-excitation SiO transitions have been
found to originate in the high velocity gas (e.g., Leurini et al. 2014).

van Gelder et al. (2021) do not predict any significant increase in the CH3OH
abundance except for the highest preshock densities above $\sim$108 cm-3 and
$V_{\rm s}\gtrsim 10$ km s-1 where ice mantles containing CH3OH can be
thermally desorbed from warm dust grains. Our observations are consistent with
this prediction since only faint CH3OH lines have been detected toward the 1.2
mm continuum and SO2 peaks in 1 A. The CH3OH emission is cold with the
temperature of $\sim$12 K/$\sim$22 K toward the continuum/SO2 peak (Sewiło et
al. 2022), thus its origin by thermal desorption from ices is unlikely. The
possible origin of both methanol and SO2 emission in the star-forming region
is discussed in detail in Sewiło et al. (2022).

####  5.3.4 Shock Activity vs. Photoionization

We explore other shock tracers covered by our observations and those available
in the literature to search for strong shock activity toward N 105–1 A and
identify potential contribution(s) from the shock ionization to the observed
ionized gas emission. Oliveira et al. (2006) suggested the presence of the
outflow in 1 A based on the Br$\alpha$ data (see Section 3), but we do not
find any compelling evidence supporting their conclusion. Out of all available
data discussed in this paper, CO (3–2) is the best outflow tracer. The CO line
profiles (or those for any other species) do not show high-velocity ($>$10 km
s-1) wings, the outflow signature. Other classical diagnostics of shocks
include the grain-destruction products (SiO and S-bearing species such as SO2
and SO; e.g., Caselli et al. 1997; Schilke et al. 1997; Gusdorf et al. 2008)
and grain sputtering products (e.g., CH3OH; e.g., Jørgensen et al. 2020 and
references therein). As mentioned above, the observations of SiO (6–5) toward
1 A resulted in a non-detection (Sewiło et al. 2022). Faint CH3OH emission is
detected south of the continuum peak, but it is cold and not in the outflow
direction proposed by Oliveira et al. (2006).

The optical [S ii] $\lambda\lambda$6716, 6731 Å emission (MCELS, see Section
5.3.1) is a high-velocity shock tracer; however, a non-detection toward 1 A
can be the result of high extinction. Near-IR molecular hydrogen emission
lines that are thermally excited through shocks, are also commonly used as
outflow tracers. However, H2 lines are also excited by fluorescence through UV
photons from OB stars, therefore additional information is often needed to
confirm the origin of the H2 emission. Sewiło et al. (2022) reported the
detection of the H2 v = 1–0 S(1) line at 2.12 $\mu$m toward 1 A with the Very
Large Telescope/$K$-band Multi Object Spectrograph (VLT/KMOS). The H2 emission
is observed over the entire $2\hbox to0.0pt{.\hss}{{}^{\prime\prime}}8\times
2\hbox to0.0pt{.\hss}{{}^{\prime\prime}}8$ field of view with the bright peak
coincident with the ionized gas emission peak (at the continuum peak),
indicating that excitation from UV photons is the dominant excitation
mechanism at this location. Due to a lack of other evidence for the presence
of the outflow, it is unlikely that the extended H2 emission detected toward 1
A traces high-velocity shocks. The H2 2.12 $\mu$m emission gets brighter
toward the west of the KMOS field that, similarly to the region of the peak
emission, may be the results of the photoexcitation from nearby OB stars.

The spectra discussed in Sewiło et al. (2022) extracted from the area around
the 1 A continuum emission peak (as the mean within the contour corresponding
to 50% of the 1.2 mm continuum emission peak) and around the SO2 emission peak
(as the mean within the contour corresponding to 50% of the SO2 emission peak)
show rather narrow molecular lines ($\sim$3–5 km s-1; including H13CO+, SO2,
SO, CS, CH3OH, see Appendix A), indicating that no significant bulk motions
are present. However, all the molecular and hydrogen recombination lines have
asymmetric profiles often with multiple peaks, depending on the position
within N 105–1 A. Figure 11 shows examples of the CO, HCO+, H13CO+, and
H40$\alpha$ lines in several locations in 1 A. The shape of line profiles in
some regions suggest the presence of an infall (two line peaks with the
brighter one at lower velocity); however, the overall distribution of line
profiles is inconsistent with the large-scale infall motions. It is possible
that the complex kinematics at the site of the cloud-cloud collision, the
projection effects, or unresolved multiple sources (or a combination of
thereof) make it difficult to disentangle signatures of different kinematic
motions.

##  6 Summary and Conclusions

Using the ALMA Band 3 and Band 6 observations, we detected the H40$\alpha$,
H36$\beta$, H50$\beta$, H41$\gamma$, H57$\gamma$, H49$\epsilon$, H53$\eta$,
H54$\eta$, and tentatively H55$\theta$ mm-RLs toward the 1.2 mm continuum
source N 105–1 A in the N 105 star-forming region in the LMC. The detection of
the hydrogen $\gamma$, $\epsilon$, and $\eta$ mm-RLs in the LMC constitutes
the first extragalactic detection of these higher-order transitions. The
$\eta$-transitions are detected for the first time at millimeter wavelengths.

The (H40$\alpha$, H50$\beta$, H57$\gamma$)/(H36$\beta$, H41$\gamma$,
H49$\epsilon$, H53$\eta$, H54$\eta$, H55$\theta$) Band 3/6 observations with
the spatial resolution of $\sim$2$\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}$1/$\sim$0$\hbox
to0.0pt{.\hss}{{}^{\prime\prime}}$49 probe the $\sim$0.51 pc/$\sim$0.12 pc
scales, similar to single-dish studies on Galactic massive star-forming
regions.

  1. –

We performed the LTE and non-LTE analysis of the mm-RLs with $\Delta n\leq 8$
covered by our observations of N 105–1 A using the XCLASS software. We were
unable to obtain reliable measurements of the physical properties of the
source ($T_{\rm e}$ and EM) in any of the three sets of mm-RLs we modeled
(Band 3 only, Band 6 only, and all mm-RLs across both bands). The XCLASS
analysis showed that the observed mm-RL intensities are consistent with those
predicted by the LTE model, and revealed considerable beam dilution effect in
the Band 3 observations.

  2. –

Using the H40$\alpha$ line and the 3 mm continuum data, we measured the LTE
electron temperature of $\sim$10,900 K toward N 105–1 A. This temperature is
similar to the electron temperature observed toward Galactic H ii regions at
large Galactocentric distances ($\sim$18 kpc) where the oxygen abundance /
metallicity is comparable to that found in the LMC.

  3. –

We determined the number of Lyman continuum photons ($N_{\rm Ly}$) required to
ionize the source of $\sim$$1.3\times 10^{49}$ s-1 toward N 105–1 A, based on
the derived $T_{\rm e}$ and the 3 mm free-free continuum emission. This value
of $N_{\rm Ly}$ corresponds to the star with the O5.5 V spectral type with the
luminosity of $2.6\times 10^{5}$ $L_{\odot}$ and mass of 34.4 $M_{\odot}$
(Martins et al. 2005).

  4. –

The physical properties of N 105–1 A: FWHM size of $(0.13\pm 0.02)$ pc, EM of
$(8.9\pm 1.3)\times 10^{7}$ pc cm-6, and $n_{\rm e}$ of $(2.6\pm 0.3)\times
10^{4}$ cm-3, are within the size, EM, and $n_{\rm e}$ ranges observed toward
a sample of 178 Galactic compact and UC H ii regions with the detection of
(sub)mm-RLs (Kim et al. 2017, 2018). The physical properties of 1 A and its
association with the IR and radio emission are consistent with it being an UC
H ii region.

  5. –

The H40$\alpha$ line integrated intensity observed toward 1 A is among the
highest (for both H40$\alpha$ and H$n\alpha$, the stacked
$\alpha$-transitions) reported for Galactic H ii regions by Kim et al. (2017).
The sources with the highest values of the integrated H40$\alpha$ and/or
H$n\alpha$ line intensities are among the most luminous in the Kim et al.
(2017) sample ($L_{\rm bol}>10^{5}$ $L_{\odot}$). They include well-known
massive star-forming regions such as the central part of W 49, G34.26$+$0.15,
and G10.62$-$0.38 (unresolved in Kim et al. 2017’s observations).

  6. –

The results of the SED fitting with the Robitaille (2017)’s 2D radiative
transfer YSO SED models indicate an earlier evolutionary stage for 1 A where
both the circumstellar envelope and the disk are present (Class I YSO). It is
possible that 1 A is in fact an unresolved protocluster, containing the UC H
ii region (the most massive star that is evolving the fastest), and one or
more less massive protostars. We find no clear outflow signatures in the
available data (including reliable outflow tracers), but our view may be
obscured by geometry and overlapping velocity components.

  7. –

There is evidence for large scale motions of the ionized gas in N 105–1 A,
including broadened and asymmetric RL profiles, and the H40$\alpha$ velocity
gradient. However, we are unable to unambiguously identify a single phenomenon
responsible for the observed kinematic structure. In the absence of outflow
signatures, the core rotation and/or higher-velocity ionized gas flows in the
region are likely candidates.

  8. –

The high-resolution CO (3–2) data reveal two distinct velocity components in N
105–1 that appear to intersect at the position of 1 A, hinting on a
possibility that star formation in the region was triggered by the cloud-cloud
collision. Such a scenario is supported by the CO gas distribution in the PV-
diagram, a high relative velocity between the clouds ($>$10 km s-1), the
filamentary structure of the molecular gas observed in multiple tracers, and
the presence of the shocked gas (traced by [S ii]) at a proposed collision
site.

  9. –

The HCO+ data reveal multiple velocity components in N 105–1 A that could
explain multi-peaked and asymmetric line profiles observed for multiple
molecular species. We identified three low-velocity peaks and one high-
velocity HCO+ peak in 1 A with velocities consistent with the two (blueshifted
and redshifted) colliding CO clouds.

  10. –

We identified HCO+ gas filaments and streamers toward N 105–1 A indicating an
ongoing accretion onto the clump harboring the UC H ii region and (likely) an
associated embedded (proto)cluster, formed as a result of the cloud-cloud
collision. Chemistry of S-bearing species (SO and SO2) observed toward 1 A is
consistent with the accretion shock model predictions.

  11. –

Our observations demonstrate that ALMA’s high resolution and high sensitivity
enable detailed Galactic-like studies of compact and UC H ii regions and their
molecular environments outside the Galaxy.

We thank the anonymous referee for the constructive report and insightful
comments that helped us improve the manuscript. The material is based upon
work supported by NASA under award number 80GSFC21M0002. We thank Robert
Gruendl and You-Hua Chu for providing the MCELS2 H$\alpha$ image of N 105, and
Sean Points for providing the flux-calibrated, continuum-subtracted MCELS
H$\alpha$, [S ii], and [O iii] images of the LMC. S.B.C. acknowledges support
from the NASA Planetary Science Division Internal Scientist Funding Program
through the Fundamental Laboratory Research work package (FLaRe). K.T. was
supported by a NAOJ ALMA Scientific Research grant No. 2022-22B, Grants-in-Aid
for Scientific Research (KAKENHI) of Japan Society for the Promotion of
Science (JSPS; grant Nos. JP21H00049, JP20H05645, and 21K13962). A.S.M.
acknowledges support from the RyC2021-032892-I grant funded by
MCIN/AEI/10.13039/501100011033 and by the European Union ‘Next
GenerationEU’/PRTR, as well as the program Unidad de Excelencia María de
Maeztu CEX2020-001058-M. The National Radio Astronomy Observatory is a
facility of the National Science Foundation operated under cooperative
agreement by Associated Universities, Inc. This paper makes use of the
following ALMA data: ADS/JAO.ALMA#2019.1.01720.S,
ADS/JAO.ALMA#2017.1.00093.S, and
ADS/JAO.ALMA#2019.1.01770.S. ALMA is a partnership of ESO (representing its
member states), NSF (USA) and NINS (Japan), together with NRC (Canada), NSC
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.

##  Appendix A ALMA Band 3 and Band 6 Spectra of N 105–1 A

In Fig. A.1, we show the ALMA Band 3 spectrum of N 105–1 A covering the
H57$\gamma$, H40$\alpha$, and H50$\beta$ recombination lines (see Table 1),
and the SO ${}^{3}\Sigma$ 32–21 line with the rest frequency of 99,299.87 MHz
and the upper energy level ($E_{\rm U}$) of 9.2 K.

For completeness, we provide the ALMA Band 6 spectra of N 105–1 A in Fig. A.2
that were analyzed in detail in Sewiło et al. (2022). The H13CO+ (3–2), H13CN
(3–2), CS (5–4), SO ${}^{3}\Sigma$ 66–55, and multiple transitions of CH3OH
and SO2 (see Table 3 in Sewiło et al. 2022) are detected toward 1 A. The H2CS
71,6–61,5 line is tentatively detected. Hydrogen recombination lines detected
(H36$\beta$, H41$\gamma$, H49$\epsilon$, H53$\eta$, and H54$\eta$) and
tentatively detected (H55$\theta$) in our Band 6 observations are indicated.

Figure A.1: ALMA Band 3 spectrum of N 105–1 A: a full spectral window covering
hydrogen recombination lines. The detected spectral lines are labeled.

Figure A.2: ALMA Band 6 spectra of N 105–1 A. The detected (black) and
tentatively detected (gray) spectral lines are labeled. The results of the
molecular line analysis are presented in Sewiło et al. (2022).

##  Appendix B XCLASS Modeling: Technical Details, Example Model Fit and Error
Estimation

As mentioned in Section 4.2, a contribution of each mm-RL to a model spectrum
is described in XCLASS by a certain number of components with each component
characterized by a set of physical parameters. These are the source size
$\theta_{\rm source}$ (in arcsec), electron temperature $T_{\rm e}$ (K),
emission measure EM (pc cm-6), line width $\Delta v$ (km s-1), velocity offset
$v_{\rm offset}$ (km s-1), and the position along the line of sight.

XCLASS models the mm-RLs/RRLs by solving the 1D radiative transfer equation
assuming LTE conditions and an isothermal source,

$\displaystyle T_{\rm mb}(\nu)=\sum_{m,c\in i}$
$\displaystyle\Bigg{[}\eta\left(\theta_{\rm
source}^{m,c}\right)\left[S^{m,c}(\nu)\left(1-e^{-\tau_{\rm
total}^{m,c}(\nu)}\right)\right.$ $\displaystyle\left.\,\,+I_{\rm
bg}(\nu)\left(e^{-\tau_{\rm total}^{m,c}(\nu)}-1\right)\right]\Bigg{]}$
$\displaystyle\,\,+\left(I_{\rm bg}(\nu)-J_{\mathrm{CMB}}\right),$

where the sums go over the indices $m$ for molecule / atom, and $c$ for
component, respectively. In addition, $\eta(\theta^{m,c})$ describes the beam
filling (dilution) factor, $S^{m,c}(\nu)$ the source function, $I_{\rm bg}$
the background intensity, and $J_{\mathrm{CMB}}$ the intensity of the cosmic
microwave background. The optical depth $\tau_{\rm total}^{m,c}(\nu)$ of mm-
RL/RRLs (in LTE) is given by Gordon & Sorochenko (2002)

$\displaystyle\tau_{{\rm RRL},\nu}$
$\displaystyle=\frac{\pi\,h^{3}\,e^{2}}{(2\pi\,m_{e}\,k_{B})^{3/2}\,m_{e}\,c}\cdot{\rm
EM}\cdot\frac{n_{1}^{2}\,f_{n_{1},n_{2}}}{T_{e}^{3/2}}$ (B2)
$\displaystyle\quad\times\exp\left[\frac{Z^{2}\,E_{n_{1}}}{k_{B}T_{e}}\right]\,\left(1-e^{-h\nu_{n_{1},n_{2}}/k_{B}T_{e}}\right)\,\phi_{\nu},$

where the oscillator strength444For the absorption oscillator strength
$f_{n_{1},n_{2}}$, we applied Menzel’s approximation (Menzel 1968)
$f_{n_{1},n_{2}}\approx n_{1}\,M_{\Delta n}\left(1+1.5\frac{\Delta
n}{n_{1}}\right)$ with $M_{\Delta n}=0.190775$, $0.026332$, $8.105620\cdot
10^{-3}$, $3.491680\cdot 10^{-3}$, $1.811850\cdot 10^{-3}$, $1.058470\cdot
10^{-3}$, $6.712780\cdot 10^{-4}$, and $4.521900\cdot 10^{-4}$ for $\Delta
n=n_{2}-n_{1}$ = 1, …, 8. $f_{n_{1},n_{2}}$, the transition frequency555The
transition frequency $\nu_{n_{1},n_{2}}$ is given as $\nu_{n,n+\Delta
n}=R_{M}\cdot c\cdot\left[n^{-2}-(n+\Delta n)^{-2}\right]$.
$\nu_{n_{1},n_{2}}$, and the energy of the lower state666In Eq. B2,
$E_{n_{1}}$ is the energy of level $n_{1}$ below the continuum and is given by
$E_{n_{1}}=R_{M}\cdot n_{1}^{-2}$, where $R_{M}$ is the Rydberg constant.
$E_{n_{1}}$ of a particular mm-RL/RRL with main quantum number $n_{1}$ are
obtained from an embedded database describing mm-RL/RRL transitions up to
$\Delta n=8$ ($\theta$-transitions). This means that in our analysis, we
consider all mm-RL transitions within the given frequency range up to $\theta$
transitions. In Eq. (B2), the term $\phi_{\nu}$ represents the line profile
function.

In contrast to the LTE case, where the Saha-Boltzmann distribution was used to
determine the electron population, radiative transitions dominate over
collisional transitions in the non-LTE case. To describe the departures from
LTE, one introduces for each electronic level $n$ a so-called departure
coefficient $b_{n}$, relating the electron population in non-LTE ($N_{n}$)
with that in LTE ($N_{n}^{*}$),

$N_{n}=b_{n}\,N_{n}^{*}.$ (B3)

These departure coefficients have been derived by Storey & Hummer (1995) and
depend on both collisional and radiative processes. XCLASS uses their
tabulated coefficients777For non-tabulated temperatures and densities a linear
interpolation procedure is used. No extrapolation is used and the
corresponding departure coefficients are set to one. for all mm-RLs/RRLs. Due
to the fact that the departure coefficients depends on the electron density,
the non-LTE description of the mm-RLs/RRLs requires an additional parameter,
the electron density $N_{\rm e}$ (in cm-3) for each component.

The general results of the XCLASS analysis of our Band 3 and Band 6 data for N
105–1 A are described in Section 4.2. Here, we present the example model fit
and error estimation for the XCLASS fit to the Band 6 transitions. Figure B.1
shows spectra centered on rest frequencies of the mm-RL transitions (up to
$\theta$-transitions, $\Delta n=8$) covered by our Band 6 (H36$\beta$,
H41$\gamma$, H49$\epsilon$, H53$\eta$, H54$\eta$, and H55$\theta$) and Band 3
(H40$\alpha$, H50$\beta$, H57$\gamma$, H67$\epsilon$, H74$\eta$, and
H77$\theta$) observations, with the XCLASS “LTE with one component” model fit
overlaid (‘Scenario 1’ in Section4.2). The corner plot illustrating the error
estimation for this model is shown in Fig. B.2.

Figure B.1: Spectra of N 105–1 A (shown in blue) centered on rest frequencies
of the mm-RL transitions (up to $\theta$-transitions, $\Delta n=8$) covered by
our Band 3 (top two panels) and Band 6 (bottom two panels) observations,
ordered by increasing frequency (from left to right, top to bottom). The red
solid line shows the XCLASS synthetic spectrum fitted to Band 6 mm-RL
transitions; it represents the pure mm-RL contribution assuming LTE with a
single component. The XCLASS spectrum is extended to lower frequencies (shown
as the red dashed line) to compare line intensities predicted by the LTE model
to our observed Band 3 spectra, revealing the beam dilution effects in Band 3.
The XCLASS spectrum taking into account both mm-RLs and all molecular species
detected toward 1 A (discussed in Sewiło et al. 2022; see also Appendix A) is
shown as the gray solid line. The gray vertical dashed line indicates the
source velocity.  Figure B.2: Corner plot of the MCMC error estimate for the
XCLASS fit of the LTE model with one component shown in Fig. B.1. The 16% and
84% quantiles are plotted as black dashed lines, the 50% quantile (median) is
shown in blue, while the lowest $\chi^{2}$ value is shown in orange. The plot
demonstrates that we cannot reliably determine either the electron temperature
($T_{\rm e}$) or emission measure (EM); it is likely due to the signal-to-
noise ratio being too low for most of the Band 6 transitions.

##  Appendix C Multiwavelength Images and Photometry of N 105–1 A

We provide the multiwavelength photometry (from $\sim$1 $\mu$m to 6 cm) for N
105–1 A in Table C. The image cutouts are shown in Fig. C.1. The references
and technical details are provided in the Table C footnotes, Fig. C.1 caption,
and the text below. The spectral energy distribution (SED) of N 105–1 A for a
wavelength range from $\sim$1.3 $\mu$m to 3 mm is presented in Fig. 7.

In Table C, we provide the following photometric data:

  1. –

the $JHK_{S}$ data from the Infrared Survey Facility (IRSF) Magellanic Clouds
Point Source Survey (Kato et al. 2007); we treat the $J$-band flux density as
an upper limit for the SED fitting since no source is visible in the image at
the position of 1 A. The Visible and Infrared Survey Telescope for Astronomy
(VISTA) $YJK_{S}$ survey of the Magellanic Clouds system (VMC; Cioni et al.
2011) only provides the $K_{s}$-band photometry and it is consistent with the
IRSF measurement;

  2. –

Spitzer Space Telescope Infrared Array Camera (IRAC) 3.6–8.0 $\mu$m data from
the LMC-wide Spitzer “Surveying the Agents of Galaxy Evolution” (SAGE; Meixner
et al. 2006; SAGE Team 2006) survey (e.g., Whitney et al. 2008; Gruendl & Chu
2009; Carlson et al. 2012). Both the point spread function (PSF; SAGE Team
2006) and aperture photometry (Gruendl & Chu 2009) are available (see Table
C). The PSF and aperture photometry agree within the uncertainties. We have
adopted the aperture photometry fluxes for the fitting since it better
accounts for the extended emission surrounding the source at longer IRAC
bands. No Multiband Imaging Photometer for Spitzer (MIPS) 24 or 70 $\mu$m
photometry is available in the existing catalogs. Source 1 A is unresolved
from the neighboring YSO in these bands (see Fig. C.1);

  3. –

the five-band Herschel Space Observatory photometric measurements from Seale
et al. (2014) covering 100–500 $\mu$m; the catalog is based on the
Photoconductor Array Camera and Spectrometer (PACS; 100 and 160 $\mu$m) and
Spectral and Photometric Imaging Receiver (SPIRE; 250–500 $\mu$m) data from
the “Herschel Inventory of the Agents of Galaxy Evolution” (HERITAGE; Meixner
et al. 2013; HERITAGE Team 2013) survey. Source 1 A is unresolved from the
neighboring source between 160 and 500 $\mu$m, thus we use the Herschel flux
densities in these bands as upper limits;

  4. –

the ALMA 870 $\mu$m continuum flux density based on the reprocessed archival
data;

  5. –

ALMA 1.2 mm and 3 mm flux densities reported in Sewiło et al. (2022) and
measured in the present paper, respectively, after removing the contribution
from the free-free emission.

In addition to using the Spitzer catalog data, we have used the Spitzer
InfraRed Spectrograph (IRS) spectrum from Seale et al. (2009) to better
constrain the SED between 5.2 and 37.9 $\mu$m. We extracted 11 data points
from the IRS spectrum that were selected at wavelengths free of fine-structure
emission lines to delineate silicate features and the underlying continuum:
6.89, 8.96, 9.44, 9.95, 10.75, 11.64, 13.63, 15.95, 19.68, 25.58, 29.95
$\mu$m.

In the analysis of the IRS spectra, the scaling factors were applied to match
smoothly the spectrum segments taken under different modules across the full
wavelength range (Seale et al. 2009). For the fitting, we have reverted these
IRS fluxes to their original values for three affected spectrum segments by
removing the corresponding scaling factors, i.e., dividing the fluxes within
SL1 (short wavelength, low resolution; 7.6–14.6 $\mu$m), SH (short wavelength,
high resolution; 9.9–19.3 $\mu$m), and LH (long wavelength, high resolution;
18.9–36.9 $\mu$m) modules by 0.9704, 0.8937, and 1.4173, respectively.

The SED fitting procedure and the fitting results are discussed in Section
4.4.

Figure C.1: Multiwavelength images of N 105–1 (from left to right, top to
bottom): the H$\alpha$ image from the MCELS2 survey (PI: You-Hua Chu), VMC
$Y$\- and $J$-band (Cioni et al. 2011), IRSF $H$-band (Kato et al. 2007), VMC
$K_{s}$-band, Spitzer/SAGE IRAC 3.6–8.0 $\mu$m and MIPS 24 and 70 $\mu$m
(Meixner et al. 2006; SAGE Team 2006), Herschel/HERITAGE PACS 100 and 160
$\mu$m and SPIRE 250–500 $\mu$m (Meixner et al. 2013; HERITAGE Team 2013),
ALMA 1.2 mm (Sewiło et al. 2022), ALMA 3 mm (this paper), and ATCA 3 and 6 cm
(Indebetouw et al. 2004). The 1.2 mm continuum contours with contour levels of
(3, 30) $\times$ the image rms of $6.9\times 10^{-5}$ mJy beam-1 are overlaid
on selected images for reference. Red circles in the 4.5 and 100 $\mu$m images
indicate the positions of Spitzer YSOs associated with 1 A and 1 B (Whitney et
al. 2008; Gruendl & Chu 2009; Seale et al. 2009), while the blue circle shows
the position of the Herschel YSO candidate (Seale et al. 2014).  Table 2:
Multiwavelength Photometry for N 105–1 A

Instruments and Bands | Flux density (uncertainty) | FWHMaaBoth the Band 3 and Band 6 mm-RL spectra for 1 A were extracted for the analysis as the mean within the contour corresponding to 50% of the 1.2 mm continuum emission peak, using data cubes with the native angular resolution. No correction for beam dilution was applied to $T_{\rm b}$ and $\int T_{b}dv$. | Source ID | Ref.
---|---|---|---|---
| (mJy) | (′′) |  |
VMC $Y$ (1.02), $J$ (1.25), $K_{s}$ (2.15 $\mu$m) | -999.9, -999.9, 1.78 (0.10) |  $\lesssim$1 | 558354728325 | 1
IRSF $J$ (1.25), $H$ (1.63) , $K_{s}$ (2.14 $\mu$m) | 0.04 (0.01), 0.20 (0.02), 1.85 (0.07) |  $\sim$1.3 | 05095050-6853052 | 2
2MASS $J$ (1.25), $H$ (1.65) , $K_{s}$ (2.16 $\mu$m) | -999.9, -999.9, 2.26 (0.10) |  $\sim$2.5 | J05095054$-$6853052 | 3
Spitzer/IRAC 3.6, 4.5, 5.8, 8.0 $\mu$mbbBand 6: The fitting was performed on the spectra after applying three iterations of Hanning smoothing, resulting in a channel width of $\Delta v_{\rm instr}$ = (4.51, 4.54, 4.83, 4.54, 4.78) km s-1 for (H36$\beta$, H41$\gamma$, H49$\epsilon$, H53$\eta$, H54$\eta$) lines. Band 3: The fitting was performed on the spectra after applying one iteration of Hanning smoothing, resulting in a channel width of $\Delta v_{\rm instr}$ = 5.91 km s-1 for the H40$\alpha$, H50$\beta$, H57$\gamma$ lines. | 24.45 (0.94), 64.85 (2.54), 156.50 (3.25), 504.70 (14.69) | 1.7, 1.7, 1.9, 2.0 | SSTISAGE1C J050950.53$-$685305.4 | 4
| 28.61 (1.58), 65.25 (3.00), 177.30 (9.80), 514.12 (23.68) | —"— | 050950.53$-$685305.5ccThe magnitudes provided in Gruendl & Chu (2009) were converted to fluxes using the zero-magnitude fluxes of (280.9, 179.7, 115.0, 64.13) Jy for the IRAC (3.6, 4.5, 5.8, 8.0) $\mu$m bands (the same as those used in the SAGE survey). | 5
Herschel/PACS 100, 160 $\mu$m | 41390 (3013), 23080 (1612) | 8.6, 12.6 | HSOBMHERICC J77.459799$-$68.884809 | 7
Herschel/SPIRE 250, 350, 500 $\mu$m | 10060 (652), 4967 (309), 7710 (552) | 18.3, 26.7, 40.5 | —"— | 7
ALMA 870 $\mu$m ( GHz) | 161.7 (0.1)ddMeasured within a 3$\sigma$ contour using the CASA task imstat. Both the dust and free-free emission contribute to the measured flux densities at 1.2 and 3 mm. The contribution of the free-free emission was estimated to be $\sim$50% at 242.4 GHz (Sewiło et al. 2022) and $\sim$3.7% at 99 GHz (see Section 4.3.1), respectively. |  $0.47\times 0.39$, 37$\hbox to0.0pt{.\hss}^{\circ}$0 | N 105–1 A | 9
ALMA 1.2 mm (242.4 GHz) | 71.8 (2.2)ddMeasured within a 3$\sigma$ contour using the CASA task imstat. Both the dust and free-free emission contribute to the measured flux densities at 1.2 and 3 mm. The contribution of the free-free emission was estimated to be $\sim$50% at 242.4 GHz (Sewiło et al. 2022) and $\sim$3.7% at 99 GHz (see Section 4.3.1), respectively. |  $0.51\times 0.47$, 37$\hbox to0.0pt{.\hss}^{\circ}$2 | —"— | 8, 9
ALMA 3 mm (99 GHz) | 48.6 (0.3)ddMeasured within a 3$\sigma$ contour using the CASA task imstat. Both the dust and free-free emission contribute to the measured flux densities at 1.2 and 3 mm. The contribution of the free-free emission was estimated to be $\sim$50% at 242.4 GHz (Sewiło et al. 2022) and $\sim$3.7% at 99 GHz (see Section 4.3.1), respectively. |  $2.49\times 1.83$, 46$\hbox to0.0pt{.\hss}^{\circ}$1 | —"— | 9
ATCA 3 cm (8.6 GHz) | 39 (1) |  $1.82\times 1.24$, 14$\hbox to0.0pt{.\hss}^{\circ}$8 | B0510$-$6857 W | 10
ATCA 6 cm (4.8 GHz) | 26 (1) |  $2.19\times 1.70$, 16$\hbox to0.0pt{.\hss}^{\circ}$7 | —"— | 10

aafootnotetext: The image angular resolution at full-width half maximum
(FWHM). For the near-IR ground-based surveys VMC, IRSF, and 2MASS, the angular
resolution depends on the atmospheric seeing. For the ALMA and ATCA
observations, we provide the synthesized beam sizes and position angles:
$\theta_{\rm min}\times\theta_{\rm maj},PA$.

bbfootnotetext: No MIPS photometry is available for 1 A in Whitney et al.
(2008) and Gruendl & Chu (2009).

References. — (1) Cioni et al. (2011); (2) Kato et al. (2007); (3) Skrutskie
et al. (2006); (4) Whitney et al. (2008), Spitzer PSF photometry from the SAGE
Epoch 1 Catalog; (5) Gruendl & Chu (2009), Spitzer aperture photometry; (6)
Wright et al. (2010); (7) Seale et al. (2014); (8) Sewiło et al. (2022); (9)
this paper; (10) Indebetouw et al. (2004)

##  Appendix D ALMA Integrated Intensity, Velocity, and Line Width Images of N
105–1

We present the integrated intensity (moment 0), intensity weighted velocity
(moment 1), and line width (the full width at half maximum, FWHM, calculated
from the intensity weighted velocity dispersion, or moment 2, image and
deconvolved from the instrumental broadening arising from a finite channel
width) images of the ALMA field N 105–1 centered on the continuum source 1 A
for three molecular species discussed in Sewiło et al. (2022, CS, SO, and
CH3OH), and for 13CO, the data reported in this paper. All three images are
shown for the CS (5–4) transition in Fig. D.2, for SO ${}^{3}\Sigma$ 66–55 in
Fig. D.4, and for 13CO (1–0) in Fig. D.1. The CH3OH $5_{0,5}$–$4_{0,4}$
emission is the faintest and spatially most limited out of the four species
and no reliable line width image was obtained; the CH3OH integrated intensity
and velocity images are shown in Fig. D.6. In addition, in Fig. D.3 and D.5,
we show the channel maps for CS and SO, respectively. The kinematics of the