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

The hydrogen radio and millimeter/submillimeter recombination lines (RRLs and mm/submm-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.

RLs 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 RLs) 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 α (Lyα) spectral line in the UV, the strongest emission line observed toward astrophysical objects. Atoms other than hydrogen produce RLs 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). RLs 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 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 (Δn) indicated with successive letters in the Greek alphabet. The RLs of H from level n + Δn to n are denoted (Hn α, Hn β, Hn γ, Hn δ, Hn , Hn ζ, Hn η...) for Δn = (1, 2, 3, 4, 5, 6, 7...). The α-transitions (Δn = 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 α- and β-transitions were first detected outside the Galaxy in the 1970s. The higher-order (Δn > 2) millimeter/radio transitions had not been detected until our observations reported in this paper.

The first extragalactic detection of RRLs was reported by Mezger et al. (1970). The H109α, He109α, and H137β 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 beamwidth, HPBW ≈ 4'). The H109α 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 ∼ 15''), Peck et al. (1997) detected the H90α, H113β, and He90α (∼8.9 GHz), H92α (8.3 GHz), and again the H109α (5 GHz) RRLs. 30 Dor has been the LMC star-forming region best-studied in RRLs. The H30α and H40α 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 (C30α in N 79) is uncertain since it may in fact be a helium RL (He30α).

The first extragalactic detections of RRLs beyond the Magellanic Clouds were made toward M82 (H166α, Shaver et al. 1977; H102α, Bell & Seaquist 1977; H92α, Chaisson & Rodriguez 1977) and NGC 253 (H102α, Seaquist & Bell 1977). After these first observations of extragalactic RRLs, no new detections were reported until the H53α mm-RL was detected toward NGC 2146 over 10 yr 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 Hn α mm-RLs toward nearby galaxies (e.g., Bendo et al. 2015, 2016, 2017; Michiyama et al. 2020) and the first detection of the fainter Hn β and Hen α 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 α-transitions with principal quantum numbers n = 266–301, detected in the spectrum of the radio quasar 3C 190 (z = 1.1946) with the Low Frequency Array (van Haarlem et al. 2013).

In this paper, we report the detection of the H40α, H36β, H50β, H41γ, H57γ, H49, H53η, and H54η mm-RLs and a tentative detection of the H55θ 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 ± 0.09 (statistical) ± 0.54 (systematic) kpc (Pietrzyński et al. 2019). The mm-RL γ-, -, η-, and θ-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_LMC ∼ 0.3–0.5 Z_⊙, 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 RLs. 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∼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 Sections 4 and 5, respectively. In Section 6, we provide the summary and conclusions of our study.

The N 105–1 field hosting the source with the detection of mm-RLs (N 105–1 A) was observed with ALMA's 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 Common Astronomy Software Applications (CASA; CASA Team et al. 2022). The observations were executed twice on 2019 October 21 with 43 antennas and baselines from 15 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 2019 October 23 with 43 antennas, baselines from 15 to 782 m, and the same calibrators. The total on-source integration time was ∼13.1 minutes including both executions.

The spectral setup included four 1875 MHz spectral windows centered on frequencies of 242.4, 244.8, 257.9, and 259.7 GHz. The continuum was subtracted in the uv domain from the line spectral windows 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 automultithresh masking. The rms sensitivity of 0.05 mJy per 051 × 047 beam (4.4 mK) was achieved in the continuum. A sensitivity of 1.97 mJy (0.15 K) per 054 × 050 beam was achieved in the 242.4 GHz spectral cube; 1.88 mJy (0.15 K) per 053 × 049 beam was achieved in the 244.8 GHz cube; 2.05 mJy (0.16 K) per 051 × 047 beam was achieved in the 257.9 GHz cube; 2.28 mJy (0.18 K) per 051 × 047 beam was achieved in the 259.7 GHz cube; all four data cubes have 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 12-m and 7-m Array observations of two young stellar objects (YSOs) located in N 105–1 (one coinciding with 1 A; see Figure C1) 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 H40α line (a rest frequency of 99.02295 GHz; 12-m Array data only to use with our Band 6 12-m Array data) and ¹³CO (1–0) (110.20135 GHz, an upper state energy, E_U = 5.3 K; combined 12-m and 7-m Array data).

The 12-m Array observations toward N 105–1 were executed on 2018 April 2 and April 3 with 41 and 43 antennas and baselines between 15 and 543 m, and 15 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 ∼4 minutes. The 1875 MHz spectral window centered on the H40α line was divided into 1920 channels of ∼976.6 kHz (∼2.96 km s⁻¹) each. The 59 MHz spectral window centered on the ¹³CO (1–0) line was divided into 1920 channels of ∼30.7 kHz (∼0.084 km s⁻¹) each. The data were calibrated with CASA 5.1.1-5.

The continuum was subtracted in the uv domain from the H40α line spectral window. A sensitivity of 4.4 mJy per 256 × 195 beam (0.11 K) was achieved in the H40α data cube. The H40α spectral cube has a cell size of 03 ×03 × 2.96 km s⁻¹. A sensitivity of 0.25 mJy per 249 × 183 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.

The 7-m Array observations were executed eight times between 2017 December 29 and 2018 January 10 with 10 (one execution) and 11 (seven executions) antennas, with baselines between 8.9 and 48.9 m. The total on-source integration time was ∼19.7 minutes for each target YSO. Both bandpass and flux calibrators were J0522−3627 (seven executions) and J0006−0623 (one execution). J0450−8101 (four executions) and J0529−7245 (four executions) were used as phase calibrators. The 62 MHz spectral window centered on the ¹³CO (1–0) line was divided into 2048 channels of ∼30.3 kHz (∼0.082 km s⁻¹) each. We combined the 12-m and 7-m Array data to image the ¹³CO (1–0) line. The ¹³CO (1–0) spectral cube has a cell size of 03 × 03 × 0.2 km s⁻¹, and it is corrected for primary beam attenuation. The sensitivity of 16.9 mJy per 229 × 167 beam (0.44 K) was achieved in the ¹³CO (1–0) data cube.

The 1.2 and 3 mm continuum images are shown in Figure 1. The Band 3 spectrum from the spectral window covering H RLs 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.

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2.1. ALMA Archival Data

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; Alcaino & Liller 1986; Vallenari et al. 1994). 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 ATCA (a resolution of ∼15 and ∼2'' at 3 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 μm spectrum from the ISAAC instrument on ESO's Very Large Telescope (VLT) displays a very red continuum, very strong Brα and Pfγ RL emission, and a nondetection of the Pfδ line, which was postulated to be the result of a high dust column density in front of the source (A_V ∼ 40 mag). Further evidence supporting the interpretation of the source as an embedded protostar includes broad Brα 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_S – $L^{\prime}$ = 3.9 mag). Strong hydrogen RLs indicate that the massive YSO has started ionizing its immediate surroundings. The recent VLT/K-band Multi Object Spectrograph (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, No. 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, hereafter IRS, 5–38 μm, 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 μm, [Ne ii] 12.8 μm, [Ne iii] 15.5 μm, [S iii] 18.7 μm and 33.5 μm, and [S iii] 34.8 μm, and weak emission lines from the polycyclic aromatic hydrocarbons. 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 ± 2.6 M_⊙ and (1.4 ± 0.2) × 10⁵ 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 Hα, CS, SO, H36β, 1.2 mm, and 6 cm emission are compared in Figure 2. The high-resolution Hα image shows that 1 A is associated with the extended low-intensity Hα emission around the millimeter/radio continuum peak, while the millimeter continuum emission extending to the north from the main peak overlaps with what appears to be the Hα-dark region. The CS emission has a filamentary structure and extends even farther to the north (∼8'' or 1.9 pc), filling in the Hα-dark region. The H36β mm-RL emission (and the emission from all other mm-RLs except H57γ; see below) coincides with the millimeter/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 millimeter/radio continuum peak.

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N 105–1 A is associated with cold CH₃OH (12 ± 1 K) and hot SO₂ (96 ± 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 ≲ 0.1 pc), hot (T_kin ≳ 100 K), and dense (n_H ≳ 10⁶⁻⁷ 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 six 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 Magellanic Clouds (LMC: Shimonishi et al. 2016; Sewiło et al. 2018, 2019; Shimonishi et al. 2020; Sewiło et al. 2022; the Small Magellanic Cloud: Shimonishi et al. 2023).

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 subthermally 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 ∼06 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.

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: H36β (260.03278 GHz), H41γ (257.63549 GHz), H49 (241.86116GHz), H53η (257.19399 GHz), and H54η (243.94239 GHz). An additional transition, H55θ (258.52592 GHz), is tentatively detected. These are all hydrogen RLs with Δn ≤ 8 (up to θ-transitions) falling within the frequency range of our observations (e.g., Gordon & Sorochenko 2002). Three more mm-RLs (H40α, H50β, and H57γ at 99.02295, 99.22521, 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).

We verified the identification of the faintest reliably detected mm-RLs, H53η and H54η, by comparing their observed peak intensities relative to the peak intensity of the H36β 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 Figure 3. There is a good agreement for H41γ and H49 lines. For the H53η and H54η 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 (∼6σ detections). The spatial distribution of the H53η and H54η line emission is in good agreement with other mm-RLs (see Figures 4 and 5), and the lines are too broad for molecular lines (Figure 6).

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To our knowledge, only the α- and β-transitions have been observed outside the Galaxy to date, making the γ-, -, and η-transitions observed toward N 105–1 A the first extragalactic detections. The η-transitions are the first to be detected at millimeter wavelengths.

The H36β, H41γ, H49, H53η, and H54η integrated intensity images are shown in Figure 4 (∼05 resolution), and the H40α, H50β, and H57γ integrated intensity images are shown in Figure 5 (∼22). The mm-RL emission peaks for all observed transitions except H57γ (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 ∼1–1.4 times the beam size at 1.2 mm; see Figure 2 and Figures 4–5). The H57γ emission peak appears to be offset to the east from the continuum peak by ∼06; since this offset corresponds to a relatively small fraction of the 3 mm beam (∼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 (ʋ_LSR), line width (Δʋ_FWHM, the full width at half maximum, FWHM), peak intensity (T_b ), and the integrated intensity (∫T_b dʋ) for each transition. No correction for beam dilution was applied. The observed mm-RLs with the fitted Gauss profiles are shown in Figure 6.

Table 1. Gauss Fitting Results for Hydrogen Recombination Lines Detected toward N 105–1 A ^{a , b}

Transition	Frequency	ʋLSR	Δʋc	Tb	∫Tb dʋ
	(GHz)	(km s−1)	(km s−1)	(K)	(K km s−1)
Band 6
H36β	260.03278	231.9 (0.6)	32.3 (1.3)	0.84 (0.03)	29.0 (1.5)
H41γ	257.63549	231.1 (0.8)	33.5 (1.8)	0.39 (0.02)	14.2 (1.0)
H49	241.86116	227.7 (2.7)	35.7 (6.4)	0.14 (0.02)	5.5 (1.3)
H53η	257.19399	239.3 (1.9)	25.7 (4.5)	0.12 (0.02)	3.3 (0.7)
H54η	243.94239	233.9 (1.4)	24.4 (3.4)	0.12 (0.02)	3.3 (0.6)
Band 3
H40α	99.02295	230.1 (0.6)	31.1 (1.5)	0.91 (0.04)	30.6 (1.9)
H50β	99.22521	229.9 (2.9)	29.3 (6.9)	0.26 (0.05)	8.4 (2.6)
H57γ	98.67189	228.6 (5.7)	25.5 (13.5)	0.13 (0.06)	3.6 (2.4)

Notes.

^a Both 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_b and ∫T_b dʋ. ^b For Band 6 transitions, the fitting was performed on the spectra after applying three iterations of Hanning smoothing, resulting in a channel width of ${\rm{\Delta }}{{\unicode{x0028B}}}_{\mathrm{instr}}$ = (4.51, 4.54, 4.83, 4.54, 4.78) km s⁻¹ for (H36β, H41γ, H49, H53η, H54η) lines. For Band 3 transitions, the fitting was performed on the spectra after applying one iteration of Hanning smoothing, resulting in a channel width of ${\rm{\Delta }}{{\unicode{x0028B}}}_{\mathrm{instr}}$ = 5.91 km s⁻¹ for the H40α, H50β, H57γ lines. ^c The line FWHM corrected for instrumental broadening: ${\rm{\Delta }}{\unicode{x0028B}}=\sqrt{{\rm{\Delta }}{{\unicode{x0028B}}}_{\mathrm{obs}}^{2}-{\rm{\Delta }}{{\unicode{x0028B}}}_{\mathrm{instr}}^{2}}$, where Δʋ_obs is the observed line width, and ${\rm{\Delta }}{{\unicode{x0028B}}}_{\mathrm{instr}}$ is the channel width after Hanning smoothing.

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A comparison of the observed intensity ratios of Hn α 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 H40α and H50β lines detected toward 1 A are an ideal line pair to investigate the conditions under which the observed RLs are formed. The theoretical value of the H50β/H40α 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 H50β/H40α line ratio toward 1 A of 0.29 ± 0.06 is in good agreement with the theoretical value, indicating that the region is likely to be in LTE.

4.2.  XCLASS Modeling of Recombination Lines

4.3. Physical Properties of the Ionized Gas in N 105–1 A Based on the H40α 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 Hnα RL emission. The H40α line is the only α-transition currently available for 1 A.

4.3.1. The Electron Temperature

Assuming the RLs are emitted under pure LTE conditions and pressure broadening of the lines by electron impacts is negligible, the observed RL 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 Hn α lines (Wilson et al. 2009):

$$\begin{eqnarray}\begin{array}{rcl}\displaystyle \frac{{I}_{{\rm{L}}}}{{I}_{{\rm{C}}}}\left(\displaystyle \frac{{\rm{\Delta }}{\unicode{x0028B}}}{\mathrm{km}\,{{\rm{s}}}^{-1}}\right) & \approx & \displaystyle \frac{6985}{a({T}_{{\rm{e}}},\nu )}{\left(\displaystyle \frac{{\nu }_{{\rm{L}}}}{\mathrm{GHz}}\right)}^{1.1}\\ & \cdot & {\left(\displaystyle \frac{{T}_{{\rm{e}}}^{* }}{{\rm{K}}}\right)}^{-1.15}\left(\displaystyle \frac{1}{1+{y}^{+}}\right),\end{array}\end{eqnarray} \tag{ 1 }$$

where I_L is the line intensity, I_C is the free–free continuum intensity, Δʋ is the FWHM line width, ν_L is the Hn α line rest frequency, y⁺ is a singly ionized helium-to-hydrogen abundance ratio (${N}_{{\mathrm{He}}^{+}}/{N}_{{{\rm{H}}}^{+}}$), a(T_e , ν) is a dimensionless factor of the order of 1 slowly varying with T_e and ν, tabulated in Mezger & Henderson (1967). At the frequency of ∼100 GHz, a∼0.9 for T_e from 8000 to 12,000 K, typical for H ii regions. I_L/I_C ratio allows I_L and I_C to be given in any convenient units.

We utilize the H40α line (ν_L = 99.02295 GHz) to estimate ${T}_{{\rm{e}}}^{* }$ in N 105–1 A. The y⁺ parameter cannot be determined for 1 A in this work since no He RL has been detected toward this source in our observations. A typical value of y⁺ in the Galaxy is ∼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⁺ 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⁺ of 0.1 specifically for N 105 using optical spectroscopy.

We obtain I_L Δʋ from the H40α integrated intensity, ∫I_L dʋ; for Gaussian line profiles, ∫I_L dʋ ≈ 1.064 I_L Δʋ (e.g., Brown et al. 1978). We determine ∫I_L dʋ of 1110 ± 10 mJy km s⁻¹ from the H40α integrated intensity (moment 0) map by measuring the emission within the 3σ contour. We measured the 99 GHz continuum flux density of 43.2 ± 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 μm continuum image to estimate the dust emission at 3 mm under the assumption that all the emission at 870 μm originates from dust. We measured the 870 μm 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 following formula: ${S}_{{\rm{B}}3}\,={S}_{{\rm{B}}7}{({\nu }_{{\rm{B}}3}/{\nu }_{{\rm{B}}7})}^{2+\beta }$, where β is a dust emissivity index. We adopt β 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 Equation (1), we determine ${T}_{{\rm{e}}}^{* }$ of 10,940 ± 1350 K toward 1 A, which is consistent with T_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 H50β/H40α 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_e).

4.3.2. Emission Measure and Electron Density

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

$$\begin{eqnarray}\begin{array}{rcl}{\rm{EM}}({\rm{pc}}\,{{\rm{cm}}}^{-6}) & = & 1.74\times {10}^{-3}{\left(\displaystyle \frac{{T}_{{\rm{e}}}}{{\rm{K}}}\right)}^{3/2}\left(\displaystyle \frac{{\nu }_{{\rm{L}}}}{{\rm{GHz}}}\right)\\ & & \cdot \,{\eta }^{-1}\left(\displaystyle \frac{{T}_{{\rm{L}}}}{{\rm{K}}}\right)\left(\displaystyle \frac{{\rm{\Delta }}{\unicode{x0028B}}}{{\rm{km}}\,{{\rm{s}}}^{-1}}\right),\end{array}\end{eqnarray} \tag{ 2 }$$

where T_L is the Hn α line intensity in kelvins, η is the beam-filling factor, and other parameters are the same as in Equation (1). We use T_e = 10,940 K (Section 4.3.1). We calculate T_L Δʋ from the integrated intensity of the H40α line (∫I_L dʋ converted to K km s⁻¹) as described above. The beam-filling factor $\eta ={\theta }_{s}^{2}/({\theta }_{s}^{2}+{\theta }_{\mathrm{beam}}^{2})$, where ${\theta }_{s}^{}$ and ${\theta }_{\mathrm{beam}}^{}$ are the FWHM size of the source and the synthesized beam, respectively. For 1 A, we obtain θ_s of 055 ± 010 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 ± 1.7) × 10⁷ pc cm⁻⁶ toward N 105–1 A.

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

$$\begin{eqnarray}&&\begin{array}{l}{\mathrm{EM}}_{\mathrm{cont}}(\mathrm{pc}\,{\mathrm{cm}}^{-6})=\displaystyle \frac{{\tau }_{{\rm{C}}}}{8.235\times {10}^{-2}\,a({T}_{{\rm{e}}},\nu )\,{T}_{{\rm{e}}}^{-1.35}\,{\nu }^{-2.1}},\end{array}\end{eqnarray} \tag{ 3 }$$

where τ_C is the continuum optical depth, ν is the continuum frequency in GHz, and a(T_e, ν) is the same as in Equation (1). We calculate the peak τ_C and EM_cont based on the 3 mm continuum data.

For the 99.023 GHz/3 mm free–free continuum peak intensity of 43.64 ± 2.10 mJy per 249 × 183 beam (96.4% of the observed continuum peak of 45.27 mJy beam⁻¹), we obtain the brightness temperature (T_b) of 1.19 ± 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 (τ_C) from the formula ${T}_{{\rm{b}}}={T}_{{\rm{e}}}(1-{e}^{-{\tau }_{{\rm{C}}}})$ after correcting T_b for the beam dilution effect: ${\tau }_{{\rm{C}}}=-\mathrm{ln}(1-{\eta }^{-1}\,{T}_{{\rm{b}}}/{T}_{{\rm{e}}})\,=(1.7\pm 0.2)\times {10}^{-3}$. From Equation (3), we obtain the peak EM_cont of (1.0 ± 0.2) × 10⁸ 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_e, we use the expression ${n}_{{\rm{e}}}=\sqrt{\mathrm{EM}/{\rm{\Delta }}s}$, where Δs (in parsecs) 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 (θ_s ) as Δs (0.13 ± 0.02 pc). The resulting n_e is (2.6 ± 0.3) × 10⁴ cm⁻³ for N 105–1 A.

The values of (FWHM size, EM, n_e) of (0.13 pc, ∼9 × 10⁷ pc cm⁻⁶, ∼3 × 10⁴ 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 ≲0.1 pc, n_e ≳ 10⁴ cm⁻³, and EM ≳ 10⁷ 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 10 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, outflow, 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_Ly) required to ionize the source can be expressed by the following formula (e.g., Goudis 1977 and references therein):

$$\begin{eqnarray}\begin{array}{rcl}{N}_{\mathrm{Ly}}({{\rm{s}}}^{-1}) & = & 4.76\times {10}^{48}\left(\displaystyle \frac{{S}_{\nu }^{\mathrm{ff}}}{\mathrm{Jy}}\right){\left(\displaystyle \frac{D}{\mathrm{kpc}}\right)}^{2}\\ & \cdot & {\left(\displaystyle \frac{\nu }{\mathrm{GHz}}\right)}^{0.1}{\left(\displaystyle \frac{{T}_{{\rm{e}}}}{{\rm{K}}}\right)}^{-0.45},\end{array}\end{eqnarray} \tag{ 4 }$$

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

For 1 A, Equation (4) results in N_Ly of 1.32 × 10⁴⁹ photons per second (log N_Ly = 49.12), which corresponds to an O5.5 V star (log N_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 × 10⁵ L_⊙ and 34.4 M_⊙, 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_⊙. In their models, the ionizing flux of log N_Ly = 49.12 corresponds to the O6 V spectral type and luminosity of ∼3 × 10⁵ L_⊙ for the metallicity within the range observed in the LMC (Z = 0.4 Z_⊙). 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

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 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 μm) 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 ∼1.3 μm to 3 mm. The SED is shown in Figure 7, while flux densities with references and technical details are provided in Table 2 in Appendix C. Figure C1 shows the images of N 105–1, from optical (Hα) to radio (6 cm) wavelengths.

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We do not include the Herschel photometry except the 100 μm 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 Figure C1). 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 μm; see also Figure C1). We also use the 870 μm 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 3σ contour; see Table 2 in Appendix 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 total continuum emission at this wavelength.

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 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 χ² 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 ${\chi }^{2}-{\chi }_{\mathrm{best}}^{2}\lt F\,{n}_{\mathrm{data}}$ where ${\chi }_{\mathrm{best}}^{2}$ is χ² of the best-fit model across all the model sets and all models, n_data is the number of data points used for the fitting, and F is a threshold parameter that is determined empirically. In the 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 χ² across all the model sets and all models. The SED with the best-fit model overlaid is shown in Figure 8. The best-fit model comes from the model set spubsmi and includes a central source, passive disk, Ulrich (rotationally flattened) envelope, bipolar cavities, and an ambient interstellar medium (see an inset cartoon in Figure 8; adapted 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.

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We determine the luminosity of the central source of 1.3 × 10⁵ L_⊙ from the stellar radius and effective temperature returned by the fitter for the best-fit YSO model using the Stefan–Boltzmann equation. To obtain the 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 stage; the evolutionary tracks were calculated for the initial stellar masses from 0.1 to 350 M_⊙, 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_⊙ and 1.1 × 10⁵ yr, 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 2 times lower, and the mass is ∼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 RL values for luminosity and mass to be more reliable, as they are based on direct observations of the ionizing photons.

4.5. Ionized Gas Kinematics

There are three groups of the line broadening mechanisms that can affect the RL widths: the natural, Doppler, and pressure (collisional) broadenings (e.g., Gordon & Sorochenko2002). The effects of the natural and collisional broadening are negligible for mm/submm-RLs (e.g., Brocklehurst & Seaton1972; Roelfsema & Goss 1992; Gordon & Sorochenko 2002); thus, we only consider the Doppler thermal and nonthermal contributions to the line widths of mm-RLs detected toward N 105–1 A.

The thermal line width can be determined from the following formula:

$$\begin{eqnarray}&&\begin{array}{l}{\rm{\Delta }}{{\unicode{x0028B}}}_{\mathrm{th}}=\sqrt{8\,\mathrm{ln}2\,\tfrac{{k}_{{\rm{B}}}{T}_{{\rm{e}}}}{\mu {m}_{{\rm{H}}}}},\end{array}\end{eqnarray} \tag{ 5 }$$

where k_B is the Boltzmann constant, T_e is the gas electron temperature, μ is the molecular weight in atomic units (for atomic hydrogen, μ_H = 1.008 amu), and m_H is the mass of the hydrogen atom. For the temperature of the hydrogen gas of 10,940 ± 1350 K, Δʋ_th is 22.3 ± 1.4 km s⁻¹.

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

$$\begin{eqnarray}&&{\rm{\Delta }}{{\unicode{x0028B}}}_{\mathrm{nth}}=\sqrt{{\rm{\Delta }}{{\unicode{x0028B}}}^{2}-{\rm{\Delta }}{{\unicode{x0028B}}}_{\mathrm{th}}^{2}}\end{eqnarray} \tag{ 6 }$$

where Δʋ is the measured line width corrected for instrumental broadening (see Table 1). The average line width for the three brightest RLs is 32.3 km s⁻¹, resulting in Δʋ_nth of ∼23.4 km s⁻¹, similar to Δʋ_th.

The relatively large nonthermal 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 RLs detected toward 1 A: H40α (Band 3) and H36β (Band 6; ∼5× higher angular resolution than the Band 3 observations); all the images are shown in Figure 9. The H40α velocity map reveals the velocity gradient of 17 km s⁻¹ over 35 (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 H40α 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 H40α data have to be treated with caution since the signal-to-noise ratio of the H40α line is low in the outermost regions. The H36β emission traces much smaller scales close to the central star; the overall trend in velocity seems to be similar to that seen in H40α. The H36β lines are the broadest at the continuum peak.

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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 H40α and H36β velocity gradients are not homogenous; therefore, the rotation is not the most likely interpretation, but cannot be ruled out.

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 versus Galactic H ii Regions

5.2. N 105–1 A versus Galactic H ii Regions with the Detection of Millimeter 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 Δn = 1, 2, 3, and 4 (α-, β-, γ-, and δ-transitions) toward 178 sources using the IRAM 30 m (HPBW ∼ 29'') and Mopra 22 m (HPBW ∼ 36'') telescopes. Based on the known distances to 120 sources, the physical scales probed by the Kim et al. (2017) single-dish observations range from 0.1 to ∼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 data set 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 Δn = 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 μm continuum data, covering the pre-stellar, protostellar, and H ii region massive star formation stages (Contreras et al. 2013; Csengeri et al. 2014; Urquhart et al. 2014). The Kim et al. (2017) IRAM/Mopra sample includes roughly the same number of the mid-IR bright and mid-IR quiet clumps that are among the brightest 870 μm 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 α-transition, (65, 23, 22) sources were also detected in (Hn β, Hn γ, Hn δ).

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 μm 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 μm, 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 (${S}_{6\,\mathrm{cm}}\times {D}_{\mathrm{LMC},\ \mathrm{kpc}}^{2}$) of ∼6.4 × 10⁴ mJy kpc² lies well within the distribution of the 6 cm radio luminosities of the Kim et al. (2017) H ii regions with mm-RLs, with the median of 2.4 × 10⁵ mJy kpc² (see Figure 10 in Kim et al. 2017).

In the top panel of Figure 10, we show the integrated Hn α line intensities of Galactic H ii regions with the detection of mm-RLs from Kim et al. (2017, stacked H39α, H40α, H41α, and H42α transitions; not corrected for beam dilution) as a function of bolometric luminosity (L_bol), with the position of N 105–1 A indicated. Since the H40α transition is the only α-transition observed toward 1 A to date, we also show the plot only including Galactic H ii regions with the detection of the H40α transition. The plots demonstrate that 1 A lies within the parameter space covered by the Galactic sources. The integrated H40α line intensity of ∼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 H40α 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), and AGAL012.804−00.199 (63.24). In the case of the stacked α-transitions (Hn α), the integrated Hn α 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 H40α. The sources with the highest values of the integrated H40α and/or Hn α line intensities are among the most luminous in the Kim et al. (2017) sample (L_bol > 10⁵ L_⊙).

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For example, clump AGAL043.166+00.011 is the central part of one of the most active massive star-forming regions in the Galaxy (M = 4.3 × 10⁴ 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 the Kim et al. (2017) single-dish observations.

N 105–1 A lies within the physical parameter space (FWHM size, EM, and n_e; see Section 4.3) covered by the Kim et al. (2017) sample of Galactic H ii regions with mm-RLs (see their Figure 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).

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 submillimeter wavelengths with the APEX 12 m telescope (HPBW∼16''–27'', tracing linear scales of ∼0.1–1 pc with the 16'' beam and ∼1.2–1.7 pc with the 27'' beam). The observations included submm-RLs H25α, H27α, H28α, H29α, H30α, and H35β. Submm-RLs were detected toward 93 clumps. Kim et al. (2018) found that the clumps associated with mm/submm-RLs are the most massive and luminous (L_bol > 10⁴ L_⊙) 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 (H36β, H41γ, H53η) at (260.03, 257.635, 257.194) GHz. The properties of 1 A are consistent with the mm/submm-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 Hn α 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 blueshifted or redshifted 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 H40α lines detected toward N 105–1 A, shown in Figure 11, also reveals asymmetries in line profiles, providing an additional evidence for the presence of bulk motions in the region (see Section 4.5).

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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 ³Σ 6₆–5₅, CH₃OH 5_0,5–4_0,4 A, and H¹³CO⁺ (3–2) data. These observations trace the physical scales of ∼0.12 pc or ∼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 ∼0.47 pc/∼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, ∼18,000 au).

We present the integrated intensity (moment 0), intensity weighted velocity (moment 1), and line FWHM (calculated from the intensity weighted velocity dispersion, moment 2) maps in Appendix D: Figures D1, D2, D4, D6, D7, and D9 for ¹³CO, CS, SO, CH₃OH, CO, and HCO⁺, respectively. The integrated intensity and velocity maps are available for all the species, while the line 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 (Figures D3, D5, D8, and D10, 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

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 7-m Array data. The ¹³CO velocity map shows that the molecular gas velocity gradient extends over ∼28'' (∼6.7 pc) in R.A. and ∼18'' (∼4.3 pc) in decl. (see Figure D1). The velocity increases with increasing R.A. from the area west of 1 A to the continuum source 1 B in the northeast. At the higher-velocity end of the gradient, there is an additional U-shaped filamentary structure extending ∼16'' (∼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 (Figure D2), SO (Figure D4), and even CH₃OH (Figure D6).

The sharpest view of the diffuse gas kinematics toward N 105–1 (although least extensive; see Figures D1 and D7) 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 (Figure D8). These velocity components are highlighted in the two-color image shown in the upper panel of Figure 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.

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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'' (∼4.8 pc) and centered on the proposed collision site (see Figure 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 parsecs and line width σ in kilometers per second of the form σ ∝ R^α with α ≈ 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 parsecs in size under typical Galactic and LMC conditions.

The S-shaped redshifted component is distributed over ∼4.5 pc in the CO PV diagram with the central velocity of ∼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. 2015a, 2015b; 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] λ λ 6716, 6731 Å data from the Magellanic Cloud Emission Line Survey (MCELS, ≲5'' 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 the southeast from 1 A, near the proposed site of the cloud–cloud collision. In addition to the [S ii] image, Figure 13 also shows the three-color mosaic combining the MCELS Hα, [O iii] λ5007 Å, [S ii] λ λ 6716, 6731 Å images that allows us to further explore the stellar radiation feedback in the region. Hα traces 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.

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The CO (3–2) peak integrated intensity around 1 A is ∼400 K km s⁻¹, which corresponds to the H₂ peak column density (${N}_{{{\rm{H}}}_{2},\mathrm{peak}}$) of 2 × 10²³ cm⁻², assuming a CO-to-H₂ conversion factor in the LMC of 5 × 10²⁰ 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 ${N}_{{{\rm{H}}}_{2},\mathrm{peak}}\gt {10}^{23}$ cm⁻² and a collision velocity of ∼10 km s⁻¹ can form up to 10 OB-type stars (Enokiya et al. 2021; Abe et al. 2022). The empirical relation between ${N}_{{{\rm{H}}}_{2},\mathrm{peak}}$ and the number of formed OB stars predicts the formation of 7 OB stars as a result of the collision of clouds with ${N}_{{{\rm{H}}}_{2},\mathrm{peak}}$ similar to that observed toward N 105–1 A (Enokiya et al. 2021). Only one O-type star is observed toward N 105–1 A, which is inconsistent with that prediction. The compact nature of colliding clouds (a few parsecs 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 ${N}_{{{\rm{H}}}_{2},\mathrm{peak}}$ as large as 1 × 10²³ 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 are lower than those 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 ∼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 (12-m Array, 7-m Array, 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 Hα-dark region (see Section 3 and Figure 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 (Figure D6). Dust in the filaments is traced by the 1.2 mm continuum emission.

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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 Figure 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 multipeaked and asymmetric line profiles. The HCO⁺ integrated intensity (moment 0) image in Figure D9 reveals three emission peaks, forming an incomplete ring roughly around the 1.2 mm continuum peak and extending over ∼0.37 pc (or ∼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⁺ channel maps in Figure D10 reveal an additional, significantly fainter lower-velocity component with the velocity peak at ∼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 ∼0.15 pc (∼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 Figure 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 and 241–243 km s⁻¹ separate two CO velocity components—the ring-like and S-shaped clouds, respectively (see Figure 12). The CO emission peaks corresponding to the HCO⁺ peaks in Figure 15 can easily be identified in Figure 12, indicating that the lower-velocity/higher-velocity gas traced by HCO⁺ is associated with the ring-like/S-shaped cloud traced by CO.

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In the lower panel in Figure 15, we compare the blueshifted and redshifted HCO⁺ velocity components to the distribution of the ionized gas emission traced by H36β, and to the SO, SO₂, and CS emission in the region. The images clearly show the offset between the peak of the ionized gas emission (coincident with the 1.2 mm continuum peak) and the HCO⁺ peaks; the latter are shifted to the north and are associated with the SO, SO₂, and CS emission peaks.

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_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; Artur de la Villarmois et al. 2019; Oya 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 (∼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, n₀ (=n_H = 2 n(H₂) +n(H) + n(H⁺)) of 10⁵ − 10⁸ cm⁻³, shock velocities V_s in the range 1–10 km s⁻¹, and irradiation by multiples of the (Mathis) UV radiation field, G₀. van Gelder et al. (2021) discussed the viable formation routes for warm gas-phase SO and SO₂ (and the resulting abundances) under variation of n₀, V_s, and G₀.

For the peak H₂ column density of 2.0 × 10²³ 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(H₂) toward 1 A of ∼5 × 10⁵ 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_H ≈ 2 n(H₂), resulting in n_H toward 1 A of ∼10⁶ cm⁻³, within the range of n₀ 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 (${N}_{{\mathrm{SO}}_{2}}$ and N_SO, respectively) toward 1 A. ${N}_{{\mathrm{SO}}_{2}}$ and N_SO are $({1.6}_{-0.2}^{+0.3})\times {10}^{14}$ cm⁻² and $({3.5}_{-0.3}^{+0.4})\times {10}^{14}$ cm⁻², respectively, resulting in the ${N}_{{\mathrm{SO}}_{2}}$/N_SO ratio of ${0.46}_{-0.3}^{+0.4}$ (Sewiło et al. 2022). For a UV radiation field strength of G₀ = 1 (in units of the Mathis interstellar radiation field), this ratio implies low preshock densities of n₀ < 10⁶ cm⁻³ and shock speeds of either V_s < 2 km s⁻¹ or V_s > 6.5 km s⁻¹ (Figure 10 of van Gelder et al. 2021). At G₀ = 100, the ${N}_{{\mathrm{SO}}_{2}}$/N_SO ratio observed toward 1 A can again be similarly reproduced with either n₀ < 10⁶ cm⁻³ and V_s < 2 km s⁻¹ or n₀ > 10⁷ cm⁻³ and V_s > 6 km s⁻¹.

We only consider the ${N}_{{\mathrm{SO}}_{2}}$/N_SO ratio to compare our observations to the shock model predictions. We do not compare the absolute values of ${N}_{{\mathrm{SO}}_{2}}$ and N_SO measured toward 1 A to the maximum values predicted by the van Gelder et al. (2021) 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 that in the Galaxy by factors of 2.6 and 2.2, respectively (Acharyya & Herbst 2015 and references therein). However, as discussed in van Gelder et al. (2021), for n₀ ≲ 10⁶ 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 times 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 ${N}_{{\mathrm{SO}}_{2}}$/N_SO ratio in 1 A is consistent with shocks propagating with V_s > 6.5 km s⁻¹ in gas with low preshock densities (n₀ < 10⁶ 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 ± 20) K/(98 ± 20) K, indicating that the observed composition of 1 A could be explained by simple models of postevaporation chemistry. For shocks propagating with velocities ≳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

$$\begin{eqnarray*}&&{\rm{S}}\mathop{\longrightarrow }\limits^{\mathrm{OH}}\mathrm{SO}\mathop{\longrightarrow }\limits^{\mathrm{OH}}{\mathrm{SO}}_{2}.\end{eqnarray*}$$

It is formed through photodissociation of H₂O that was formed before the shock cooled down to below ∼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):

$$\begin{eqnarray*}&&{{\rm{H}}}_{2}{\rm{S}}\mathop{\longrightarrow }\limits^{{\rm{H}}}\mathrm{SH}\mathop{\longrightarrow }\limits^{{\rm{H}}}{\rm{S}}.\end{eqnarray*}$$

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

$$\begin{eqnarray*}&&\mathrm{SH}\mathop{\longrightarrow }\limits^{{\rm{O}}}\mathrm{SO}\mathop{\longrightarrow }\limits^{\mathrm{OH}}{\mathrm{SO}}_{2}.\end{eqnarray*}$$

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) 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 ∼10⁸ cm⁻³ and V_s ≳ 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 ∼12 K/∼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 N 105 is discussed in detail in Sewiło et al. (2022).

5.3.4. Shock Activity versus 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 Brα 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 nondetection (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] λ λ 6716, 6731 Å emission (MCELS; see Section 5.3.1) is a high-velocity shock tracer; however, a nondetection 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 ʋ = 1–0 S(1) line at 2.12 μm toward 1 A with the VLT/KMOS. The H₂ emission is observed over the entire 28 × 28 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 μm 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 (∼3–5 km s⁻¹; including H¹³CO⁺, SO₂, SO, CS, CH₃OH; see Appendix A), indicating that no significant molecular gas bulk motions are present. However, all the molecular and hydrogen RLs 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 H40α lines in several locations in 1 A. The shape of line profiles in some regions suggests 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.

Using the ALMA Band 3 and Band 6 observations, we detected the H40α, H36β, H50β, H41γ, H57γ, H49, H53η, H54η, and tentatively H55θ 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 γ, , and η mm-RLs in the LMC constitutes the first extragalactic detection of these higher-order transitions. The η-transitions are detected for the first time at millimeter wavelengths.

The (H40α, H50β, H57γ)/(H36β, H41γ, H49, H53η, H54η, H55θ) Band 3/6 observations with the spatial resolution of ∼21/∼049 probe the ∼0.51 pc/∼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 Δn ≤ 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_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 Band 6 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 H40α line and the 3 mm continuum data, we measured the LTE electron temperature of ∼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 (∼18 kpc) where the oxygen abundance (thus metallicity) is comparable to that found in the LMC.
• 3.  
  We determined the number of Lyman continuum photons (N_Ly) required to ionize the source of ∼1.3 × 10⁴⁹ s⁻¹ toward N 105–1 A, based on the derived T_e and the 3 mm free–free continuum emission. This value of N_Ly corresponds to the star with the O5.5 V spectral type with the luminosity of 2.6 × 10⁵ L_⊙ and mass of 34.4 M_⊙ (Martins et al. 2005).
• 4.  
  The physical properties of N 105–1 A—FWHM size of (0.13 ± 0.02) pc, EM of (8.9 ± 1.3) × 10⁷ pc cm⁻⁶, and n_e of (2.6 ± 0.3) × 10⁴ cm⁻³—are within the size, EM, and n_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 H40α line integrated intensity observed toward 1 A is among the highest (for both H40α and Hn α, the stacked α-transitions) reported for Galactic H ii regions by Kim et al. (2017). The sources with the highest values of the integrated H40α and/or Hn α line intensities are among the most luminous in the Kim et al. (2017) sample (L_bol > 10⁵ L_⊙). 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 observations).
• 6.  
  The results of the SED fitting with the Robitaille (2017) 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 H40α 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 multipeaked 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. The 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 No. 80GSFC21M0002. We thank Robert Gruendl and You-Hua Chu for providing the MCELS2 Hα image of N 105, and Sean Points for providing the flux-calibrated, continuum-subtracted MCELS Hα, [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). C.-H.R.C. acknowledges support from the Deutsches Zentrum für Luft- und Raumfahrt (DLR) grant NS1 under contract No. 50 OR 2214. 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 CASSIS interactive spectrum analyzer (http://cassis.irap.omp.eu; Vastel et al. 2015) was used for the data analysis. CASSIS has been developed by IRAP-UPS/CNRS. 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.

In Figure A1, we show the ALMA Band 3 spectrum of N 105–1 A covering the H57γ, H40α, and H50β RLs (see Table 1), and the SO ³Σ 3₂–2₁ line with the rest frequency of 99,299.87 MHz and the upper energy level (E_U) of 9.2 K.

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For completeness, in Figure A2 we provide the ALMA Band 6 spectra of N 105–1 A that were analyzed in detail in Sewiło et al. (2022). The H¹³CO⁺ (3–2), H¹³CN (3–2), CS (5–4), SO ³Σ 6₆–5₅, and multiple transitions of CH₃OH and SO₂ (see Table 3 in Sewiło et al. 2022) are detected toward 1 A. The H₂CS 7_1,6–6_1,5 line is tentatively detected. Hydrogen RLs detected (H36β, H41γ, H49, H53η, and H54η) and tentatively detected (H55θ) in our Band 6 observations are indicated.

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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 θ_source (in arcseconds), electron temperature T_e (kelvins), EM (pc cm⁻⁶), line width Δʋ (kilometers per second), velocity offset ʋ_offset (kilometers per second), 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,

$$\begin{eqnarray}\begin{array}{rcl}{T}_{\mathrm{mb}}(\nu ) & = & \displaystyle \displaystyle \sum _{m,c\in i}\left[\eta \left({\theta }_{\mathrm{source}}^{m,c}\right)\left[{S}^{m,c}(\nu )\left(1-{e}^{-{\tau }_{\mathrm{total}}^{m,c}(\nu )}\right)\right.\right.\\ & & \left.\left.\,+{I}_{\mathrm{bg}}(\nu )\left({e}^{-{\tau }_{\mathrm{total}}^{m,c}(\nu )}-1\right)\right]\right]\\ & & +\ \left({I}_{\mathrm{bg}}(\nu )-{J}_{\mathrm{CMB}}\right),\end{array}\end{eqnarray} \tag{ B1 }$$

where the sums go over the indices m for molecule/atom, and c for component, respectively. In addition, η(θ^m,c ) describes the beam-filling (dilution) factor, S^m,c (ν) describes the source function, I_bg describes the background intensity, and J_CMB describes the intensity of the cosmic microwave background. The optical depth ${\tau }_{\mathrm{total}}^{m,c}(\nu )$ of mm-RLs/RRLs (in LTE) is given by Gordon & Sorochenko (2002):

$$\begin{eqnarray}&&\begin{array}{c}{\tau }_{\mathrm{RRL},\nu }=\frac{\pi \,{h}^{3}\,{e}^{2}}{{(2\pi \,{m}_{e}\,{k}_{{\rm{B}}})}^{3/2}\,{m}_{e}\,c}\cdot \mathrm{EM}\cdot \frac{{n}_{1}^{2}\,{f}_{{n}_{1},{n}_{2}}}{{T}_{e}^{3/2}}\\ \quad \times \ \exp \left[\frac{{Z}^{2}\,{E}_{{n}_{1}}}{{k}_{{\rm{B}}}{T}_{e}}\right]\left(1-{e}^{-h{\nu }_{{n}_{1},{n}_{2}}/{k}_{{\rm{B}}}{T}_{e}}\right)\,{\phi }_{\nu },\end{array}\end{eqnarray} \tag{ B2 }$$

where the oscillator strength ¹⁹ ${f}_{{n}_{1},{n}_{2}}$, the transition frequency ²⁰ ${\nu }_{{n}_{1},{n}_{2}}$, and the energy of the lower state ²¹ ${E}_{{n}_{1}}$ of a particular mm-RL/RRL with main quantum number n₁ are obtained from an embedded database describing mm-RL/RRL transitions up to Δn = 8 (θ-transitions). This means that, in our analysis, we consider all mm-RL transitions within the given frequency range up to θ transitions. In Equation (B2), the term ϕ_ν 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 (${N}_{n}^{* }$),

$$\begin{eqnarray}&&{N}_{n}={b}_{n}\,{N}_{n}^{* }.\end{eqnarray} \tag{ B3 }$$

These departure coefficients have been derived by Storey & Hummer (1995) and depend on both collisional and radiative processes. XCLASS uses their tabulated coefficients ²² for all mm-RLs/RRLs. Due to the fact that the departure coefficients depend on the electron density, the non-LTE description of the mm-RLs/RRLs requires an additional parameter, the electron density N_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 B1 shows spectra centered on rest frequencies of the mm-RL transitions (up to θ-transitions, Δn = 8) covered by our Band 6 (H36β, H41γ, H49, H53η, H54η, and H55θ) and Band 3 (H40α, H50β, H57γ, H67, H74η, and H77θ) observations, with the XCLASS "LTE with one component" model fit overlaid. The corner plot illustrating the error estimation for this model is shown in Figure B2.

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We provide the multiwavelength photometry (from ∼1 μm to 6 cm) for N 105–1 A in Table 2. The image cutouts are shown in Figure C1. The references and technical details are provided in the Table 2 footnotes, Figure C1 caption, and the text below. The SED of N 105–1 A for a wavelength range from ∼1.3 μm to 3 mm is presented in Figure 7.

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In Table 2, 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 VISTA survey of the Magellanic Clouds system (VMC, YJ K _S bands; Cioni et al. 2011) only provides the K_s -band photometry, and it is consistent with the IRSF measurement.
• 2.  
  The Spitzer Space Telescope Infrared Array Camera (IRAC) 3.6–8.0 μm data from the LMC-wide Spitzer "Surveying the Agents of Galaxy Evolution" (SAGE; Meixner et al. 2006; SAGE 2006) survey (e.g., Whitney et al. 2008; Gruendl & Chu 2009; Carlson et al. 2012). Both the point-spread function (PSF; SAGE 2006) and aperture photometry (Gruendl & Chu 2009) are available (see Table 2). 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 μm photometry is available in the existing catalogs. Source 1 A is unresolved from the neighboring YSO in these bands (see Figure C1).
• 3.  
  The five-band Herschel Space Observatory photometric measurements from Seale et al. (2014) covering 100–500 μm; the catalog is based on the Photoconductor Array Camera and Spectrometer (PACS; 100 and 160 μm) and Spectral and Photometric Imaging Receiver (SPIRE; 250–500 μm) data from the "Herschel Inventory of the Agents of Galaxy Evolution" (HERITAGE; HERITAGE Team 2013; Meixner et al. 2013) survey. Source 1 A is unresolved from the neighboring source between 160 and 500 μm; thus, we use the Herschel flux densities in these bands as upper limits.
• 4.  
  The ALMA 870 μm continuum flux density based on the reprocessed archival data.
• 5.  
  The ALMA 1.2 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.

Table 2. Multiwavelength Photometry for N 105–1 A

Instruments and Bands	Flux Density (Uncertainty)	FWHM a	Source ID	Ref.
	(mJy)	('')
VMC Y (1.02), J (1.25), Ks (2.15 μm)	−999.9, −999.9, 1.78 (0.10)	≲1	558354728325	(1)
IRSF J (1.25), H (1.63), Ks (2.14 μm)	0.04 (0.01), 0.20 (0.02), 1.85 (0.07)	∼1.3	05095050-6853052	(2)
2MASS J (1.25), H (1.65), Ks (2.16 μm)	−999.9, −999.9, 2.26 (0.10)	∼2.5	J05095054−6853052	(3)
Spitzer/IRAC 3.6, 4.5, 5.8, 8.0 μm b	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)	—"—	050950.53−685305.5 c	(5)
Herschel/PACS 100, 160 μm	41,390 (3013), 23,080 (1612)	8.6, 12.6	HSOBMHERICC J77.459799−68.884809	(7)
Herschel/SPIRE 250, 350, 500 μm	10,060 (652), 4967 (309), 7710 (552)	18.3, 26.7, 40.5	—"—	(7)
ALMA 870 μm (345.798 GHz)	161.7 (0.1) d	0.47 × 0.39, 370	N 105–1 A	(9)
ALMA 1.2 mm (242.4 GHz)	71.8 (2.2) d	0.51 × 0.47, 372	—"—	(8), (9)
ALMA 3 mm (99 GHz)	48.6 (0.3) d	2.49 × 1.83, 461	—"—	(9)
ATCA 3 cm (8.6 GHz)	39 (1)	1.82 × 1.24, 148	B0510−6857 W	(10)
ATCA 6 cm (4.8 GHz)	26 (1)	2.19 × 1.70, 167	—"—	(10)

Notes.

^a The image angular resolution at FWHM. For the near-IR ground-based surveys VMC, IRSF, and Two Micron All Sky Survey (2MASS), the angular resolution depends on the atmospheric seeing. For the ALMA and ATCA observations, we provide the synthesized beam sizes and position angles: ${\theta }_{{\rm{maj}}}\times {\theta }_{{\rm{\min }}},{PA}$. ^b No MIPS photometry is available for 1 A in Whitney et al. (2008) and Gruendl & Chu (2009). ^c The 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) μm bands (the same as those used in the SAGE survey). ^d Measured within a 3σ 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 ∼50% at 242.4 GHz (Sewiło et al. 2022) and ∼96.4% at 99 GHz (see Section 4.3.1), respectively.

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).

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In addition to using the Spitzer catalog data, we have used the Spitzer IRS spectrum from Seale et al. (2009) to better constrain the SED between 5.2 and 37.9 μm. 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 μm.

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 short wavelength, low resolution (SL1; 7.6–14.6 μm), short wavelength, high resolution (SH; 9.9–19.3 μm), and long wavelength, high resolution (LH; 18.9–36.9 μm) modules by 0.9704, 0.8937, and 1.4173, respectively.

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

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We present the integrated intensity (moment 0), intensity weighted velocity (moment 1), and line-width (the 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 discussed in the present paper for the first time. All three images are shown for the CS (5–4) transition in Figure D2, for SO ³Σ 6₆–5₅ in Figure D4, and for ¹³CO (1–0) in Figure D1. 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 Figure D6. In addition, in Figures D3 and D5, we show the channel maps for CS and SO, respectively.

For the highest-resolution ¹²CO (3–2) and HCO⁺ (4–3) data with the smallest field of view, we present the moment 0 and moment 1 images in Figures D7 and D9, respectively, and channel maps in Figures D8 and D10, respectively. The kinematics of the molecular gas in N 105–1 is discussed in Section 5.3.

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