ALMA Observations of Molecular Complexity in the Large Magellanic Cloud: The N 105 Star-forming Region

Methanol (CH₃OH), methyl cyanide (CH₃CN), and larger COMs have been found in the LMC toward hot cores (Sewiło et al. 2018; Shimonishi et al. 2020): small (D ≲ 0.1 pc), hot (T_kin ≳ 100 K), and dense (n_H ≳ 10⁶⁻⁷ cm⁻³) regions around forming massive stars where ice mantles have recently been removed from dust grains as a result of thermal evaporation and/or sputtering in shock waves (e.g., Garay & Lizano 1999; Kurtz et al. 2000; Cesaroni 2005; Palau et al. 2011). A typical Galactic hot core has a very rich spectrum at submillimeter wavelengths including lines from many complex organics—the products of interstellar grain-surface chemistry or postdesorption gas chemistry (e.g., Herbst & van Dishoeck 2009; Oberg 2016; Jørgensen et al. 2020). Methanol has also been detected toward a handful of other locations in the LMC but outside hot cores ("cold methanol"; see e.g., Sewiło et al. 2019).

Observational, theoretical, and laboratory studies indicate that COMs are mainly formed on dust grains through ice chemistry in the young stellar object's (YSO's) envelope, which has an inward temperature gradient due to heating from the central protostar (e.g., Herbst & van Dishoeck 2009; van Dishoeck 2014, 2018; Oberg 2016). During the YSO accretion phase, the composition of the icy grain mantles change as they approach the protostar and eventually sublimate when the grains reach the inner hot core region in its immediate surroundings. (1) Initially, they only contain simple ices that are formed in the molecular cloud phase by a condensation of atoms and molecules from the gas phase and by subsequent grain surface chemistry (e.g., H₂O, CH₄, NH₃, CO₂, H₂CO, CH₃OH). First, the ices on the grain surface are formed though hydrogenation (adding H atoms that are the most mobile ice constituents at T ∼ 10 K), and then also through chemical reactions involving CO. (2) When exposed to UV radiation (e.g., cosmic-ray interactions with H₂), simple ices can partially dissociate into radicals. (3) The radicals become mobile when the temperature increases with decreasing distance from the protostar, and they combine to form new species, including more complex molecules. (4) As the dust grain approaches the central protostar, the temperature becomes high enough for ice mantles to sublimate (T ∼ 100–150 K). The molecules released to the gas phase include simple ices from the original ice mantles, as well as newly formed complex organics. Gas-phase chemistry following ice sublimation can also contribute to the formation of some COMs (Taquet et al. 2016).

YSOs are also associated with jets and outflows at a range of velocities, which interact with the envelope and cloud material and produce shocks, enabling shock and hot-gas chemistry. Shocks can sublimate or sputter icy grain mantles, releasing ice chemistry products (molecules such as CH₃OH and other COMs) into the gas (e.g., Arce et al. 2008). At high velocities, shocks can also sputter the grain cores, releasing the Si and S atoms and as a consequence, enhancing the production of Si- and S-bearing species such as SiO, SO₂, and SO (e.g., Schilke et al. 1997; Gusdorf et al. 2008; van Dishoeck 2018 and references therein). In addition to jets and outflows, low-velocity shocks can be produced in YSOs at the envelope-disk interface where sublimation and sputtering of ices can occur (e.g., Aota et al. 2015; Miura et al. 2017). In summary, the formation of COMs is mainly a result of the chemical processes taking place in icy grain mantles in the protostellar envelope. The ice chemistry products (including COMs) become observable after icy grain mantles are sublimated close to the protostar or sublimated and sputtered in shocks in the jets/outflows or at the envelope-disk interface. Some COMs may be the result of the gas-phase chemistry following ice sublimation.

Prior to the present study, COMs with more than six atoms had only been detected toward two hot cores in the LMC: A1 and B3 in star-forming region N 113 (N 113 A1 and N 113 B3; Sewiło et al. 2018, 2019). Sewiło et al. (2018) reported the detection of methyl formate (HCOOCH₃) and dimethyl ether (CH₃OCH₃), together with their likely parent species CH₃OH, with fractional abundances with respect to H₂ (corrected for a reduced metallicity in the LMC with respect to the Milky Way) at the lower end, but within the range measured toward Galactic hot cores.

This was a surprising result, because the previous theoretical and observational studies indicated a deficiency of CH₃OH in the LMC (e.g., Acharyya & Herbst 2015; Nishimura et al. 2016; Shimonishi et al. 2016a, 2016b). For example, Shimonishi et al. (2016b) claimed a detection of a hot core toward the massive YSO ST11 in the LMC based on the derived physical conditions and the presence of simple molecules connected to the gas chemistry (e.g., SO₂), but no CH₃OH or other COMs were detected. They concluded that CH₃OH is depleted by 2–3 orders of magnitude as compared with Galactic hot cores. The underabundance of the CH₃OH ice and a low detection rate of CH₃OH masers were also reported in the LMC (e.g., Sinclair et al. 1992; Green et al. 2008; Shimonishi et al. 2016a).

The fact that molecules whose formation requires the hydrogenation of CO on grain surfaces were not detected (e.g., CH₃OH, HNCO) or were underabundant (e.g., H₂CO) in ST11 and that the CH₃OH ice is underabundant in the LMC YSOs led Shimonishi et al. (2016a) to propose a "warm ice chemistry" model in which the observed differences between the chemistry of the LMC and Galactic sources are a consequence of the dust being warmer in the LMC due to the strong interstellar radiation field. High dust temperatures in the LMC (T ≳ 20 K) suppress the hydrogenation of CO on grain surfaces due to the decrease in available hydrogen atoms, leading to inefficient production of CH₃OH. At the same time, the model also predicts an enhancement in CO₂ production due to the increased mobility of the parent species, which explains the increased CO₂/H₂O ice column density ratio observed toward LMC YSOs (e.g., Shimonishi et al. 2008; Oliveira et al. 2009, 2011); alternatively, this increased ratio can be explained by the underabundance of H₂O (e.g., Oliveira et al. 2011). The predictions of the warm ice chemistry model are consistent with astrochemical simulations for the appropriate elemental depletions (Acharyya & Herbst 2015, 2018; Pauly & Garrod 2018). There is evidence that CH₃OH and other complex organics observed toward YSOs (hot cores and outflow shocks) might have formed in the cold (∼10 K) molecular cloud phase preceding the onset of star formation (see discussion in Section 7.2).

A picture of a chemically diverse hot core population in the LMC has recently started emerging with the detection of hot core ST16 exhibiting CH₃OH and CH₃CN emission, but no larger COMs, and a general underabundance of organic species compared with Galactic hot cores (Shimonishi et al. 2020). Shimonishi et al. (2020) suggested that LMC hot cores can be divided into "organic-poor" and "organic-rich." This classification, however, is based on only a handful of objects and needs verification.

If confirmed based on a larger sample of LMC hot cores, organic-rich hot cores would be those sources that are associated with larger COMs and have molecular abundances roughly scaled with metallicity (as in N 113 A1 and B3). In the organic-poor hot cores, the low abundances of organic molecules cannot be explained by the decreased abundance of C and O. No COMs (ST11) or only CH₃OH and CH₃CN (ST16) are detected. In ST11 and ST16, H₂CO, CH₃OH, HNCO, CS, H₂CS, and SiO are significantly less abundant, while HCO⁺, SO, SO₂, and NO are comparable with or more abundant than Galactic hot cores, after being corrected for metallicity. The organic-poor hot cores are unique to the low-metallicity environment of the LMC. Shimonishi et al. (2020) argue that a large chemical diversity of organic molecules seen in the LMC hot cores can be a consequence of the different grain temperature at the initial (ice-forming) stage of star formation. They support their conclusions with astrochemical simulations. The analysis of a larger sample of hot cores in the LMC is required to verify the hot core classification scheme suggested by Shimonishi et al. (2020) and get a better understanding of the complex chemistry in the metal-poor environment.

In this paper, we report the results of our observations of three fields in star-forming region N 105 in the LMC with the Atacama Large Millimeter/submillimeter Array (ALMA), which include a detection of two hot cores that increase a previously known very small sample of four hot cores in the LMC. The LHA 120–N 105 (hereafter N 105, Henize 1956; or DEM L86, Davies et al. 1976) nebula is the star-forming region located at the western edge of the LMC bar (e.g., Ambrocio-Cruz et al. 1998). The Hα image of N 105 reveals a bright central region (referred to in literature as N 105A) surrounded by a faint extended emission (see Figures 1 and 2). A sparse cluster NGC 1858 (e.g., Bica et al. 1996) with age estimates in a range 8–17 Myr (Vallenari et al. 1994; Alcaino & Liller 1986) and an associated OB association LH 31 (e.g., Lucke & Hodge 1970) are embedded within N 105A. LH 31 contains 18 OB stars and two Wolf–Rayet stars, and coincides with the strongest X-ray emission in the region (e.g., Vallenari et al. 1994; Dunne et al. 2001). Despite the presence of the OB association, the dense cloud N 105A shows little evidence for feedback from massive stars (e.g., Ambrocio-Cruz et al. 1998; Oliveira et al. 2006). The N 105 optical nebula is associated with the thermal radio continuum source MC 23 or B0510–6857 (e.g., McGee et al. 1972; Ellingsen et al. 1994; Filipovic et al. 1998).

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Ongoing star formation in N 105A is traced by H₂O (e.g., Scalise & Braz 1982; Whiteoak et al. 1983; Lazendic et al. 2002; Oliveira et al. 2006; Ellingsen et al. 2010), OH (e.g., Haynes & Caswell 1981; Brooks & Whiteoak 1997), and CH₃OH masers (e.g., Green et al. 2008; Ellingsen et al. 2010), ultracompact (UC) H ii regions (Indebetouw et al. 2004), and YSOs. About 40 YSOs have been identified based on the Spitzer Space Telescope (3.6–70 μm; Carlson et al. 2012 and references therein) and the Herschel Space Observatory (100–500 μm; Sewiło et al. 2010; Seale et al. 2014) data within the Hα nebula. The Spitzer images shown in Figure 1 reveal a complex structure of the dust and polycyclic aromatic hydrocarbon (PAH) emission in N 105, with the brightest emission coinciding with the position of the most massive YSOs in N 105A.

Active star-forming sites in N 105 coincide with the position of the molecular cloud detected in single-dish observations of ¹²CO and ¹³CO (1–0), tracing gas densities of ∼10²–10³ cm⁻³ (e.g., Israel et al. 1993, HPBW∼45''at 115 GHz; Chin et al. 1997, 45''; Fukui et al. 1999 and Fukui et al. 2008, 26; Wong et al. 2011, 45''—see Figure 2). High-resolution (∼67) interferometric observations of HCN and HCO⁺ (1–0) toward the peak of the CO emission with the Australia Telescope Compact Array (ATCA) revealed the densest gas in N 105 (Seale et al. 2012; see Figure 2). Two of the three fields we observed with ALMA are located in this region and are associated with H₂O and OH masers, while the third field covers a lower density region to the south and is associated with a CH₃OH maser.

The paper is organized as follows. In Section 2, we describe the observations and the archival data used in the paper. In Sections 3–5, we present the analysis of the 1.2 mm continuum and spectral line data. In Section 6, we investigate the physical characteristics of the observed fields and chemical properties of selected sources in the N 105 star-forming region based on the data ranging from the optical to radio wavelengths. The discussion is presented in Section 7, while in Section 8 we provide the summary and conclusions of our study.

We selected six fields in the LMC for Cycle 7 observations that have common characteristics with those hosting N 113 A1 and B3, at that time, the only known LMC hot cores with COMs: they are associated with massive Spitzer YSOs, H₂O/OH masers, and SO emission, a well-known hot core and shock tracer (e.g., Chernin et al. 1994; Mookerjea et al. 2007). We also observed an additional field centered on a Stage 0/I protostar (e.g., Sewiło et al. 2010) associated with one of four 6.67 GHz and the only 12.2 GHz CH₃OH maser known in the LMC (Sinclair et al. 1992), making it a good hot core candidate. In total, seven fields were observed with the ALMA 12 m Array in Band 6 (with a single pointing each) as part of the Cycle 7 project 2019.1.01720.S (PI M. Sewiło).

The SO 3₂–2₁ line emission toward N 113 A1 and B3 hot cores was serendipitously detected in our ALMA Cycle 3 observations (2015.1.01388.S, PI M. Sewiło; see also Sewiło et al. 2018). Enhanced SO emission can occur in hot cores following reactions S + OH and O + SH where the radicals and atoms are produced from the gas-phase destruction of H₂O and H₂S molecules evaporated/sputtered from ices (e.g., Charnley 1997).

A similar Band 3 correlator setup as for N 113 that covered the SO line was used in an unrelated project targeting massive YSOs in the LMC (2017.1.00093.S, PI T. Onishi), providing us with an opportunity to search for sources with a serendipitous SO detection. We have identified four Band 3 fields with SO detections and associated with masers (three with H₂O masers and one with an OH maser; e.g., Ellingsen et al. 2010; J. Ott 2022, private communication), resembling the A1 and B3 hot cores in N 113, which are also associated with masers. All the Band 3 fields were observed with the same setup, resulting in an ALMA synthesized beam of ∼213 ×157 and a channel width of 2.96 km s⁻¹. These observations also include the (1–0) transitions of ¹³CO and C¹⁸O, CS (2–1), and the 3 mm continuum; all four Cycle 5 Band 3 fields are associated with dense gas tracers (C¹⁸O and CS). Six out of seven fields included in our Cycle 7 Band 6 observations are centered on regions with SO emission and H₂O/OH masers within these four Cycle 5 Band 3 fields and thus are most likely to host hot cores.

Here, we present the results for three fields observed in Cycle 7, all located in the N 105 star-forming region. Two of the fields are associated with SO emission and H₂O/OH masers; we have dubbed them "N 105–1" and "N 105–2". The third field is associated with methanol masers, and we will refer to it as "N 105–3"; no prior ALMA observations are available for N 105–3. All three ALMA fields in N 105 hosting hot core candidates are shown in Figure 1, and their positions are listed in Table 2.

The observations of all fields 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. The targets were observed again on 2019 October 23, with 43 antennas and baselines from 15 to 782 m. The calibrators were the same as for the first run on October 21. The total on-source integration time for all seven fields was 91.8 minutes. for all three executions. The maximum recoverable scale calculated from the 5th percentile baseline length for the final data set combining all executions varied between 56 and 52 for a sky frequency range covered by our observations (∼241.3–260.4 GHz). The spectral setup included four 1875 MHz spectral windows centered on frequencies of 242.4, 244.8, 257.85, and 259.7 GHz, each with 3840 channels, providing a spectral resolution of 1.21–1.13 km s⁻¹. Henceforth, we will refer to the spectral windows as the "242 GHz/245 GHz/258 GHz/260 GHz spectral window."

The data were calibrated and imaged with version 5.6.1–8 of the ALMA pipeline in Common Astronomy Software Applications (CASA; McMullin et al. 2007). The continuum in each spectral window was identified and subtracted before cube imaging. The CASA task tclean was used for imaging using the Hogbom deconvolver, standard gridder, Briggs weighting with a robust parameter of 0.5, and masking using the "auto-multithresh" algorithm. The spectral cubes have a cell size of 0092 × 0092 × 0.56 km s⁻¹. Additional information on the data cubes is included in Table 1. The 242.4 GHz (1.2 mm) continuum image parameters are listed in Table 2. All the images have been corrected for primary beam attenuation.

Three near-IR sources in the ALMA Cycle 7 fields in N 105 were observed with the VLT/KMOS (Sharples et al. 2013) as part of a survey of YSO candidates under program 0101.C-0856(A) (PI J. L. Ward) using the H + K grating with a spectral resolving power of 2000 and a spatial pixel scale of 02. The observations took place on the night of 2018 August 28–29, with seeing ranging from 055 to 166. The measured FWHM of one of the sources in N 105 (ID 558354728325) in the K band is 8.2 pixels, corresponding to approximately 16. KMOS is able to perform integral field spectroscopy in the near-IR bands for 24 targets simultaneously using 24 configurable arms. The KMOS observations were carried out using a standard nod-to-sky procedure with an integration time of 150 s, four detector integration times (DITs) and three dither positions, yielding a total on-source integration time of 1800 s. Telluric absorption correction, response curve correction, and absolute flux calibration were carried out using observations of telluric standard stars using three integral-field units (IFUs). The data were reduced with the standard VLT/KMOS pipeline using the esoreflex data reduction package (Davies et al.2013).

The K-band continuum image is produced by integrating over a third-order polynomial fit to the data for every spatial pixel (spaxel) over the spectral range 2.028–2.290 μm. The Brγ and H₂ line emission images are produced by fitting a Gaussian profile to the emission lines at every position in the image. Each IFU has a square field of view of 28 × 28.

The measured KMOS H- and K-band fluxes are found to be significantly lower than those determined by the near-IR surveys covering this region, the Two Micron All Sky Survey (2MASS) JHK_S , and the Visible and Infrared Survey Telescope for Astronomy (VISTA) YJK_S survey of the Magellanic Clouds system (VMC; Cioni et al. 2011), on average by a factor of 24. Thus, the KMOS fluxes are not reliable enough to be used directly. Instead, for the subsequent analysis, we have scaled the extracted spectra so that the sum of the spectral region from 2.028–2.295 μm is consistent with the K-band magnitude of the corresponding point source from the VMC survey catalog. The publicly available VMC catalog was queried using the VISTA Science Archive (VSA; ¹⁹ Cross et al. 2012) to obtain aperture photometry in YJK_S (VMCDR4). For reference, N 105 is located in VMC tile LMC 6_4.

To improve astrometry of the KMOS images, we computed the cross-correlation functions for all the KMOS fields with the VMC survey and calculated the R.A. and decl. values that the KMOS data should be shifted by to match the VMC data. First, the KMOS K-band continuum images were flipped, rotated by 4.9918 degrees, and rescaled to match the orientation and pixel scale of the VMC data using the WCSTools package. The CORREL_IMAGES function in IDL was then used to compute the 2D cross-correlation function between the KMOS and VMC images. The 2D Gaussian profiles were fitted to the cross-correlation functions giving the most probable offsets (the centroid position of the Gaussian) in the VMC survey coordinate frame. Flipping and rotating the offsets between the KMOS and VMC images then converts them from the VMC coordinate system into R.A. and decl. There is still a small shift between the KMOS and VMC data; however, this shift is subpixel and thus not significant, and the association between the VMC and KMOS sources can be established reliably. The precision of the KMOS astrometry is limited by the relatively poor resolution of the KMOS data due to seeing. The results of the VLT/KMOS observations are discussed in Section 6.2.

Each of our ALMA fields contains high-mass YSO candidates identified based on the Spitzer 3.6–70 μm data from the LMC-wide Spitzer Surveying the Agents of Galaxy Evolution (SAGE; Meixner et al. 2006; SAGE Team 2006) survey (e.g., Whitney et al. 2008; Gruendl & Chu 2009; Carlson et al. 2012). Spitzer is mostly sensitive to Stage I YSOs with disks and envelopes and some more evolved Stage II YSOs with disks and remnant or no envelopes. Subsets of YSO candidates were followed up with near- to far-IR spectroscopic observations, which confirmed their nature and allowed for investigating their physical and chemical characteristics (e.g., Oliveira et al. 2009; Seale et al. 2009; Ward et al. 2016; Jones et al. 2017; and Oliveira et al. 2019).

Four out of six YSO candidates in the ALMA fields of view in N 105 were confirmed spectroscopically as bona fide YSOs by Seale et al. (2009) using the Spitzer Infrared Spectrograph (IRS) observations (5–37 μm). Two sources were classified as "Group P" and another two as "Group PE" YSOs. Both Group P and PE sources show strong PAH emission features. More evolved Group PE sources also show 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. The sources from both groups may show some absorption from silicates, particularly at 10 μm; the silicate absorption features are difficult to identify unambiguously in the presence of strong PAH emission features at 6.2, 7.7, 8.6, and 11.3 μm. Group P and PE sources can also exhibit the CO₂ 15.2 μm ice absorption feature in the Spitzer/IRS spectra (Seale et al. 2011; see also Section 6.1).

Below we provide a more detailed discussion on YSO candidates, spectroscopically confirmed YSOs, and masers (H₂O, OH, CH₃OH) in individual ALMA fields in N 105.

N 105–1: N 105–1 hosts two spectroscopically confirmed YSOs from Seale et al. (2009): 050950.53−685305.5 (source #318 or SSTISAGE1C J050950.53−685305.4 from Whitney et al. 2008) and 050952.73−685300.7 (see Figure 4). Source 050950.53−685305.5 is associated with the bright 1.2 mm continuum source N 105–1 A, while 050952.73−685300.7 coincides with N 105–1 B and an extended emission to the east (see Section 6). YSO 050950.53−685305.5 was classified by Seale et al. (2009) as a Group P and 050952.73−685300.7 as a Group PE source.

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No maser detection has been reported in literature toward N 105–1.

N 105–2: Three Spitzer YSO candidates from Gruendl & Chu (2009) are in the N 105–2 field (050952.26−685327.3, 050953.89−685336.7, and 050951.31−685335.6; Figure 4), two of which were spectroscopically confirmed as YSOs by Seale et al. (2009) and are associated with the 1.2 mm continuum emission. The YSO 050953.89−685336.7 (Group P source in Seale et al. 2009) coincides with N 105–2 E, while the Gruendl & Chu (2009)'s catalog position of 050952.26−685327.3 (Group PE) lies between the 2 A and 2 B continuum peaks. The inspection of the Spitzer images shows that no source in the image is visible at this position, but there are two Spitzer sources in the vicinity. The Spitzer Infrared Array Camera (IRAC) resolution is just about resolving these two sources corresponding to the 1.2 mm continuum peaks 2 A and 2 B, separated by only ∼1.5 pixels. It is likely that the source-finding routine used by Gruendl & Chu (2009) found only 2 A (the brighter peak at 4.5 μm) but not 2 B in the shorter IRAC bands, and only 2 B (the brighter peak at 8.0 μm) but not 2 A in the longer IRAC bands. The sources are close enough to be identified as a single object during the band-merging process, resulting in a catalog photometry and position (a weighted mean of the positions found in individual bands) being a combination of these two sources. The catalog position roughly in between the two Spitzer sources supports this interpretation. The Spitzer/IRS spectrum of 050952.26−685327.3 analyzed by Seale et al. (2009) most likely includes contributions from both nearby sources as well. The SAGE IRAC point-source catalog does not include any sources in the central part of the ALMA field.

No 1.2 mm continuum emission has been detected toward the position of the YSO candidate 050951.31−685335.6 in N 105–2 located to the southwest from 2 A and 2 B (Figure 4).

Two sources in N 105–2 are associated with masers. N 105–2 A and 2 B coincide with the 22 GHz H₂O masers (Whiteoak et al. 1983; Whiteoak & Gardner 1986; Lazendic et al. 2002; Ellingsen et al. 2010; J. Ott 2022, private communication; see also Schwarz et al. 2012). Source 2 A is also associated with the 1665/1667 MHz OH maser (Haynes & Caswell 1981; Gardner & Whiteoak 1985; Brooks & Whiteoak1997). No methanol masers have been detected toward N 105–2 (Green et al. 2008 and Ellingsen et al. 2010).

The maser positions used for investigating correlations with the ALMA and infrared emission (e.g., in Figure 4) come from Green et al. (2008) who summarize previous observations of different types of masers and provide accurate positions (within subarcsecond) obtained using the interferometric observations with ATCA.

One H₂O maser spot in N 105–2 detected with ATCA is offset toward the southeast from the 1.2 mm continuum source 2 B (J. Ott 2022, private communication; Schwarz et al. 2012; see e.g., Figure 4). Imai et al. (2013) reported a detection of another H₂O maser spot at a distance of ∼34 from the N 105–2 B 1.2 mm continuum peak toward north—northeast. However, they incorrectly associate this source with that reported in Oliveira et al. (2006) who was not able to estimate an accurate position of the maser spot based on their Parkes 64–m telescope observations, but argued that it is likely related to the maser detected by Lazendic et al. (2002); the accurate position of the H₂O maser provided by Lazendic et al. (2002) indicates that the maser emission originates in the vicinity of N 105–2 A. Due to this positional uncertainty, we do not show the position of the H₂O maser from Imai et al. (2013), which is not reported in other surveys, in the images.

N 105–3: One Spitzer YSO candidate lies within the N105 –3 field (050958.52−685435.5, Gruendl & Chu 2009; Carlson et al. 2012; see Figure 4), with the Spitzer catalog position corresponding to the 1.2 mm continuum peak of source N 105–3 A. No follow-up spectroscopic observations exist for 050958.52−685435.5.

In N 105–3, the 1.2 mm continuum source 3 B is associated with CH₃OH masers: 6.7 and 12.2 GHz (Green et al. 2008; Ellingsen et al. 2010). No H₂O masers have been detected toward this field (e.g., Ellingsen et al. 2010).

Three of the YSOs in N 105 have been well fit with the Robitaille et al. (2006) YSO radiation transfer models by Carlson et al. (2012): 050950.53−685305.5 (source SSTISAGEMAJ050950.53−685305.4 in Carlson et al. 2012; 1 A), 050953.89−685336.7 (SSTISAGEMA J050953.91−685337.1; 2 E), and 050958.52−685435.5 (SSTISAGEMA J050958.52−685435.2; 3 A). All sources were found to be massive with stellar masses and luminosities for the best-fit YSO models of (31.3 ± 2.6, 23.0 ± 3.2, 17.9 ± 1.4) M_⊙ and (14 ± 2, 6.6 ± 2.2, 3.6 ± 0.5) × 10⁴ L_⊙ for (050950.53−685305.5, 050953.89−685336.7, 050958.52−685435.5).

The initial spectral line identification was carried out in the CASA task Viewer, which uses the NRAO's spectral line database Splatalogue. ²⁰ The line identification was later verified by comparing the spectra to the predictions of models for the subset of molecules detected in Galactic hot cores, assuming local thermodynamic equilibrium (LTE) conditions. All the detectable lines predicted by the spectral model in the observed frequency ranges must be present in the observed spectrum with relative intensities for different transitions approximately consistent with the model predictions. The spectral analysis and modeling results are described in detail in Section 4.3. In our analysis, we use molecular data from the Cologne Database for Molecular Spectroscopy (CDMS ²¹ ; Müller et al. 2005) where available; otherwise, we use the Jet Propulsion Laboratory (JPL) Millimeter and Submillimeter Spectral Line Catalog ²² (Pickett et al. 1998; see Section 4.3 for details). We use the CDMS quantum number notation from the Splatalogue throughout the paper.

Table 3 lists all the molecular lines detected toward the continuum sources in N 105. We detected S-bearing species: SO, ³³SO, SO₂, ³⁴SO₂, CS, C³³S, OCS, H₂CS; N-bearing species: HNCO, HC₃N, HC¹⁵N, H¹³CN; three deuterated molecules: HDO, HDCO, and HDS; as well as SiO, H¹³CO⁺, and CH₂CO.

Table 3. Spectral Lines Detected toward Continuum Sources in N 105–1, N 105–2, and N 105–3 ^a

Species	Transition	Frequency	EU	1 A	1 B	1 C	2 A	2 B	2 C	2 D	2 E	2 F	3 A	3 B	3 C
		(MHz)	(K)
COMs
CH3OH	5−0,5–4−0,4 E, vt = 0	241,700.159	47.94	✓?	✓	✓?	✓	✓	✓	✓	✓?	✓	✓?	✓ ? d	✓?
CH3OH	51,5–41,4 E, vt = 0	241,767.234	40.39	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓ ? d	✓
CH3OH	50,5–40,4 A, vt = 0	241,791.352	34.82	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓ ? d	✓
CH3OH	54,2–44,1 A, vt = 0	241,806.524	115.17	−	−	−	✓	✓	✓?	−	−	✓?	−	✓ ? d	−
CH3OH	54,1–44,0 A, vt = 0	241,806.525	115.17	−	−	−	✓	✓	✓?	−	−	✓?	−	✓ ? d	−
CH3OH	54,2–44,1 E, vt = 0	241,813.255	122.73	−	−	−	✓	✓	✓?	−	−	✓?	−	✓ ? d	−
CH3OH	5−4,1–4−4,0 E, vt = 0	241,829.629	130.82	−	−	−	✓	✓	✓	−	−	−	−	✓ ? d	−
CH3OH	53,3–43,2 A, vt = 0	241,832.718	84.62	−	−	−	✓	✓	✓	✓?	−	✓?	−	✓ ? d	−
CH3OH	53,2–43,1 A, vt = 0	241,833.106	84.62	−	✓	−	✓	✓	✓	✓?	−	✓?	−	✓ ? d	−
CH3OH	52,4–42,3 A, vt = 0	241,842.284	72.53	−	✓	✓?	✓	✓	✓	✓	−	✓?	−	✓ ? d	−
CH3OH	5−3,3–4−3,2 E, vt = 0	241,843.604	82.53	−	−	−	✓	✓	✓	✓	−	✓?	−	✓ ? d	−
CH3OH	53,2–43,1 E, vt = 0	241,852.299	97.53	−	−	−	✓	✓	−	−	−	✓?	−	✓ ? d	−
CH3OH	5−1,4–4−1,3 E, vt = 0	241,879.025	55.87	✓? b	✓	✓?	✓	✓	✓	✓	✓?	✓	✓?	✓ ? d	−
CH3OH	52,3–42,2 A, vt = 0	241,887.674	72.54	−	−	✓?	✓	✓	✓	✓?	−	−	−	✓ ? d	−
CH3OH	52,3–42,2 E, vt = 0	241,904.147	60.73	✓? b	✓	✓	✓	✓	✓	✓	✓?	✓	−	✓ ? d	−
CH3OH	5−2,4–4−2,3 E, vt = 0	241,904.643	57.07	✓? b	✓	✓	✓	✓	✓	✓	✓?	✓	−	✓ ? d	−
CH3OH	141,14–132,11 E, vt = 0	242,446.084	248.94	−	−	−	✓	✓?	−	−	−	−	−	−	−
CH3OH	51,4–41,3 A, vt = 0	243,915.788	49.66	−	✓	−	✓	✓	✓	✓	−	✓	−	−	−
CH3OH	223,19–222,20 A, vt = 0	244,330.372	636.78	✓?	−	−	−	−	−	−	−	−	−	−	−
CH3OH	9−1,9–8−0,8 E, vt = 1	244,337.983	395.66	✓?	−	−	✓?	−	−	−	−	−	−	−	−
CH3OH	213,18–212,19 A, vt = 0	245,223.019	585.76	−	−	−	✓	−	−	−	−	−	−	−	−
CH3OH	183,16–182,17 A, vt = 0	257,402.086	446.55	−	−	−	✓	✓? c	−	−	−	−	−	−	−
CH3OH	193,17–192,18 A, vt = 0	258,780.248	490.58	−	−	−	✓	−	−	−	−	−	−	−	−
CH3OH	203,18–202,19 A, vt = 0	260,381.463	536.97	−	−	−	✓	−	−	−	−	−	−	−	−
CH3OCH3	131,13–120,12 EA	241,946.249	81.13	−	−	−	✓?	−	−	−	−	−	−	−	−
CH3OCH3	131,13–120,12 AE	241,946.249	81.13	−	−	−	✓?	−	−	−	−	−	−	−	−
CH3OCH3	131,13–120,12 EE	241,946.542	81.13	−	−	−	✓?	−	−	−	−	−	−	−	−
CH3OCH3	131,13–120,12 AA	241,946.835	81.13	−	−	−	✓?	−	−	−	−	−	−	−	−
CH3OCH3	141,14–130,13 EA	258,548.819	93.33	−	−	−	✓	✓	−	−	−	−	−	−	−
CH3OCH3	141,14–130,13 AE	258,548.819	93.33	−	−	−	✓	✓	−	−	−	−	−	−	−
CH3OCH3	141,14–130,13 EE	258,549.063	93.33	−	−	−	✓	✓	−	−	−	−	−	−	−
CH3OCH3	141,14–130,13 AA	258,549.308	93.33	−	−	−	✓	✓	−	−	−	−	−	−	−
CH3CN	146–136	257,349.180	349.73	−	−	−	✓	− e	−	−	−	−	−	−	−
CH3CN	145–135	257,403.585	271.23	−	−	−	✓	✓? c	−	−	−	−	−	−	−
CH3CN	144–134	257,448.128	206.98	−	−	−	✓	✓?	−	−	−	−	−	−	−
CH3CN	143–133	257,482.792	156.99	−	−	−	✓	✓	−	−	−	−	−	−	−
CH3CN	142–132	257,507.562	121.28	−	−	−	✓	✓	✓?	−	−	−	−	−	−
CH3CN	141–131	257,522.428	99.85	−	−	−	✓	✓	✓?	−	−	−	−	−	−
CH3CN	140–130	257,527.384	92.71	−	−	−	✓	✓	✓	−	−	−	−	−	−
NH2CHO	122,10–112,9	260,189.090	92.36	−	−	−	✓	−	−	−	−	−	−	−	−
Other Molecules
HNCO	110,11–100,10	241,774.032	69.63	−	✓	−	✓	✓	✓	−	−	−	−	−	−
HNCO	111,10–101,9	242,639.705	113.15	−	−	−	✓	✓	−	−	−	−	−	−	−
HC3N	27–26	245,606.320	165.04	−	−	−	✓	✓	−	−	−	−	−	−	−
HC15N	3–2	258,156.996	24.78	−	✓	✓	✓	✓	✓	✓	−	✓?	−	−	−
H13CN	3–2	259,011.798	24.86	✓	✓	✓	✓	✓	✓	✓	−	✓	−	−	−
H13CO+	3–2	260,255.339	24.98	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓?
CH2CO	121,11–111,10	244,712.269	89.40	−	−	−	✓	−	−	−	−	−	−	−	−
CH2CO	131,13–121,12	260,191.982	100.47	−	−	−	✓	−	−	−	−	−	−	−	−
SO 3Σ	66–55	258,255.826	56.50	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓
OCS	20–19	243,218.036	122.58	−	−	−	✓	✓	✓	−	−	−	−	−	−
H2CS	71,6–61,5	244,048.504	60.03	✓?	✓	✓	✓	✓	✓	✓	−	✓	−	−	−
CS	5–4	244,935.557	35.27	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓	✓
C33S	5–4	242,913.610	34.98	−	✓	✓?	✓	✓	✓	✓	−	✓	−	−	−
SO2	52,4–41,3	241,615.797	23.59	✓?	✓	−	✓	✓	✓	−	−	✓	✓	−	−
SO2	54,2–63,3	243,087.647	53.07	−	−	−	✓	✓	−	−	−	−	−	−	−
SO2	268,18–277,21	243,245.435	479.58	−	−	−	−	✓	−	−	−	−	−	−	−
SO2 f	140,14–131,13	244,254.218	93.90	✓	✓?	−	✓	✓	✓	−	−	✓	−	−	−
SO2	263,23–254,22	245,339.233	350.79	−	−	−	✓	✓	−	−	−	−	−	−	−
SO2 f	103,7–102,8	245,563.422	72.72	✓	✓	−	✓	✓	✓?	−	−	−	✓	−	−
SO2	73,5–72,6	257,099.966	47.84	✓	✓?	−	✓	✓	✓	−	−	−	−	−	−
SO2	324,28–323,29	258,388.716	531.12	−	−	−	✓	✓	−	−	−	−	−	−	−
SO2	207,13–216,16	258,666.969	313.19	−	−	−	✓	✓	−	−	−	−	−	−	−
SO2 f	93,7–92,8	258,942.199	63.47	✓	✓	−	✓	✓	−	−	−	−	−	−	−
SO2	304,26–303,27	259,599.448	471.52	−	−	−	✓	✓	−	−	−	−	−	−	−
34SO2	161,15–152,14	241,509.046	130.31	−	−	−	✓	✓	−	−	−	−	−	−	−
34SO2	83,5–82,6	241,985.449	54.38	−	−	−	✓	✓	−	−	−	−	−	−	−
34SO2	181,17–180,18	243,936.052	162.59	−	−	−	−	✓	−	−	−	−	−	−	−
34SO2	140,14–131,13	244,481.517	93.54	−	−	−	✓	✓	−	−	−	✓?	−	−	−
34SO2	152,14–151,15	245,178.587	118.72	−	−	−	✓	✓?	−	−	−	−	−	−	−
34SO2	63,3–62,4	245,302.239	40.66	−	−	−	✓	✓?	−	−	−	✓	−	−	−
34SO2	133,11–132,12	259,617.203	104.91	−	−	−	✓	✓	−	−	−	−	−	−	−
33SO	67,6–56,5	259,280.331	47.12	−	−	−	✓	✓	−	−	−	−	−	−	−
33SO	67,7–56,6	259,282.276	47.12	−	−	−	✓	✓	−	−	−	−	−	−	−
33SO	67,8–56,7	259,284.027	47.12	−	−	−	✓	✓	−	−	−	−	−	−	−
33SO	67,9–56,8	259,284.027	47.12	−	−	−	✓	✓	−	−	−	−	−	−	−
SiO	6–5	260,518.009	43.76	−	✓?	−	✓	✓	✓	✓	−	✓	−	−	−
HDO	21,1–21,2	241,561.550	95.22	−	−	−	✓	✓	−	−	−	−	−	−	−
HDCO	42,3–32,2	257,748.701	62.78	−	−	−	✓?	−	−	−	−	−	−	−	−
HDCO	42,2–32,1	259,034.910	62.87	−	−	−	✓	−	−	−	−	−	−	−	−
HDS	10,1–00,0	244,555.580	11.74	−	−	−	−	−	✓	−	−	−	−	−	−
HDS	21,1–20,2	257,781.410	47.04	−	−	−	−	−	✓?	−	−	−	−	−	−

Notes.

^a Spectroscopic parameters were taken from the CDMS catalog for all species except HDO, for which the data were taken from the JPL database (see Section 4.3). The symbols "✓", "✓?", and "−" indicate, respectively, a detection, a tentative detection, and a nondetection of a given transition. ^b Methanol transitions likely blended with the H49 recombination line toward 1 A. ^c The 257,402.086 MHz CH₃OH and 257,403.5848 MHz CH₃CN transitions are blended in the spectrum of 2 B. ^d In the spectrum of 3 B, the J = 5 − 4 CH₃OH Q-branch as a whole is securely detected, even though the individual lines are not clearly visible. ^e Even though this is a <3σ detection for 2 B, this CH₃CN transition is used in the rotational diagram analysis as it still contributes useful information (see Section 4.3 and Figure 16). ^f The SO₂ transitions likely suffering from significant opacity effects (as defined by E_U < 100 K and $\mathrm{log}(S{\mu }^{2})\gt 1.0$) and thus excluded from the fitting in the rotational diagram for 2 A and 2 B (see Section 4.3 and Figure 16).

Download table as:  ASCIITypeset images: 1 2

We detected COMs CH₃OH, CH₃CN, and CH₃OCH₃ in N 105. All three COMs are observed toward sources 2 A and 2 B. CH₃CN is also detected toward 2 C. CH₃OH is identified in the spectra of all the continuum sources, making it the most widespread COM in our observations.

We also report a tentative detection of formamide (NH₂CHO) toward N 105–2 A. We detected a single NH₂CHO transition (260.18984820 GHz): this is a ∼3.2σ detection of the brightest NH₂CHO transition within the frequency range covered by our ALMA observations. The NH₂CHO line is blended with a ketene (CH₂CO) line (260.19198200 GHz); the lines are separated by ∼2.13 MHz, which corresponds to ∼2.45 km s⁻¹ or 2.19 channel widths. The significance of this detection in the low-metallicity environment is discussed in Section 7.4.

Extragalactic detection of deuterated species were first reported in the LMC star-forming regions by Chin et al. (1996) who detected DCO⁺ toward three (N 113, N 44 BC, N 159 HW) and DCN toward one star-forming region (N 113; see also Heikkilä et al. 1997 for N 159, and Sewiło et al. 2018 and Wang et al. 2009 for N 113). Martín et al. (2006) reported a tentative detection of DNC and N₂D⁺ in the nucleus of the starburst galaxy NGC 253. Most recently, Muller et al. (2020) reported the detection of ND, NH₂D, and HDO with ALMA at redshift z = 0.89 in the spiral galaxy intercepting the line of sight to quasar PKS 1830−211. We detected deuterated formaldehyde (HDCO), deuterated hydrogen sulfide (HDS), and deuterated water (HDO) toward hot cores 2 A (HDCO and HDO) and 2 B (HDO), and a hot core candidate 2 C (HDS). These are the first extragalactic detections of HDCO and HDS, and the first detection of HDO in an extragalactic star-forming region. A detailed discussion on the detection of HDO in the LMC will be included in a separate paper. Our observations did not cover any H₂CO, H₂O, or H₂S transitions, preventing us from calculating the deuterium fractionation (the abundance ratio of deuterated over hydrogenated isotopologues, D/H) toward N 105–2.

Several hydrogen recombination lines are observed toward source 1 A: H49 (241.86116 GHz), H54η (243.94239 GHz), H53η (257.19399 GHz), H41γ (257.63549 GHz), and H36β (260.03278 GHz). This is the first extragalactic detection of the , η, and γ transitions of the hydrogen recombination lines and will be reported elsewhere.

Figures 7–15 show the integrated intensity (moment 0) images for molecular species detected toward N 105 with ALMA. Methanol is detected toward all the continuum sources, with the faintest emission associated with those in N 105–3. Other species detected toward all the sources are SO, CS, and H¹³CO⁺.

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4.2.1. N 105–1

4.2.2. N 105–2

4.2.3. N 105–3

We performed an initial assessment of the physical conditions in N 105 by using a rotational diagram analysis for CH₃CN, CH₃OH, and SO₂ for sources with multiple CH₃CN, CH₃OH, and SO₂ line detections with a range of upper state energies (E_U). This analysis assumes the gas is in LTE and the lines are optically thin (Goldsmith & Langer 1999) and not blended with lines from other species. The rotational diagrams are shown in Figure 16 for the most chemically rich sources, 2 A and 2 B. For both sources, two temperature components for CH₃OH are clearly visible in the rotational diagrams, while the SO₂ rotational diagrams shows possible non-LTE effects, i.e., an apparent discontinuity in the distribution of the low- and high-E_U data points. For the fitting in the rotational diagram, we excluded the SO₂ lines likely suffering from significant opacity effects, as defined by E_U < 100 K and $\mathrm{log}(S{\mu }^{2})\gt 1.0$ (also see Shimonishi et al. 2021).

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Spectral line modeling was then performed for all the continuum sources under the assumption of LTE and taking into account line blending and opacity effects, using a least-squares approach similar to Sewiło et al. (2018) to simultaneously retrieve best-fitting column densities, rotational temperatures, Doppler shifts, and spectral line widths for each species (i): $[{N}^{i},{T}_{\mathrm{rot}}^{i},{{v}}_{i},{{d}{v}}^{i}$]. Simultaneously modeling all lines of every detectable species in our observed frequency range is crucial in order to resolve line blending issues and optical depth effects that might otherwise bias the retrieved parameters.

Spectral line models were generated for each source using a custom Python routine (based on the code used by Cordiner et al. 2017). Spectroscopic parameters were taken from the Cologne Database for Molecular Spectroscopy (CDMS; Müller et al. 2001), where available, and additional data for HDO were taken from the JPL Molecular Spectroscopy Database (Pickett et al. 1998). Gaussian spectral line opacity profiles were assumed, and the source was assumed to fill the aperture (unity beam filling factor). The model sums the radiative source terms (in the equation of radiative transfer) in each spectral channel for emission from overlapping lines of both the same and different species. The peak opacity of each spectral line was calculated using equation A2 of Turner (1991), and the final, synthetic T_B spectrum was generated by combining the emission from the full set of lines in our frequency range based on their individual contributions to the radiation source function.

Optimization of the individual $[{N}^{i},{T}_{\mathrm{rot}}^{i},{{v}}_{i},{{d}{v}}^{i}]$ parameters for each species (i) was performed using the LMFIT nonlinear least-squares package (Newville et al. 2014). The goodness of fit between the observed and synthetic spectra was monitored via the reduced chi-square statistic (${\chi }_{{\rm{R}}}^{2}$). A good fit to the observed spectra (${\chi }_{{\rm{R}}}^{2}\approx 1.0$) was obtained for all species in all sources using a single set of $[{N}^{i},{T}_{\mathrm{rot}}^{i},{{v}}_{i},{{d}{v}}^{i}]$ parameters (cloud components) for each species, apart from CH₃OH toward N105–2 A–2 D and 2 F and SO₂ for 2 A, 2 B, and 2 F, which required two cloud temperature components in order to obtain ${\chi }_{{\rm{R}}}^{2}\approx 1.0$ (the two components are hereafter designated as "hot" and "cold" due to their significantly different best-fitting temperatures).

Multiple lines of CH₃OH were detected in each source, with differing upper-state energies (in the range from 35 K up to ∼600 K for 2 A; see Table 3), enabling robust derivations of the CH₃OH rotational temperatures. Toward sources 2 A and 2 B, multiple lines of the hot core tracer CH₃CN were also detected, resulting in a more reliable estimate of the hot core gas temperature in those sources, as the hot core CH₃OH lines are more likely to be contaminated with the ambient interstellar medium (ISM) because CH₃OH is more widespread. Temperature information was also available in some cases for SO₂ and ³⁴SO₂. For the remaining species with single-line detections (H¹³CO⁺, HCN, HC¹⁵N, HC₃N, CS, C³³S, H₂CS, OCS, SO, ³³SO, SiO, NH₂CHO, HNCO, HDCO, CH₂CO, HDO, HDS), or multiline detections of insufficient strength for robust temperature determinations (e.g., CH₃OCH₃), we fixed their rotational temperatures to the best-fitting CH₃CN rotational temperature. When CH₃CN was not detected, the other molecules were assumed to follow the temperature of the (hot) CH₃OH component. In addition, when multiple lines of SO₂ were detected, the temperature of this species was obtained independently, and the temperatures of SO and ³³SO were tied to SO₂, due to the chemical similarities between these species.

Error estimates on each free parameter were generated via Monte Carlo noise resampling. This involved the generation of 300 synthetic spectra for each source, obtained by adding random (Gaussian) noise to the best-fitting model spectra (with rms equivalent to the noise in nearby line-free spectral regions), which were subsequently refit to determine the distribution of possible model parameters. 1σ errors were determined from the ±68% ranges of the resulting parameter distributions, under the assumption of Gaussian statistics.

In the spectral fitting process, we used the CDMS/JPL partition functions (Q(T_rot)); the CDMS data were used where available, i.e., for all species except HDO. In our experience, the CDMS catalog tends to have more complete partition functions than the JPL catalog, including information from higher-excitation states where available. In each case, we used the appropriate corresponding partition function from the respective catalog for each species, thus ensuring a consistent statistical weight scheme to that used in that catalog. Partition functions were tabulated at 0, 9.375, 18.75, 37.5, 75, 150, 225, and 300 K, and interpolated using a cubic spline.

For some molecules with noisy or tentative line detections, reliable parameter error estimates could not be obtained using the Monte Carlo resampling method due to the tendency of the radial velocity parameter to drift into spectral regions affected by emission from nearby species, or with zero emission, leading to erroneous, or insufficient constraints on the model. In these cases, the radial velocities were held fixed at the value given by the initial least-squares fit, with the other parameters allowed to vary freely.

Our spectral modeling procedure implicitly accounts for line opacity effects in the derivation of molecular column densities and rotational temperatures. In general, the deconvolved source sizes are larger or similar in size to the ALMA beam size (see Table 4). However, there is a possibility for additional, unresolved, high-opacity interstellar cloud components within the beam of our ALMA observations, that could not be distinguished at the resolution and signal-to-noise ratio of our data. In that case, the spectral line opacities could have been underestimated, leading to a corresponding underestimate of the column densities.

Table 4. Continuum Intensities and H₂ Column Densities

Source	R.A. (J2000) a	Decl. (J2000) a	Ipeak b	Imean c	Area c	FWHMeff	FWHMeff,deconv	F50 e	M50,gas e	N(H2) f	N(H2) g	N(H2) h
	(h m s)	(° ' '')	(mJy beam−1)	(mJy beam−1)	(arcsec2)	(''/pc)	(''/pc)	(mJy)	(M⊙)	(1023 cm−2)
N 105–1 A	05:09:50.54	−68:53:05.4	31.003 d	22.370 d	0.305	0.62/0.15	0.39/0.10	25.2 d	${3706}_{-475}^{+597}$	${93}_{-12}^{+15}$	⋯	${7.4}_{-1.8}^{+1.8}$
1 B	05:09:52.48	−68:53:00.7	2.329	1.575	0.923	1.08/0.26	0.97/0.23	5.4	${1038}_{-108}^{+109}$	${8.6}_{-1.0}^{+1.0}$	⋯	${4.8}_{-1.0}^{+1.0}$
1 C	05:09:52.82	−68:53:04.3	1.437	1.008	0.601	0.87/0.21	0.73/0.18	2.2	${390}_{-87}^{+95}$	${5.0}_{-1.2}^{+1.2}$	⋯	⋯
N 105–2 A	05:09:51.96	−68:53:28.3	6.362	4.560	0.423	0.73/0.18	0.55/0.13	7.1	${102}_{-13}^{+12}$	${1.6}_{-0.2}^{+0.3}$	${1.8}_{-0.2}^{+0.2}$	${1.7}_{-0.2}^{+0.2}$
2 B	05:09:52.56	−68:53:28.1	6.181	4.382	0.415	0.73/0.18	0.54/0.13	6.7	${170}_{-26}^{+27}$	${2.0}_{-0.2}^{+0.2}$	${3.1}_{-0.5}^{+0.5}$	${1.7}_{-0.2}^{+0.2}$
2 C	05:09:52.22	−68:53:22.6	1.799	1.217	0.474	0.78/0.19	0.60/0.15	2.1	${50}_{-12}^{+16}$	${0.8}_{-0.2}^{+0.3}$	⋯	${2.7}_{-0.4}^{+0.5}$
2 D	05:09:52.99	−68:53:31.0	1.344	0.910	0.592	0.87/0.21	0.72/0.17	2.0	${161}_{-18}^{+19}$	${2.1}_{-0.3}^{+0.3}$	⋯	⋯
2 E	05:09:53.90	−68:53:37.6	1.269	0.896	0.618	0.89/0.22	0.74/0.18	2.0	${501}_{-79}^{+95}$	${6.2}_{-1.0}^{+1.2}$	⋯	⋯
2 F	05:09:52.39	−68:53:28.1	2.397	⋯	⋯	⋯	⋯	⋯	⋯	${1.2}_{-0.3}^{+0.4}$	⋯	⋯
N 105–3 A	05:09:58.48	−68:54:35.4	0.573	0.409	0.559	0.84/0.20	0.69/0.17	0.85	${307}_{-57}^{+69}$	${4.2}_{-0.8}^{+1.0}$	⋯	${1.0}_{-0.3}^{+0.3}$
3 B	05:09:58.70	−68:54:34.1	0.342	0.241	0.381	0.70/0.17	0.50/0.12	0.33	${5}_{-2}^{+1}$	${0.09}_{-0.04}^{+0.03}$	⋯	⋯
3 C	05:09:58.41	−68:54:36.7	0.221	⋯	⋯	⋯	⋯	⋯	⋯	${2.2}_{-0.8}^{+1.1}$	⋯	⋯
N 113 A1	05:13:25.17	−69:22:45.5	13.136	9.399	0.606	0.88/0.21	0.55/0.13	10.7	${214}_{-24}^{+24}$	${2.7}_{-0.3}^{+0.3}$	⋯	⋯
B3	05:13:17.18	−69:22:21.5	6.306	4.332	1.128	1.2/0.29	0.98/0.24	9.2	${184}_{-30}^{+29}$	${1.2}_{-0.2}^{+0.2}$	⋯	⋯

Notes.

^a The continuum peak positions at ∼242.2 GHz for sources in N 105 (this paper) and ∼224.3 GHz for N 113 A1 and B3 (Sewiło et al. 2018). ^b I_peak is the observed 1.2 mm continuum intensity peak. ^c I_mean is the 1.2 mm continuum intensity averaged over the area (Area) within the contour corresponding to the 50% of the 1.2 mm continuum peak. This is the same area used to extract spectra for the analysis (see Section 4). The beam areas for the (N 105–1, N 105–2, N 105–3, N 113) observations are (0.270, 0.273, 0.270, 0.534) arcsec². ^d The observed values, i.e., not corrected for the contribution from the free–free emission (see Section 3). ^e F₅₀ and M_50,gas are flux densities and masses, respectively, calculated for the area above the 50% of the peak intensity. ^f N(H₂) calculated assuming T = T_rot(CH₃OH) (see Section 5 for details). ^g N(H₂) calculated assuming $T={T}_{\mathrm{rot}}({\mathrm{CH}}_{3}\mathrm{CN})$. ^h N(H₂) calculated assuming T = T_rot(SO₂).

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The resulting rotational temperatures (T_rot), column densities (N), velocities (v_LSR), and line widths (Δv_FWHM) are listed in Table 6, along with the estimated abundances with respect to H₂ (N(X)/N(H₂)) and CH₃OH (N(X)/N(CH₃OH)), where X represents a given species. N(H₂) was estimated from the 1.2 mm continuum, as described in Section 5.

Observed spectra with overlaid model fits are presented in Appendix B.

4.3.1. SO₂ Excitation in N 105–2 A and 2 B

The Spitzer/IRS spectra of four point sources in N 105 were first described in Seale et al. (2009). They classified these sources using a principal component analysis, resulting in P (i.e., PAH-dominated) classifications for 1 A and 2 E and PE (PAH-dominated with significant fine-structure emission) classifications for 1 B/blob A (see Section 6) and 2 A/2 B. As it can be seen from Figure 18, the spectra are very similar, and the PAH and fine structure emission are very conspicuous. Silicate in absorption was reported by Seale et al. (2009) for 2 A/2 B and 2 E, but this is not clear at all from the spectra. Jones et al. (2017) reclassified all IRS spectra obtained for LMC sources; they find no evidence of silicate absorption in any of the spectra in N 105. They classify 1 A, 1 B/blob A, 2 A/2 B, and 2 E as YSO3/H ii (1 A, 1 B/blob A, 2 A/2 B) and H ii (2 E) types, i.e., relatively evolved YSOs with emerging compact H ii regions. As pointed out by Jones et al. (2017), the IRS spectra alone are in fact not sufficient to unambiguously classify more evolved YSOs. Furthermore, given the large and variable slit width across the IRS range (36/∼ 0.9 pc at shorter to 111/∼ 2.7 pc at longer wavelengths), point-source spectra can be contaminated by the ambient emission in the wider H ii region, likely the case in N 105 (see Figures 1-2).

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Given the new ALMA data we have obtained for N 105, we have revisited the analysis of these IRS spectra. In Figure 18, there is a hint of a broad absorption feature at ∼15 μm in the spectrum of 2 A/2 B. Broad absorption features in the IRS spectra of YSOs are commonly attributed to solid-state (i.e., ice) features of abundant molecules like H₂O, CO₂, etc. (Oliveira et al. 2009, 2011; Seale et al. 2011; Oliveira et al. 2013); the feature at 15.2 μm is due to CO₂ ice.

Figure 19 shows the results of this new analysis. It closely follows that described in Oliveira et al. (2009, 2011, 2013). We first define a pseudocontinuum over the range 14–17 μm (using spectral regions free of absorption and emission lines) by fitting a low-degree polynomial (top). The spectrum and the fitted continuum are used to calculate the optical depth (bottom). The optical depth for N105–2 A/2 B is compared to that of a well-studied SMC YSO analyzed in Oliveira et al. (2013); the shape of these absorption features are very similar, except that of N 105–2 A/2 B lacks a red wing. While the shape of the CO₂ ice profile reflects composition and environmental conditions like the temperature (see the extensive discussion in Oliveira et al. 2009), in this case the presence of the strong emission line at 15.55 μm due to [Ne iii] precludes any further analysis of the profile shape. For completeness, we analyzed the spectra of the other three sources in the same way; no features attributable to CO₂ ice are detected.

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The column density has been computed using the line strength A = 1.1 × 10⁻¹⁷ cm molecule⁻¹ (Gerakines et al. 1995). Column density calculations are dependent on the exact determination of the pseudocontinuum and can vary by as much as 25% (Oliveira et al. 2009); furthermore, for this particular spectrum, the relatively low signal-to-noise ratio and the aforementioned emission line make this determination difficult. We compute a CO₂ ice column density of (2.63 ± 0.14) × 10¹⁷ cm⁻²; the uncertainty is statistical only, and it does not take into account any uncertainties in the continuum determination. This column density is at the lower end of the ranges reported in Oliveira et al. (2009) and Seale et al. (2011, ∼(1–17)×10¹⁷ cm⁻²), and this is perhaps the reason why neither Seale et al. (2011) nor Jones et al. (2017) identified CO₂ ice in the spectrum of N 105–2 A/2 B.

The KMOS H + K-band spectra are available for three VMC sources in N 105 in the area covered by our ALMA observations. Their source catalog IDs are 558354728325, 558354728333, and 558354728291 (hereafter VMC 325, VMC 333, and VMC 291) and are located nearby ALMA 1.2 mm continuum sources N 105–1 A, N 105–2 B, and N 105–2 E, respectively. The KMOS spectra are shown in Figure 20; the detected spectral lines include the H i (with the brightest Brγ at 2.166 μm), He i (2.058 μm), and H₂ (2.122 μm) lines. No CO bandhead (2.3 μm) and fluorescent Fe ii (1.688 μm) emission (tracing disks) or forbidden [Fe ii] (1.644 μm) emission (tracing outflows) were detected. Figure 21 shows three color images of VMC 325, VMC 333, and VMC 291, combining the KMOS K-band continuum, Brγ, and H₂ images, with the ALMA 1.2 mm continuum contours overlaid.

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To analyze the data, we followed the analysis outlined in Ward et al. (2016). Here, we include a brief summary of the data analysis methods, followed by a discussion of the results for individual sources. The extinction (A_V) values are calculated using the KMOS H- and K-band continuum measurements and are provided in Table 5. These were measured by fitting a third-order polynomial to the continuum spectrum, and the integrated flux of the polynomial was measured for the wavelength ranges 1.5365–1.7875 μm (H band) and 2.028–2.290 μm (K band; Equation (1) in Ward et al. 2016) and assuming an intrinsic H − K color of −0.05 mag, corresponding to a B0 V star. We were unable to estimate A_V for VMC 333 (nearby 2 B), which has a very blue spectrum, indicating that it is not embedded. The value of A_V for VMC 325 (1 A; 42 ± 4 mag, Table 5) is in excellent agreement with that estimated by Oliveira et al. (2006) using the VLT/ISAAC spectrum (∼40 mag; see Section 6). We applied the extinction correction to our measurements for VMC 325 (1 A) and VMC 291 (2 E) described below to constrain the physical properties of the observed sources.

Table 5. VMC Photometry, Visual Extinction, and Accretion Luminosity for KMOS targets

Object	VMC	R.A. (J2000)	Decl. (J2000)	KS	J	Y	AV	Lacc
	Source ID	(deg)	(deg)	(mag)	(mag)	(mag)	(mag)	(L⊙)
VMC 325	558354728325	77.4605865	−68.8848206	13.837 (0.004)	⋯	⋯	42 ± 4	$({2.9}_{-1.5}^{+0.9})\times {10}^{5}$
VMC 333	558354728333	77.4686862	−68.8912009	16.113 (0.014)	17.415 (0.026)	18.070 (0.036)	⋯	$({2.0}_{-0.1}^{+0.1})\times {10}^{2}$ a
VMC 291	558354728291	77.4746443	−68.8937331	15.357 (0.009)	⋯	⋯	12 ± 2	$({2.3}_{-0.6}^{+0.5})\times {10}^{2}$

Note.

^a Not corrected for extinction.

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The accretion luminosity (L_acc) was estimated from the Brγ luminosity (L_Brγ ) assuming the relation of Calvet et al. (2004) for intermediate-mass YSOs holds for high-mass YSOs in the LMC: $\mathrm{log}({L}_{\mathrm{acc}})=-0.7\,+\,0.9(\mathrm{log}({L}_{\mathrm{Br}\gamma })+4)$. The accretion luminosities for VMC 325, VMC 333, and VMC 291 are listed in Table 5.

VMC 325: Figure 21 shows a good positional correlation between the VMC source VMC 325 and the ALMA N 105–1 A 1.2 mm continuum peak, indicating that the near-IR source observed with KMOS is the central protostar in 1 A. With A_V of 42 ± 4 mag, VMC 325 is the most embedded source of the three KMOS sources in the ALMA fields, with by far the reddest spectrum. The spectrum of VMC 325/1 A exhibits a full Brackett series emission in the H + K bands, strong He i emission, and some considerably weaker H₂ emission.

The H₂ emission (tracing shocks) extends toward west-northwest from the continuum source (Figure 21), in agreement with the direction of the outflow reported in Oliveira et al. (2006) (see Section 6). The accretion luminosity of $({2.9}_{-1.5}^{+0.9})\times {10}^{5}$ L_⊙ (see Table 5) is 2–3 orders of magnitude larger than those measured toward similar objects in the Magellanic Clouds (Ward et al. 2016, 2017; van Gelder et al. 2020). This high value of L_acc can be explained by a contribution from the bright UC H ii region to the Brγ emission (e.g., Armand et al. 1996). This interpretation is supported by the fact that the Brγ emission is extended toward VMC 325 (see Figure 21).

VMC 333: The K-band continuum position of VMC 333 is offset by ∼04 (∼0.1 pc or ∼20,000 au) from the 1.2 mm continuum source N 105–2 B. VMC 333 has quite a blue spectrum and is unlikely to be deeply embedded. Figure 21 shows an extended Brγ and H₂ emission in the region. Following background subtraction most of this emission is removed, but there is still Brγ emission coincident with the source; thus, it seems likely that the source does contribute to the emission, although it is unclear whether it is the dominant ionizing source. Assuming it is the dominant ionizing source, the accretion luminosity of VMC 333 is (2.0 ± 0.1) × 10² L_⊙.

The position of the closest Spitzer source to 2 B is offset by ∼08, i.e., the distance is 2 times larger than that between 2 B and VMC 333 (see Figures 4 and 17). The properties of this Spitzer source are uncertain; the Spitzer/IRS spectrum analyzed in Seale et al. (2009) and Jones et al. (2017) was obtained at a position between this source and another Spitzer source nearby 2 A seen in the images, and likely both sources contribute to the emission (see Sections 3.1 and 6.1). At longer Spitzer and Herschel wavelengths, these two sources remain unresolved with offsets between the emission peaks in different bands.

No K-band source has been detected with KMOS at the position of the 1.2 mm continuum source 2 F; the central source may be below a detection limit of the KMOS observations at the K band if it is very young. Mid- and far-IR observations matching the ALMA spatial resolution are needed to learn about the central sources in 2 B and 2 F.

The extended H₂ emission throughout the KMOS field is consistent with the ALMA data showing evidence for strong shocks toward 2 B (see Section 7.5). The brightest H₂ emission is located toward east and northeast from 2 B, but it is also present in the north and southeast. The distribution of the Brγ emission is similar to H₂, but it is the brightest in the north.

VMC 291: There is a very good positional correlation between VMC 291 and the 1.2 mm continuum source N 105–2 E, indicating that VMC 291 is the central protostar in 2 E (see Figure 21). VMC 291 has quite a red spectrum (not quite as red as VMC 325/1 A). The source is embedded with A_V of 12 ± 2 mag. The signal-to-noise ratio in the H-band spectrum of VMC 291/2 E is very poor; however, there is a clear detection of the Brγ and H₂ emission lines in the K-band spectrum (see Figure 20). We have estimated the accretion luminosity for VMC 291 of $({2.3}_{-0.6}^{+0.5})\times {10}^{2}$ L_⊙.

The physical and chemical properties of the continuum sources N 105–2 A and 2 B indicate that they are bona fide hot cores. Their rotational temperatures determined based on multiple tracers exceed 100 K (see Table 5), and their effective FWHM sizes are ∼0.13 pc (Table 4), consistent with the definition of hot cores (≲0.1 pc; e.g., Kurtz et al. 2000). Slightly larger sizes of the LMC hot cores compared to Galactic hot cores are not surprising due to less dust, and as a result the radiation reaches further distances from the protostars (see also Sewiło et al. 2018). Both sources 2 A and 2 B show emission from COMs and are associated with the H₂O/OH masers, as typically observed toward Galactic hot cores.

Figure 22 shows the bar histogram of rotational temperatures and column densities for bona fide hot cores 2 A and 2 B and several other sources from our sample exhibiting high temperature in at least one of the four tracers (CH₃CN, CH₃OH, SO₂, and ³⁴SO₂): 2 C, 2 F, 3 B, and 1 A. We consider these sources as "hot core candidates."

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Source 2 C with the high CH₃OH rotational temperature of ∼95 K with an uncertainty of ∼25% is likely a hot core; however, the SO₂ temperature is warm at ∼32 K. If 2 C is indeed a hot core, the relatively low SO₂ temperature could be the result of blended hot and cold SO₂ components that could not be separated at the signal-to-noise ratio of our observations. Alternatively, the derived SO₂ temperature could be suppressed by non-LTE effects; however, it is unclear whether these non-LTE effects alone would explain the large difference in the rotational temperature between CH₃OH and SO₂. CH₃CN is present toward 2 C, but only the K = 0 component of the 14_K –13_K ladder has been reliably detected. The 2 F continuum peak located ∼1'' to the west from source 2 B is likely a separate source and a hot core with the CH₃OH temperature of 127 ± 33 K. The CH₃OH temperature determination for 3 B has the highest uncertainty, ${159}_{-61}^{+41}$ K, but the source is associated with a CH₃OH maser. Multiple transitions for other species are not available to obtain an independent temperature measurement to confirm the results for 2 F and 3 B. The sizes of 2 C (∼0.15 pc) and 3 B (∼0.12 pc) are consistent with them being hot cores; 2 F is a compact source as well, but it is blended with 2 B.

Source 1 A is an embedded YSO that has started ionizing its immediate surroundings, and it is associated with an outflow (see Section 6). N 105–1 A is associated with cold CH₃OH (∼12 K), but SO₂ has a rotational temperature of ∼96 ± 20 K. The SO₂ peak is offset from the 1 A 1.2 mm continuum peak by ∼06 (see Figure 9). The analysis of the spectrum extracted as a mean at half-peak of the SO₂ emission provided a similar result for the SO₂ temperature (98 ± 20 K) and slightly warmer CH₃OH (22 ± 4 K), which may be indicative of a more compact, hotter CH₃OH component centered on the hot SO₂ source. The SO₂ lines in the 1 A spectrum we analyzed may (at least partially) originate in the area offset from the continuum source and/or CH₃OH may be subthermally excited. The region of the hot SO₂ emission may be a separate source—externally illuminated as no infrared source has been detected at this position, possibly associated with shocks from the 1 A outflow.

Methanol is detected in all 12 continuum sources identified in the N 105 region. The bona fide hot cores, 2 A and 2 B, and the hot core candidates 2 C and 2 F, all contain both hot and cold CH₃OH components. Only hot CH₃OH is detected in the remaining candidate, 3 B, and only cold methanol is detected in the other seven sources. SO₂ also exhibits both hot and cold components in 2 A, 2 B, and 2 F.

The formation of CH₃OH by solid phase hydrogenation of CO ice is the only viable formation pathway in interstellar chemistry (e.g., Herbst & van Dishoeck 2009), thus all the detected CH₃OH originates from a prior cold phase when hydrogenation of CO molecules occurred on grain ice mantles. These molecular ices, containing primarily H₂O, CO, CO₂, CH₄, NH₃, and CH₃OH (e.g., Boogert et al. 2015), were subsequently released, wholly or partially, into the gas. For the hot cores, the mantles can be removed by thermal desorption from dust heated by the protostar and/or sputtering in shock waves associated with outflows (e.g., Jørgensen et al. 2020). Either mechanism could account for the presence of the hot CH₃OH component, but the presence of SiO emission, resulting from the sputtering of refractory dust material (e.g., Schilke et al. 1997), confirms the role of shock waves in producing the hot CH₃OH component in 2 A–C and 2 F.

In hot cores where the chemistry is initiated by shock waves, cooling of the postshock gas is rapid, the temperature approaches that expected from thermal balance determined by protostellar radiative heating, and the long-term chemical evolution closely approaches the predictions of pure thermal desorption models (Charnley & Kaufman 2000). The derived temperatures in the N 105 hot cores, ∼95–170 K, suggest that their observed composition could be explained by simple models of postevaporation chemistry. For example, H₂O and H₂S injected from grain mantles can drive the production of S-bearing molecules, such as SO and SO₂ (Charnley 1997). Hydrogen atom abstraction reactions with H₂S are endoergic but can proceed in hot gas to produce SH and atomic S

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

Protonation of H₂O by molecular ions (${\mathrm{XH}}^{+}={{\rm{H}}}_{3}^{+},{\mathrm{HCO}}^{+}$) followed by electron dissociative recombination reactions can release OH and then atomic O

$$\begin{eqnarray*}&&{{\rm{H}}}_{2}{\rm{O}}\mathop{\longrightarrow }\limits^{{\mathrm{XH}}^{+}}{{\rm{H}}}_{3}{{\rm{O}}}^{+}\mathop{\longrightarrow }\limits^{{{\rm{e}}}^{-}}\mathrm{OH}\mathop{\longrightarrow }\limits^{{\mathrm{XH}}^{+}}{{\rm{H}}}_{2}{{\rm{O}}}^{+}\mathop{\longrightarrow }\limits^{{{\rm{e}}}^{-}}{\rm{O}}\end{eqnarray*}$$

and lead to SO and SO₂ through

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

with an additional contribution to SO formation from reaction of atomic S with OH. Reformation of H₂O and H₂S in reactions of H₂, with each of O, OH, S and SH, is inhibited below about 250 K (Charnley 1997).

In 2 D, SiO is also detected but the cold CH₃OH gas is accompanied by a warm component (∼31 K), significantly less abundant and much cooler than in higher-temperature cores. However, the SiO abundance is ∼8–30 times less abundant than in the hot CH₃OH cores, perhaps indicating the presence of lower sputtering yields in weaker shocks. Less efficient sputtering of ices means less H₂S and H₂O injected into the gas, lower abundances of SH and OH, and hence may also account for the nondetection of SO₂ in 2 D.

Molecular desorption from ice mantles is most probably also the origin of the cold gas detected in CH₃OH and, indirectly, in SO₂. However, it is less clear exactly how these ice mantles were deposited into the gas, considering that the inferred dust temperatures are far too low (∼10–17 K) to allow thermal desorption. The same unresolved issue arises in Galactic dark clouds, where emission from water, methanol, and other complex molecules is detected in dense clumps at locations far from protostars or any outflows (e.g., Bacmann et al. 2012; Cernicharo et al. 2012; Vastel et al. 2014; Wirström et al. 2014; Taquet et al. 2017; Soma et al. 2018; Agúndez et al. 2021).

Several desorption mechanisms have been proposed and include photodesorption, explosion of UV-irradiated ices, grain heating following cosmic-ray impact (Leger et al. 1985), cosmic-ray sputtering (Wakelam et al. 2021), as well as "reactive desorption" in which the energy released in exothermic grain-surface reactions is sufficient to overcome the physisorption binding energy (Minissale et al. 2016; Chuang et al. 2018). Alternatively, a localized kinematic origin has been proposed for the origin of methanol and other putative ice-mantle molecules in Galactic dark clouds. This involves transient heating of dust grains in low-velocity grain–grain collisions. The center-of-mass kinetic energy resulting from drift velocities of ≲1 km s⁻¹ heats the grains, which then cool by evaporation of surface molecules. Relative grain–grain streaming could occur through wave motions (Markwick et al. 2001), in merging collisions between small clumps or filaments (Dickens et al. 2001; Buckle et al. 2006; V. Taquet 2022, in preparation), or in fluid dynamical instabilities (Harju et al.2020).

These mechanisms could also lead indirectly to SO and SO₂ in cold gas via the neutral processes described above: if H₂S is also desorbed from the ices, SH and S can still be produced through protonation and dissociative electron recombination

$$\begin{eqnarray*}&&{{\rm{H}}}_{2}{\rm{S}}\mathop{\longrightarrow }\limits^{{\mathrm{XH}}^{+}}{{\rm{H}}}_{3}{{\rm{S}}}^{+}\mathop{\longrightarrow }\limits^{{{\rm{e}}}^{-}}\mathrm{SH}\mathop{\longrightarrow }\limits^{{\mathrm{XH}}^{+}}{\mathrm{SH}}^{+}\mathop{\longrightarrow }\limits^{{{\rm{e}}}^{-}}{\rm{S}},\end{eqnarray*}$$

even though H₂S destruction by H atoms is inefficient. Model calculations based on the grain-streaming picture are able to explain the observed close spatial correspondence between CH₃OH and SO in maps of several dark clouds (Buckle et al.2006).

Thus, the origin of the cold methanol, as well as of SO and SO₂, could in principle be due to several processes. For the cold CH₃OH found in N 105, photodesorption is very unlikely due to the high visual extinctions and because ice photolysis experiments indicate that CH₃OH dissociates on desorption (Bertin et al. 2016; Martín-Doménech et al. 2016). Processes involving cosmic rays or reactive desorption may be expected to produce extended, almost homogenous, CH₃OH distributions. We find that cold CH₃OH emission is widespread in N 105–1 A relative to other sources, and this may reflect a higher cosmic-ray flux in this field. It is difficult to estimate the cosmic-ray flux necessary to explain the cold CH₃OH emission as this requires knowledge of the cosmic-ray composition, the manner in which the grain is heated (whole-grain versus spot heating), as well as whether a radical explosion can be initiated (e.g., Leger et al. 1985). A markedly higher cosmic-ray flux in N 105–1 should have consequences for other molecules. In N 105–1, we only detect simpler molecules that most likely originate in gas-phase reactions. Of these, H¹³CO⁺ would be most sensitive to the cosmic-ray ionization rate. The fact that the derived H¹³CO⁺ abundances in N105-1 are lower than toward most of the sources in N 105–2 and N 105–3 probably rules out the cosmic-ray desorption mechanism as the important contributor to the production of the cold CH₃OH emission.

If reactive desorption of CH₃OH upon surface formation produced the widespread cold CH₃OH emission in N 105–1 A, it is difficult to understand why it is also not evident in N 105–2 and N 105–3. For cores where the cold CH₃OH emission is more compact, grain collisions could be responsible for mantle desorption; evaluating this contribution will require a detailed understanding of the kinematics and higher angular resolution observations to probe structures (clumps, filaments) not resolved in our ALMA observations.

In Figures 23 and 24, we compare the fractional abundances with respect to H₂ for all the 1.2 mm continuum sources in N 105, calculated as described in Section 5. Figure 23 shows the bar histogram for the species detected toward all the sources: CH₃OH, SO, CS, and H¹³CO⁺. The largest differences between the sources are observed in H¹³CO⁺.

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In Figure 24, we separately compare fractional abundances of COMs, N-bearing species, S-/O-bearing species and SiO, and C-/S-bearing species. In general, N 105–2 is the most chemically rich region with the detection of COMs other than CH₃OH and the highest fractional abundances of other species; the only exception is SO and CS, which have comparable abundances toward the hot core candidate 3 B. The fewest number of species has been detected toward the continuum sources in N 105–3, but those detected have fractional abundances comparable to those observed toward other fields, indicating a smaller size or lower density for these sources, resulting in reduced emission line strengths. N-bearing species have lower abundances than S-bearing species in all three ALMA fields.

Sources 2 A, 2 B, and 2 C are the only sources in the N 105 ALMA fields with detection of CH₃CN. Recently, Mininni et al. (2021) analyzed the CH₃CN data for a sample of high-mass star-forming regions at different evolutionary stages (high-mass starless cores, high-mass protostellar objects, and UC H ii regions) and concluded that the mean abundance of CH₃CN is a good tracer of the early stages of high-mass star formation; it shows 1 order of magnitude increase from starless cores to later evolutionary stages. In the LMC, we can use this result to investigate the relative ages of sources with the detection of CH₃CN. The lowest CH₃CN abundance was measured toward 2 C: 2 times lower than that observed toward 2 B and 5 times lower than for 2 A. Thus, 2 C appears to be the youngest of the three sources, which is consistent with its location in the optical dark lane and a lack of an IR match in the available data.

The continuum source A in N 105–1 is the only source in our ALMA fields with the detection of H recombination lines, indicating the presence of ionized gas. It is still an embedded protostar, but it has already started ionizing its surroundings. N 105–2 appears to be the site of the most vigorous ongoing star formation with multiple maser sites, a YSO with the detection of the CO₂ ice band, and the presence of COMs, hot cores, and deuterated species. In the south, hot core candidate N 105–3 B is associated with 6.7 and 12.2 GHz CH₃OH masers. These are radiatively excited Class ii masers known to be tracers of a very early phase of massive star formation (e.g., Cragg et al. 1992; Ellingsen 2006). It is plausible that the source is at the early hot core phase and thus not all hot core tracers (such as SO₂) have achieved detectable levels.

Formamide (NH₂CHO) is the simplest naturally occurring amide and has been proposed as a precursor of prebiotic molecules with a key role in the emergence of life on Earth (see López-Sepulcre et al. 2019 for a review). We detected the NH₂CHO 12_2,10–11_2,9 transition at 260.189 GHz (the strongest NH₂CHO line in the observed spectral range) at a 3.2σ level toward hot core N 105–2 A. The statistical uncertainty of our NH₂CHO detection takes into account the uncertainty in the emission due to the contribution of the overlapping CH₂CO transition at 260.192 GHz to the detected line flux. Since our identification of NH₂CHO is based on a single, low signal-to-noise ratio transition, which is blended with another line, we can only consider this detection as tentative. If the presence of NH₂CHO in the LMC is confirmed, it will constitute the first detection of this astrobiologically relevant molecule in the extragalactic low-metallicity environment.

Based on the CH₃OH rotational temperature, N 105–2 A is the hottest of the sources observed in N 105, which is likely the reason for the detection of gas-phase NH₂CHO only in this core. NH₂CHO has one of the highest physisorption binding energies of the most common interstellar COMs (including CH₃OH; Penteado et al. 2017). Lower maximum dust temperatures could account for its nondetection in the other sources in N 105.

The formation routes of NH₂CHO are still debated. It may proceed via gas phase pathways involving H₂CO and NH₂ (e.g., Barone et al. 2015; Skouteris et al. 2017) and/or grain surface reactions involving hydrogenation of HNCO (e.g., Charnley & Rodgers 2008); although laboratory studies indicate that the latter route may not be viable (e.g., Noble et al. 2015; López-Sepulcre et al. 2019 and references therein). Observations show that NH₂CHO is most abundant in the hot parts of protostellar envelopes (T > 100 K; hot cores: e.g., Allen et al. 2017; Bisschop et al. 2007; hot corinos, similar to hot cores but formed around low- and intermediate-mass protostars: e.g., López-Sepulcre et al. 2015; Marcelino et al. 2018; Bianchi et al. 2019) and regions dominated by shocks (e.g., in protostellar outflows; e.g., Yamaguchi et al. 2012; Mendoza et al. 2014; Ceccarelli et al. 2017; Codella et al. 2017). It has not been established yet whether the formation route of NH₂CHO depends on the environment, although gas-phase reactions seem to be the dominant pathway leading to NH₂CHO in the protostellar outflows (e.g., Codella et al. 2017). Observations of NH₂CHO in the LMC offer unique tests of its formation models in the metal-poor environments with lower dust content and higher UV radiation fields. While our observations do not provide enough information to draw reliable conclusions, we have attempted a preliminary investigation of the NH₂CHO formation routes in the LMC.

Our observations did not cover any H₂CO lines, but we detected HNCO toward 2 A. The NH₂CHO and HNCO fractional abundances with respect to H₂ for 2 A are $X({\mathrm{NH}}_{2}\mathrm{CHO})=({1.7}_{-0.6}^{+0.5})\times {10}^{-10}$ and X(HNCO) = (5.4 ± 0.7) × 10⁻¹⁰, respectively (see Table 6). One of the methods used to test the dominant formation route of NH₂CHO is to investigate its abundance correlation with other molecules that are thought to be chemically linked (i.e., forming from a common precursor or one forming from the other). Almost a linear correlation over several orders of magnitude in the fractional abundances was found for NH₂CHO and HNCO in the Milky Way (e.g., López-Sepulcre et al. 2019). We can test whether the fractional abundances observed toward 2 A follow the X(NH₂CHO) − X(HNCO) correlation found for Galactic sources by estimating the expected X(NH₂CHO) from the observed X(HNCO) and comparing it to the observed value (after correcting both observed values for a difference in metallicity, assuming Z_LMC = 0.5 Z_⊙). The observational correlation (power-law fit) that holds for sources with NH₂CHO detection (no upper limits) is available in literature: $X{({\mathrm{NH}}_{2}\mathrm{CHO})=0.04X(\mathrm{HNCO})}^{0.93}$ (López-Sepulcre et al. 2015), resulting in an expected X(NH₂CHO) for 2 A of ∼1.8 × 10⁻¹⁰ based on the observed, metallicity-corrected X(HNCO). The observed, metallicity-corrected X(NH₂CHO) (∼3.4 × 10⁻¹⁰) is a factor of ∼1.9 higher than the expected value but still within the scatter of the X(NH₂CHO) −X(HNCO) relation. The X(NH₂CHO) −X(HNCO) correlation observed toward Galactic sources may not come from a direct chemical link between HNCO and NH₂CHO, but rather be the result of their similar response to the temperature of their environment (e.g., Quénard et al. 2018).

Using the ALMA Band 6 and Band 7 observations, Shimonishi et al. (2021) have recently detected multiple transitions of NH₂CHO toward the submillimeter continuum source SMM1 in the WB89–789 star-forming region in the extreme outer Galaxy. WB89–789 SMM1 is located at the galactocentric distance of 19 kpc where metallicity (traced by the oxygen abundance) is expected to be a factor of 4 lower than in the solar neighborhood (e.g., Fernández-Martín et al. 2017; see Shimonishi et al. 2021 and references therein). Shimonishi et al. (2021) found that even though WB89–789 SMM1 and hot cores in the LMC represent low-metallicity environments, there is no resemblance between the extreme outer Galaxy and the LMC sources. The authors suggest that the dissemblance might be a result of differences in the environments such as the strength of the interstellar radiation field, which is significantly higher in the LMC (see Section 1).

Shimonishi et al. (2021) determined X(NH₂CHO) of (1.8 ± 0.1) × 10⁻¹¹ and X(HNCO) of (2.7 ± 0.8) × 10⁻¹⁰ for WB89–789 SMM1 for physical scales corresponding to those probed by the ALMA observations of the LMC hot cores (0.1 pc), which is 1 order of magnitude lower than X(NH₂CHO) estimated for 2 A. The expected value of X(NH₂CHO) based on the X(NH₂CHO) − X(HNCO) correlation of López-Sepulcre et al. (2015) for metallicity-corrected X(HNCO) measured toward WB89–789 SMM1 is ∼5.3 × 10⁻¹¹. This value is a factor of ∼1.4 lower than the metallicity-corrected observed X(NH₂CHO) (∼7.2 × 10⁻¹¹)—a result similar to that we obtained for 2 A.

Figure 27 shows that the abundance of NH₂CHO in 2 A is higher than those measured toward the Orion Hot Core and Sgr B2(N) in single-dish observations (see a discussion in Section 7.7). There are, however, single-dish measurements of X(NH₂CHO) toward Galactic hot cores, which are 1 order of magnitude higher than X(NH₂CHO) in 2 A. For example, Bisschop et al. (2007) measured X(NH₂CHO) of a few times 10⁻⁹ for six hot cores: G24.78+0.08, G75.78+0.34, NGC 6334 IRS1, NGC 7538 IRS1, W3(H₂O), and W33 A; these values are 1–2 orders of magnitude higher than X(NH₂CHO) reported in the literature for the same sources.

The NH₂CHO transition in the spectrum of 2 A is blended with a CH₂CO 13_1,13–12_1,12 line at 260.192 GHz (a frequency/velocity shift of ∼2.9 MHz/∼3.3 km s⁻¹; see Table 3), making the measurements of the NH₂CHO column density and abundance less reliable. We have investigated the NH₂CHO–CH₂CO line blending issue in detail based on other CH₂CO transitions detected in our observations: the CH₂CO line at 244.712 GHz (Table 3) and two weak lines at 242.375 and 242.398 GHz (<3σ detections), which helped constrain the fit.

While the formation on grains is the most likely scenario for CH₂CO, a gas-phase formation route is also possible (e.g., Bisschop et al. 2007). The NH₂CHO emission peak is offset from the HNCO peak in 2 A (∼014 or ∼0.034 pc/∼7000 au at 50 kpc), but it is also not coincident with the CH₂CO peak. Observations of additional transitions of NH₂CHO are needed to confirm our tentative detection and to conduct a more reliable investigation of the spatial correlation between the NH₂CHO and HNCO emission that would allow us to test the NH₂CHO grain-surface formation scenario in the significantly different chemical laboratory of the LMC. Observations of H₂CO and NH₂, when compared to the NH₂CHO data, would help test the gas-phase NH₂CHO formation route proposed by Barone et al. (2015) in which H₂CO and NH₂ are its precursors. Interstellar NH₂ was detected from the ground in Sgr B2 (van Dishoeck et al. 1993) and could be observed with ALMA. In addition, the abundance ratio of the deuterated forms of NH₂CHO can provide a strong constraint on its formation route (e.g., Coutens et al. 2016; López-Sepulcre et al. 2019).

The integrated intensity images of N 105–2 and the spectra of individual continuum components reveal differences in the chemical makeup of hot cores 2 A and 2 B, which are separated by ∼32 (∼0.78 pc/∼160,000 au). N-bearing species and deuterated species have higher abundances toward 2 A (up to a factor of ∼2 for CH₃CN and HDO), while S-bearing species and SiO have higher abundances toward source 2 B (by a factor of ∼3–4). The hot SO₂ abundance is larger in 2 B than in any other LMC hot core observed to date. NH₂CHO and HDCO have only been detected toward source 2 A. Finally, the kinematic structure observed toward 2 B is much more complex than that in 2 A, indicating the presence of at least two velocity structures.

In general, spectral lines for all the species are broader toward 2 B than 2 A, indicating the presence of more significant large-scale motions in this region. The difference in line widths, as well as line intensities, is particularly striking for SO, ³³SO, SO₂, ³⁴SO₂, and SiO (see Figure 25). The enhanced abundance of S-bearing species and SiO toward 2 B indicates the shock origin of some of the emission, consistent with the observed greater kinematic complexity, which is reflected in the broad line profiles. The enhanced production of the S- and Si-bearing molecules in shocks is a result of sputtering or destruction of refractory grain cores that release the Si and S atoms to the gas, making them available for chemical reactions (e.g., Schilke et al. 1997; Gusdorf et al. 2008; van Dishoeck 2018 and references therein; see also a discussion in Section 7.2).

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Other species that are often detected in low-velocity outflows have broader lines toward 2 B than 2 A; these include the ice chemistry products such as CH₃OH, CH₃CN, HNCO, and HDO that can be be released to the gas by shock-driven sublimation in addition to thermal sublimation in the hot core region (e.g., van Dishoeck 2018; Öberg & Bergin 2021), as well as H¹³CN that can be produced via the hot-gas-phase chemistry in the cavity walls (e.g., Bruderer et al. 2009). Rotation of the envelope can also broaden spectral lines, but we are not able to distinguish between the outflow and rotation based on our relatively low-spatial-resolution data. Molecular species with narrower lines mostly originated in the more quiescent hot core region. All these processes most likely take place in 2 A as well but are not as much affected by complex kinematics as in 2 B; since the temperature is higher in the 2 A hot core (two times for CH₃CN), the thermal evaporation of grain ice mantles is more efficient in this region, resulting in higher or comparable fractional abundances for all molecules except the S- and Si-bearing species.

In Figure 26, we compare the fractional abundances with respect to H₂ for the LMC hot cores with COMs from literature (N 113 A1 and B3, Sewiło et al. 2018; ST16, Shimonishi et al. 2020) and the newly identified hot cores in N 105: 2 A and 2 B. We reanalyzed the spectra of N 113 A1 and B3 using the same spectral extraction and modeling techniques and the subsequent analysis for N 113 A1/B3 as for hot cores and other sources in N 105 and obtained results consistent with those reported in Sewiło et al. (2018); the results are listed in Table 6.

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As described in Section 5, we adopted $T({\mathrm{CH}}_{3}\mathrm{CN})$ for hot cores N 105–2 A and 2 B to calculate N(H₂). For ST16, we recalculated the molecular abundances from Shimonishi et al. (2020) by estimating N(H₂) using the dust temperature of 60 K provided in the paper (consistent with $T({\mathrm{CH}}_{3}\mathrm{CN})$ of ${53}_{-7}^{+10}$ K) and assuming the same LMC dust-to-gas ratio as for N 105 and N 113 (Section 5). T(CH₃OH) is the only temperature determination available for hot cores A1 and B3 in N 113, and it was used for the analysis.

Figure 26 compares fractional abundances of COMs detected toward at least two LMC hot cores (CH₃OH, CH₃OCH₃, HCOOCH₃, and CH₃CN) and simple molecules that were observed toward all of the sources (SO₂, OCS, ³³SO, and SiO). N 113 A1 and B3 hot cores represent a class of "organic-rich" hot cores, while ST16 is an example of an "organic-poor" hot core as defined by Shimonishi et al. (2020); see Section 1.

It is evident from Figure 26 that N 113 A1 and B3 remain the LMC hot cores with the highest abundance of CH₃OH with respect to H₂, the only detection of HCOOCH₃, and the most reliable detection of CH₃OCH₃. The CH₃OCH₃ lines detected toward N 105–2 A and 2 B have low signal-to-noise ratios, and thus the molecular abundance calculations are less reliable. As discussed above, the enhanced abundances of SiO and SO₂ toward N 105–2 B are likely the results of the strong shock activity in the region. The SiO abundance toward N 113 B3 is even higher; similarly to 2 B, the environment of this source is highly dynamic. There are no large variations in the OCS abundance between the known LMC hot cores. The local environment appears to have a significant impact on the observed molecular abundances.

Figure 26 demonstrates that N 105–2 A and 2 B are chemically more similar to the organic-rich hot cores A1 and B3 in N 113 than to ST16 and the other organic-poor hot core ST11 where no COMs were detected (Shimonishi et al. 2016b). In the following section, we compare the fractional abundances of the species detected toward N 105–2 A and 2 B to those detected toward several representative Galactic hot cores to confirm their classification as organic-rich hot cores. It is expected that in general the fractional abundances of the molecular species observed toward the organic-rich hot cores scale with metallicity.

To compare the fractional molecular abundances derived for the LMC hot cores to those observed toward the Galactic hot cores, we first correct them for a difference in metallicity between these two galaxies. For this analysis, we adopt a mean metallicity of −0.30 ± 0.08 dex (i.e., half of that observed in the solar neighborhood; based on the abundance of O and Si) estimated by Rolleston et al. (2002) based on the analysis of all the classes of objects used to trace the chemical composition of the ISM in the LMC (e.g., H ii regions, F-type supergiants, Cepheids, and B-type giants/bright giants), which is in agreement with a metallicity of −0.31 ± 0.04 they determined based on the young B-type dwarfs. This is also consistent with the metallicity determination based on the Fe abundance measured toward F-type supergiants (e.g., Russell & Dopita 1992 and references therein) and Cepheids (e.g., Luck et al. 1998). The overall conclusions presented in this section do not change for a lower value of Z_LMC of 0.4 Z_⊙ found in some studies (see, e.g., Maeder et al. 1999 and references therein).

To correct the fractional abundances measured toward the LMC hot cores for a difference in metallicity between the LMC and the Galaxy, we have multiplied them by a factor of 1/Z_LMC. This simple scaling does not take into account individual elemental abundances; in the LMC, the C, O, and N elemental abundances are lower than in the Galaxy by factors of 2.45, 2.19, and 4.83, respectively (Russell & Dopita 1992). While the individual elemental abundances matter, it is not clear without detailed physicochemical modeling (out of scope of the present paper) how multimetal species should be scaled with the metallicity. The scaling of these species by the product of their metal abundances may be correct for the molecules formed in the gas phase; however, for those species that form on dust grains it only depends on the grain abundance (a total surface area available for chemistry) and the abundance of the least abundant atom.

The metallicity-corrected molecular abundances for N 105–2 A and 2 B, as well as hot cores in N 113, are compared to the abundances observed toward selected Galactic hot cores in Figure 27. We compare our observations to the single-dish observations of the Orion Hot Core, W3(H₂O), and Sgr B2(N) that trace similar physical scales—within a factor of a few of the physical scales probed by our LMC ALMA observations. The data for the Orion Hot Core come from Sutton et al. (1995) where available, otherwise the fractional abundances from Blake et al. (1987) are plotted (CH₃CN, HNCO, and HDO). The data for W3(H₂O) and Sgr B2(N) come from Helmich & van Dishoeck (1997) and Nummelin et al. (2000), respectively. The W3(H₂O) observations are the most similar to our ALMA observations of N 105 in terms of the probed physical scales (∼0.16 pc versus ∼0.12 pc in N 105).

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Metallicity-scaled CH₃OH abundances with respect to H₂ observed toward N 113 A1/B3 and N 105–2 A are the most similar to those observed toward Galactic hot cores, as can been seen in Figure 27; they are at the lower end of the range of those observed toward a larger Galactic hot cores sample (e.g., Mookerjea et al. 2007). The CH₃OH abundance estimated for N 105–2 B is lower than for N 105–2 A and the two hot cores in N 113 and is more similar to that of ST16, indicating that it may be an intermediate case between the organic-rich and organic-poor types of the LMC hot cores proposed by Shimonishi et al. (2020). Gerner et al. (2014) reported a low median CH₃OH abundance of 2.6 × 10⁻⁸ (for a typical hot core temperature of 100 K), comparable (within a factor of 2) to the 2 B's metallicity-scaled CH₃OH abundance of 1.3 × 10⁻⁸, estimated based on the IRAM 30 m telescope observations tracing the 0.02–0.4 pc spatial scales of 11 Galactic hot molecular cores (G9.62+0.19, G10.47+0.03, G29.96−0.02, G31.41+0.31, G34.26+0.15, G45.47+0.05, G75.78+0.34, NGC 7538B, Orion–KL, W31 IRS5, and W3(H₂O)); the authors acknowledge that their calculations provide CH₃OH abundances up to a factor of 10 lower than those reported in literature for the same objects.

For all hot cores in N 105 and N 113, the metallicity-corrected CH₃OCH₃ (N 105 and N 113) and HCOOCH₃ (N 113) abundances are at the lower end of the range of those observed toward Galactic hot cores. The difference between the fractional abundances of the N-bearing species between the LMC (metallicity-scaled) and Galactic hot cores appears to be larger than that for COMs (e.g., CH₃CN and HC₃N; see Figure 27); however, they are still consistent with Galactic values. In Figure 27, similarly low abundances of CH₃CN, HNCO, and HC₃N as for 2 A and 2 B have been reported for the Galactic hot core W3(H₂O). The interesting result highlighted in Figure 27 is that the metallicity-corrected abundance of NH₂CHO observed toward 2 A is higher than those for the Orion and Sgr B2(N) hot cores. While intriguing, we stress that more reliable observations of NH₂CHO toward 2 A (higher sensitivity and more than one transition) are needed to draw reliable conclusions based on this result. No observations of N-bearing species exist for N 113 A1/B3 hot cores (Sewiło et al. 2018).

Molecular abundances reported in the literature for a single region can vary significantly depending on the spatial scales probed and differences in the method for determining the column densities and abundances; e.g., for W3(H₂O) there is a difference of about two orders of magnitude between the measurements reported in Helmich & van Dishoeck (1997) and Bisschop et al. (2007). The measurements from Helmich & van Dishoeck (1997) that match the metallicity-scaled LMC data well (Figure 27) are based on the James Clerk Maxwell Telescope observations, which at the distance of W3 trace similar spatial scales as our ALMA observations in the LMC; also, the data analysis is similar. These examples highlight the importance of comparing the abundances observed toward the LMC hot cores with those of a carefully selected Galactic hot cores measured based on the observations sampling similar physical scales and using a similar data analysis methods. However, a comparison to a broader population of Galactic hot cores does not change our overall conclusion that hot cores 2 A and 2 B are not significantly different than Galactic hot cores.

Figure 27 shows that hot cores 2 A and 2 B indeed belong to the organic-rich category proposed by Shimonishi et al. (2020) (yet to be confirmed as a class based on a larger sample of hot cores) since the molecular abundances roughly scale with metallicity and larger COMs are detected.

Chemically distinct hot core types observed in the LMC may be explained if there are metallicity inhomogeneities in the LMC disk. The chemical composition and dynamics of part of the LMC might have been altered by a close encounter between the LMC and SMC (Z_SMC = 0.1–0.2 Z_⊙) about 0.2 Gyr ago as suggested by the hydrodynamical simulations (e.g., Fujimoto & Noguchi 1990; Bekki & Chiba 2007a; Yozin & Bekki 2014) and the observational evidence including a highly asymmetric distribution of the H i and CO gas in the LMC and the presence of two H i velocity components separated by ∼50 km s⁻¹ in velocity, spatially connected by bridge H i features, and showing complementary spatial distributions on a kpc scale (e.g., Luks & Rohlfs 1992). One of the H i velocity components corresponds to the gas extending over the entire LMC disk and is dubbed the "D component." The lower velocity gas ("L component") is more spatially confined and introduces asymmetry in the distribution of the H i gas in the LMC.

According to the model, during the tidal interaction between the LMC and SMC, the H i gas is stripped from both galaxies, and the remnant gas falls down on each galaxy. In the LMC, this infalling H i gas is observed as the H i L component and has a relative velocity of ∼50 km s⁻¹. A collision with the H i gas in the LMC disk triggered the formation of the young massive cluster R136 in 30 Dor and massive stars in the region extending south of 30 Dor–H i Ridge; it includes two major elongated CO clouds in the southeast of the LMC: the Molecular Ridge and the CO Arc (Fukui et al. 2017). A similar scenario of massive star formation triggered by the colliding H i flows as a result of the SMC–LMC interaction was recently proposed for star-forming regions N 44 and N 159 (Fukui et al. 2019; Tokuda et al. 2019; Tsuge et al. 2019).

Numerical simulations show that the gas infalling on the LMC disk contains a large amount of metal-poor gas from the SMC (Bekki & Chiba 2007b), replenishing the higher-metallicity material in the LMC with metal-poor gas and causing metallicity inhomogeneities. Fukui et al. (2017) and Tsuge et al. (2019) found that the H i L component is metal-poor in the H i Ridge and N 44, confirming this model prediction. Tsuge et al. (2019) estimated the fraction of the SMC gas in the H i Ridge (including R136) and N 44 of 0.5 and 0.3, respectively. The metal content was estimated based on the correlation between the dust optical depth at 353 GHz (τ₃₅₃; measured using the combined Planck and IRAS data) and the H i intensity (gas/dust ∼W(H i)/τ₃₅₃).

The differences in the physical and chemical conditions in each galaxy are likely significant enough for gas mixing to result in a range of environments that would lead to variations in the hot core COM abundances. Regions affected by tidal interactions between the LMC and SMC would be characterized by a lower dust-to-gas ratio, stronger UV radiation fields, and consequently higher dust temperatures than in the unaffected LMC gas (e.g., van Loon et al. 2010; Roman-Duval et al. 2014). The lower dust content means that there would be less dust grains for surface chemistry. Higher dust temperatures may result in less efficient formation of COMs as predicted by the warm ice chemistry model and astrochemical simulations (e.g., Shimonishi et al. 2016a; Acharyya & Herbst 2018). The number of heavy elements available for chemistry would also be lower because in the SMC the elemental abundances of gaseous (N, C, O) atoms are ∼(3, 2, 2) lower than in the LMC (e.g., Russell & Dopita 1992). The cosmic-ray density in the SMC is only ∼15% of that observed in the solar neighborhood (compared to ∼25% in the LMC; Abdo et al. 2010b), making the cosmic-ray-induced UV radiation less effective than in the LMC.

Hot core ST11 with nondetection of COMs is located near the region thought to be affected by the interaction between the LMC and SMC, possibly in a lower-metallicity region, while all known hot cores with COMs are associated with the unaffected areas in the disk. This distribution of hot cores with and without COMs may be the result of a source selection effect; a larger sample of hot cores is needed to investigate the connection between the metallicity inhomogeneities and the properties of hot cores.