The Detection of a Hot Molecular Core in the Extreme Outer Galaxy

The target star-forming region is WB 89–789 (Brand & Wouterloot 1994). The region contains three Class I protostar candidates identified by near-infrared observations (Brand & Wouterloot 2007), and one of them is a main target of the present ALMA observations. The region observed with ALMA is indicated on a near-infrared two-color image shown in Figure 1. The observed position is notably reddened compared with other parts of WB 89–789.

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Observations were conducted with ALMA in 2018 and 2019 as part of the Cycle 5 (2017.1.01002.S) and Cycle 6 (2018.1.00627.S) programs (PI: T. Shimonishi). A summary of the present observations is shown in Table 1. The pointing center of the antennas is R.A. = 06^h17^m23^s and decl. = 14°54'41'' (ICRS). The total on-source integration time is 115.5 minutes for Band 6 data and 64.1 minutes for Band 7. Flux and bandpass calibrators are J0510+1800, J0854+2006, and J0725–0054 for Band 6, and J0854+2006 and J0510+1800 for Band 7. Phase calibrators are J0631+2020 and J0613+1708 for Band 6 and J0643+0857 and J0359+1433 for Band 7. Four spectral windows are used to cover the sky frequencies of 241.40–243.31, 243.76–245.66, 256.90–258.81, and 258.76–260.66 GHz for Band 6, ad 337.22–339.15, 339.03–340.96, 349.12–351.05, and 350.92–352.85 GHz for Band 7. The channel spacing is 0.98 MHz, which corresponds to 1.2 km s⁻¹ for Band 6 and 0.85 km s⁻¹ for Band 7. The total number of antennas is 45–49 for Band 6 and 43–44 for Band 7. The minimum–maximum baseline lengths are 15.1–783.5 m for Band 6 and 15.1–500.2 m for Band 7. The FWHM of the primary beam is about 25'' for Band 6 and 18'' for Band 7.

Table 1. Observation Summary

	Observation	On-source	Mean	Number	Baseline				Channel
	Date	Time	PWV a	of	Min	Max	Beam size b	MRS c	Spacing
		(min)	(mm)	Antennas	(m)	(m)	('' × '')	('')
Band 6	2018 Dec 6—	115.5	0.5–1.5	45–49	15.1	783.5	0.41 × 0.50	5.6	0.98 MHz
(250 GHz)	2019 Apr 16								(1.2 km s−1)
Band 7	2018 Apr 30—	64.1	0.6–1.0	43–44	15.1	500.2	0.46 × 0.52	5.4	0.98 MHz
(350 GHz)	2018 Aug 22								(0.85 km s−1)

Notes.

^a Precipitable water vapor. ^b The average beam size of the continuum achieved by TCLEAN with Briggs weighting and the robustness parameter of 0.5. Note that we use a common circular restoring beam size of 050 for Band 6 and 7 data to construct the final images. ^c Maximum recoverable scale.

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Raw data are processed with the Common Astronomy Software Applications (CASA) package. We use CASA 5.4.0 (Band 6) and 5.1.1 (Band 7) for the calibration and CASA 5.5.0 for the imaging. The synthesized beam sizes of 039–042 × 049–052 with a position angle of −36° for Band 6 and 045–046 × 051–052 with a position angle of −54° for Band 7 are achieved with Briggs weighting and a robustness parameter of 0.5. In this paper, we use a common circular restoring beam size of 050, which corresponds to 0.026 pc (5350 au) at the distance of WB 89–789. The synthesized images are corrected for the primary beam pattern using the impbcor task in CASA. The continuum image is constructed by selecting line-free channels. Before the spectral extraction, the continuum emission is subtracted from the spectral data using the CASA's uvcontsub task.

The spectra and continuum flux are extracted from the 050 diameter circular region centered at R.A. = 06^h17^m24073 and decl. = 14°54'4227 (ICRS), which corresponds to the submillimeter continuum center of the target and is equivalent to the hot-core position. Hereafter, the source is referred to as WB 89–789 SMM1.

Figures 2–3 show submillimeter spectra extracted from the continuum center of WB 89–789 SMM1. Spectral lines are identified with the aid of the Cologne Database for Molecular Spectroscopy ⁶ (CDMS; Müller et al. 2001, 2005) and the molecular database of the Jet Propulsion Laboratory ⁷ (JPL; Pickett et al. 1998). The detection criterion adopted here is the 3σ significance level and the velocity coincidence with the systemic velocity (V_sys) of WB 89–789 SMM1 (34.5 km s⁻¹). The lines with a significance level higher than 2.5σ but lower than 3σ are indicated as tentative detection in the tables in Appendix A. More than 85% of lines are detected above the 5σ level.

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Line parameters are measured by fitting a Gaussian profile to detected lines. We estimate the peak brightness temperature, the FWHM, the LSR velocity, and the integrated intensity for each line based on the fitting. For spectral lines for which a Gaussian profile does not fit well, their integrated intensities are calculated by directly integrating the spectrum over the frequency region of emission. Full details of the line fitting can be found in Appendix A (tables of measured line parameters) and Appendix B (figures of fitted spectra). The table also contains the estimated upper limits on important nondetection lines.

A variety of carbon-, oxygen-, nitrogen-, sulfur-, and silicon-bearing species, including COMs containing up to nine atoms, are detected from WB 89–789 SMM1 (see Table 2). Multiple high-excitation lines (upper state energy >100 K) are detected for many species. Measured line widths are typically 3–6 km s⁻¹. Most of the lines consist of a single velocity component, but SiO has Doppler-shifted components at V_sys ± 5 km s⁻¹ as indicated in Figure B1 in Appendix B.

Figures 4–5 show synthesized images of continuum and molecular emission lines observed toward the target region. The images are constructed by integrating spectral data in the velocity range where the emission is detected. Most molecular lines, except for those of the molecular radicals CN, CCH, and NO, have their intensity peak at the continuum center, which corresponds to the position of a hot core. Simple molecules such as H¹³CO⁺, H¹³CN, CS, and SO are extended compared to the beam size. Secondary intensity peaks are also seen in those species. Complex molecules and HDO are concentrated at the hot-core position. A characteristic symmetric distribution is seen in SiO. Further discussion about the distribution of the observed emission is presented in Section 4.2.

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3.3.1. Rotation Diagram Analysis

Column densities and rotation temperatures are estimated based on the rotation diagram analysis for the molecular species where multiple transitions with different excitation energies are detected (Figure 6). We here assume an optically thin condition and the local thermodynamic equilibrium (LTE). We use the following formulae based on the standard treatment of the rotation diagram analysis (e.g., Sutton et al. 1995; Goldsmith & Langer 1999):

$$\begin{eqnarray}&&\mathrm{log}\left(\displaystyle \frac{{N}_{u}}{{g}_{u}}\right)=-\left(\displaystyle \frac{\mathrm{log}e}{{T}_{\mathrm{rot}}}\right)\left(\displaystyle \frac{{E}_{u}}{k}\right)+\mathrm{log}\left(\displaystyle \frac{N}{Q({T}_{\mathrm{rot}})}\right),\end{eqnarray} \tag{ 1 }$$

where

$$\begin{eqnarray}&&\displaystyle \frac{{N}_{u}}{{g}_{u}}=\displaystyle \frac{3k\int {T}_{b}{dV}}{8{\pi }^{3}\nu S{\mu }^{2}},\end{eqnarray} \tag{ 2 }$$

and N_u is a column density of molecules in the upper energy level, g_u is the degeneracy of the upper level, k is the Boltzmann constant, ∫T_b dV is the integrated intensity estimated from the observations, ν is the transition frequency, S is the line strength, μ is the dipole moment, T_rot is the rotational temperature, E_u is the upper state energy, N is the total column density, and Q(T_rot) is the partition function at T_rot. All of the spectroscopic parameters required in the analysis are extracted from the CDMS or JPL database. Derived column densities and rotation temperatures are summarized in Table 3.

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Table 3. Estimated Rotation Temperatures, Column Densities, and Source Sizes

Molecule	Trot	N(X)	N(X) Non-LTE	Size
	(K)	(cm−2)	(cm−2)	('')
H2	⋯	1.1 × 1024	⋯	0.85 c
H13CO+	35	(7.0 ± 0.1) × 1012	(7.6 ± 0.9) × 1012	>1.5 d
HC18O+	35	(5.8 ± 0.9) × 1011	(5.7 ± 0.6) × 1011	1.18 d
CCH	35	(2.7 ± 0.1) × 1014	⋯	>2 d
c-C3H2	35	(9.5 ± 2.2) × 1013	(8.2 ± 0.9) × 1013 a	>1 d
H2CO	39	(1.1 ± 0.1) × 1014	(1.3 ± 0.1) × 1014 a	>1.5 d
HDCO	39	(5.1 ± 0.3) × 1013	⋯	>1 d
D2CO	${{39}}_{-5}^{+6}$	(2.3 ± 0.5) × 1013	⋯ ⋯	>1 d
CN	35	(3.3 ± 0.2) × 1014	(2.5 ± 0.3) × 1014	>2 d
H13CN	35	(1.2 ± 0.1) × 1013	(1.1 ± 0.1) × 1013	0.92 d
HC15N	35	(6.3 ± 0.2) × 1012	(5.8 ± 0.6) × 1012	0.75 d
HC3N	200	(2.7 ± 0.3) × 1013	(2.1 ± 0.2) × 1013	0.65
NO	35	(9.0 ± 2.5) × 1014	(8.9 ± 0.9) × 1014	>1.5 d
HNCO	${{237}}_{-15}^{+17}$	(3.0 ± 0.2) × 1014	⋯	0.54
CH3CN	${{279}}_{-11}^{+12}$	(1.8 ± 0.1) × 1014	(8.6 ± 0.8) × 1013	0.51
13CH3CN	200	<5 × 1012	⋯	⋯
C2H5CN	${{130}}_{-15}^{+20}$	(6.3 ± 1.7) × 1013	⋯	0.52
NH2CHO	${{140}}_{-7}^{+8}$	(4.2 ± 0.7) × 1013	⋯	0.56
SiO	35	(2.5 ± 0.2) × 1012	(2.5 ± 0.3) × 1012	0.65
CS	36	(1.5 ± 0.2) × 1014	(2.0 ± 0.3) × 1014	>1.5
C34S	36	(3.1 ± 0.1) × 1013	⋯	0.70
C33S	${{36}}_{-3}^{+4}$	(1.5 ± 0.2) × 1013	⋯	0.61
OCS	${{106}}_{-5}^{+6}$	(6.5 ± 0.5) × 1014	(6.4 ± 0.7) × 1014	0.55
13OCS	200	(8.7 ± 2.4) × 1013	⋯	0.45
H2CS	${{43}}_{-2}^{+3}$	(1.5 ± 0.1) × 1014	(1.4 ± 0.2) × 1014 a	0.62
SO	${{35}}_{-1}^{+1}$	(4.0 ± 0.3) × 1014	(4.5 ± 0.5) × 1014	0.70 d
34SO	35	(5.9 ± 0.1) × 1013	⋯	0.66
33SO	35	(1.1 ± 0.1) × 1013	⋯	0.53
SO2	${{166}}_{-5}^{+5}$	(1.2 ± 0.1) × 1015	⋯	0.53
34SO2	166	(5.9 ± 0.9) × 1013	⋯	0.51
CH3SH	200	<3 × 1014	⋯	⋯
HDO	${{217}}_{-12}^{+14}$	(2.2 ± 0.2) × 1015	⋯	0.52
CH3OH	${{245}}_{-4}^{+4}$	(1.9 ± 0.1) × 1016	(2.6 ± 0.1) × 1016 b	0.51
13CH3OH	${{181}}_{-9}^{+10}$	(2.8 ± 0.2) × 1015	⋯	0.46
CH2DOH	${{155}}_{-15}^{+18}$	(4.6 ± 0.3) × 1015	⋯	0.52
HCOOCH3	${{181}}_{-5}^{+6}$	(8.6 ± 0.4) × 1015	⋯	0.51
CH3OCH3	${{137}}_{-4}^{+5}$	(2.6 ± 0.1) × 1015	⋯	0.52
C2H5OH	${{136}}_{-12}^{+14}$	(9.6 ± 1.3) × 1014	⋯	0.50
CH3CHO	${{192}}_{-34}^{+52}$	(6.4 ± 0.8) × 1014	⋯	0.49
trans-HCOOH	${{71}}_{-9}^{+11}$	(2.7 ± 0.6) × 1014	⋯	0.58
cis-HCOOH	${{69}}_{-21}^{+50}$	(2.4 ± 1.2) × 1013	⋯	0.49
H2CCO	${{92}}_{-11}^{+14}$	(1.0 ± 0.2) × 1014	⋯	0.55
c-C2H4O	200	(8.9 ± 2.0) × 1013	⋯	0.47

Notes. For T_rot and N(X), those derived from rotation diagrams are shown in italics. Uncertainties and upper limits are of the 2σ level and do not include systematic errors due to adopted spectroscopic constants. See Sections 3.3.1–3.3.3 and 4.2 for details.

^a Assuming ortho/para ratio of 3. ^b Assuming E-CH₃OH/A-CH₃OH ratio of unity (Wirström et al. 2011). ^c Size of continuum emission. ^d Associated with the extended component.

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Most molecular species are well fitted by a single temperature component. Data points in diagrams of CH₃CN and C₂H₅CN are relatively scattered. For CH₃OH, CH₃CN, HNCO, SO₂, and HCOOCH₃, transitions with relatively large S μ² values at low E_u (<300 K) are excluded from the fit in order to avoid the possible effect of optical thickness (see gray points in Figure 6). Adapted threshold values are log S μ² > 1.1 for CH₃OH, log S μ² > 2.4 for CH₃CN, log S μ² > 1.6 for HNCO, log S μ² > 1.2 for SO₂, and log S μ² > 1.8 for HCOOCH₃.

Complex organic molecules, HDO, and SO₂ show high rotation temperatures (>130 K). This suggests that they originate from a warm region associated with a protostar. On the other hand, C³³S and D₂CO, and H₂CS show lower temperatures, suggesting that they arise from a colder region in the outer part of the protostellar envelope. SO also shows a low rotation temperature. Its T_rot is close to that of C³³S. However, SO lines are often optically thick in dense cores, particularly for low-E_u lines, thus the derived rotation temperature would be an upper limit.

3.3.2. Column Densities of Other Molecules

3.3.3. Column Density of H₂, Dust Extinction, and Gas Mass

A column density of molecular hydrogen (${N}_{{{\rm{H}}}_{2}}$) is estimated from the dust continuum data. We use the following equation to calculate ${N}_{{{\rm{H}}}_{2}}$ based on the standard treatment of optically thin dust emission:

$$\begin{eqnarray}&&{N}_{{{\rm{H}}}_{2}}=\displaystyle \frac{{F}_{\nu }/{\rm{\Omega }}}{2{\kappa }_{\nu }{B}_{\nu }({T}_{d})Z\mu {m}_{{\rm{H}}}},\end{eqnarray} \tag{ 3 }$$

where F_ν /Ω is the continuum flux density per beam solid angle as estimated from the observations, κ_ν is the mass absorption coefficient of dust grains coated by thin ice mantles at 1200/870 μm as taken from Ossenkopf & Henning (1994), and we here use 1.07 cm² g⁻¹ for 1200 μm and 1.90 cm² g⁻¹ for 870 μm, T_d is the dust temperature and B_ν (T_d ) is the Planck function, Z is the dust-to-gas mass ratio, μ is the mean atomic mass per hydrogen (1.41, according to Cox 2000), and m_H is the hydrogen mass. We use the dust-to-gas mass ratio of 0.002, which is obtained by scaling the Galactic value of 0.008 by the metallicity of the WB 89–789 region.

A line of sight toward a hot core contains dust grains with different temperatures because of the temperature gradient in a protostellar envelope. Representative dust temperatures (i.e., mass-weighted average temperature) would fall somewhere in between that of a warm inner region and a cold outer region. Shimonishi et al. (2020) presented a detailed analysis of effective dust temperature in the sight line of a low-metallicity hot core in the LMC, based on a comparison of ${N}_{{{\rm{H}}}_{2}}$ derived by submillimeter dust continuum with the above method, model fitting of spectral energy distributions (SEDs), and the 9.7 μm silicate dust absorption depth. The paper concluded that T_d = 60 K for the dust continuum analysis yields an ${N}_{{{\rm{H}}}_{2}}$ value that is consistent with those obtained by other different methods. This temperature corresponds to an intermediate value between a cold gas component (∼50 K) represented by SO and a warm component (∼150 K) represented by CH₃OH and SO₂ in this LMC hot core. The present hot core, WB 89–789 SMM1, harbors similar temperature components as discussed in Sections 3.3.1 and 3.3.2. We thus applied T_d = 60 K for the present source. The continuum brightness of SMM1 is measured to be 11.33 ± 0.05 mJy beam⁻¹ for 1200 μm and 28.0 ± 0.2 mJy beam⁻¹ for 870 μm (3σ uncertainty). Based on the above assumption, we obtain ${N}_{{{\rm{H}}}_{2}}=1.6\times {10}^{24}$ cm⁻² for 1200 μm and ${N}_{{{\rm{H}}}_{2}}=1.2\times {10}^{24}$ cm⁻² for the 870 μm. The ${N}_{{{\rm{H}}}_{2}}$ value changes by a factor of up to 1.6 when the assumed T_d is varied between 40 and 90 K.

Alternatively, a column density of molecular hydrogen can be determined by the model fitting of the observed SED. The best-fit SED discussed in Section 4.1 yields A_V = 184 mag. We here use a standard value of N_H/E(B − V) = 5.8 × 10²¹ cm⁻² mag⁻¹ (Draine 2003) and a slightly high A_V /E(B − V) ratio of 4 for dense clouds (Whittet et al. 2001). Taking into account a factor of 4 lower metallicity, we obtain ${N}_{{{\rm{H}}}_{2}}$/A_V = 2.9 × 10²¹ cm⁻² mag⁻¹, where we assume that all the hydrogen atoms are in the form of H₂. Using this conversion factor, we obtain ${N}_{{{\rm{H}}}_{2}}=5.3\times {10}^{23}$ cm⁻². This ${N}_{{{\rm{H}}}_{2}}$ is similar to the ${N}_{{{\rm{H}}}_{2}}$ derived from the aforementioned method assuming T_d = 150 K. Such T_d may be somewhat high as a typical dust temperature in the line of sight, but it is not a very unrealistic value given the observed temperature range of molecular gas toward WB 89–789 SMM1.

In this paper, we use ${N}_{{{\rm{H}}}_{2}}=1.1\times {10}^{24}$ cm⁻² as a representative value, which corresponds to the average of ${N}_{{{\rm{H}}}_{2}}$ derived by the dust continuum data and the SED fitting. This ${N}_{{{\rm{H}}}_{2}}$ corresponds to A_V = 380 mag using the above conversion factor. Assuming the source diameter of 0.026 pc and a uniform spherical distribution of gas around a protostar, we estimate the gas number density to be ${n}_{{{\rm{H}}}_{2}}=2.1\times {10}^{7}$ cm⁻³, where the total gas mass of 13 M_☉ is enclosed. Similarly, the mass for a 0.1 pc diameter region (i.e., a canonical size of dense cores) is estimated to be 75 M_☉ with T_d = 60 K, where Band 6 and Band 7 estimates are averaged. For the whole field shown in Figures 4–5, which roughly corresponds to a 0.5 pc diameter region, the total mass is estimated to be 800–2500 M_☉, where we assume T_d = 20–10 K for extended dust emission. Note that this is a lower limit because the maximum recoverable scale of the present observations is 54 (0.28 pc).

3.3.4. Fractional Abundances and Isotope Abundance Ratios

Fractional abundances with respect to H₂ are shown in Table 4, which are calculated based on column densities estimated in Sections 3.3.1–3.3.3. The fractional abundances normalized by the CH₃OH column density are also discussed in Sections 4.3–4.4 because of the nonnegligible uncertainty associated with ${N}_{{{\rm{H}}}_{2}}$ (see Section 3.3.3).

Table 4. Estimated Fractional Abundances

Molecule	N(X)/${N}_{{{\rm{H}}}_{2}}$
	0.026 pc area	0.1 pc area
HCO+ a	(9.5 ± 3.2) × 10−10	(1.5 ± 0.3) × 10−9
H2CO	(1.0 ± 0.3) × 10−10	(1.2 ± 0.1) × 10−10
HDCO	(4.7 ± 1.3) × 10−11	(3.9 ± 0.2) × 10−11
D2CO	(2.1 ± 0.7) × 10−11	(2.0 ± 0.3) × 10−11
C2H	(2.5 ± 0.7) × 10−10	(5.8 ± 1.2) × 10−10
c-C3H2	(8.6 ± 3.1) × 10−11	(5.9 ± 1.2) × 10−11
CN	(3.0 ± 0.8) × 10−10	(6.6 ± 1.3) × 10−10
HCN a	(1.7 ± 0.6) × 10−9	(1.2 ± 0.3) × 10−9
HC3N	(2.5 ± 0.7) × 10−11	(1.4 ± 0.1) × 10−11
NO	(8.1 ± 3.2) × 10−10	(1.6 ± 0.1) × 10−9
HNCO	(2.7 ± 0.8) × 10−10	(7.1 ± 0.6) × 10−11
CH3CN b	(4.2 ± 2.7) × 10−10	(3.7 ± 2.8) × 10−10
C2H5CN	(5.8 ± 2.2) × 10−11	(2.4 ± 0.9) × 10−11
NH2CHO	(3.8 ± 1.2) × 10−11	(1.8 ± 0.1) × 10−11
SiO	(2.2 ± 0.6) × 10−12	(1.2 ± 0.1) × 10−12
CS c	(9.7 ± 3.3) × 10−10	(6.4 ± 1.3) × 10−10
SO c	(1.9 ± 0.5) × 10−9	(1.3 ± 0.3) × 10−9
OCS a	(1.2 ± 0.5) × 10−8	(4.1 ± 1.4) × 10−9
H2CS	(1.4 ± 0.4) × 10−10	(9.0 ± 1.0) × 10−11
SO2	(1.1 ± 0.3) × 10−9	(2.9 ± 0.1) × 10−10
CH3SH	<3 × 10−10	<2 × 10−10
HDO	(2.0 ± 0.6) × 10−9	(7.7 ± 0.9) × 10−10
CH3OH a	(3.8 ± 1.3) × 10−7	(1.7 ± 0.3) × 10−7
CH2DOH	(4.2 ± 1.2) × 10−9	(1.5 ± 0.2) × 10−9
HCOOCH3	(7.8 ± 2.2) × 10−9	(3.0 ± 0.2) × 10−9
CH3OCH3	(2.3 ± 0.6) × 10−9	(1.0 ± 0.1) × 10−9
C2H5OH	(8.7 ± 2.7) × 10−10	(3.3 ± 0.8) × 10−10
CH3CHO	(5.8 ± 1.8) × 10−10	(2.1 ± 0.4) × 10−10
HCOOH d	(2.7 ± 1.0) × 10−10	(1.2 ± 0.4) × 10−10
H2CCO	(9.2 ± 3.0) × 10−11	(3.7 ± 0.9) × 10−11
c-C2H4O	(8.1 ± 2.8) × 10−11	(5.9 ± 1.2) × 10−11

Notes. Uncertainties and upper limits are of the 2σ level. Column densities of molecules for a 0.026 pc area are summarized in Table 3. An empirical uncertainty of 30% is assumed for ${N}_{{{\rm{H}}}_{2}}$.

^a Estimated from the ¹³C isotopologue with ¹²C/¹³C = 150. ^b Rotation diagram analysis of CH₃CN is used to derive a lower limit and the nondetection of ¹³CH₃CN for an upper limit. ^c Estimated from the ³⁴S isotopologue with ³²S/³⁴S = 35. ^d Sum of trans- and cis- species.

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Abundances of HCO⁺, HCN, SO, CS, OCS, and CH₃OH are estimated from their isotopologues, H¹³CO⁺, H¹³CN, ³⁴SO, C³⁴S, O¹³CS, and ¹³CH₃OH. Detections of isotopologue species for SO, CS, OCS, and CH₃OH imply that the main species would be optically thick. Isotope abundance ratios of ¹²C/¹³C = 150 and ³²S/³³S = 35 are assumed, which are obtained by extrapolating the relationship between isotope ratios and galactocentric distances reported in Wilson & Rood (1994) and Humire et al. (2020) to D_GC = 19 kpc.

Abundance ratios are derived for several rare isotopologues; we obtain CH₂DOH/CH₃OH = 0.011 ± 0.002, D₂CO/HDCO = 0.45 ± 0.10, ³⁴SO/³³SO =5 ± 1, C³⁴S/C³³S = 2 ± 1, and ³²SO₂/³⁴SO₂ = 20 ± 4. The ³²SO₂/³⁴SO₂ ratio in WB 89–789 SMM1 is similar to the solar ³²S/³⁴S ratio (22, Wilson & Rood 1994), although we expect a slightly higher value in the outer Galaxy due to the ³²S/³⁴S gradient in the Galaxy (Chin et al. 1996; Humire et al. 2020). Astrophysical implications for the deuterated species are discussed in Section 4.4.

The rotation diagram of CH₃CN is rather scattered. Although its isotopologue line is not detected, optical thickness might affect the column density estimate, as CH₃CN is often optically thick in hot cores (e.g., Fuente et al. 2014). To obtain a possible range of its column density, we use the rotation diagram of ¹²CH₃CN data to estimate a lower limit and the nondetection of the ¹³CH₃CN(19₀–18₀) line at 339.36630 GHz (E_u = 163 K) for an upper limit.

We have also repeated the analysis for the spectra extracted from a 0.1 pc (193) diameter region at the hot-core position, for the sake of comparison with LMC hot cores (see Section 4.4). Those abundances are also summarized in Table 4. The abundances for a 0.1 pc area do not drastically vary from those for a 0.026 pc area. Molecules with a compact spatial distribution (e.g., COMs) tend to decrease their abundances by a factor of ∼2–3 in the 0.1 pc data due to the beam dilution effect. In contrast, those with extended spatial distributions and intensity peaks outside the hot-core region (H¹³CO⁺, CCH, CN, and NO) increase by a factor of ∼2 in the 0.1 pc data.

The nature of WB 89–789 SMM1 is characterized by (i) the compact distribution of warm gas (∼0.03 pc, see Section 4.2), (ii) the high gas temperature that can trigger the ice sublimation (≥100 K, Section 3.3.1), (iii) the high density (2 × 10⁷ cm⁻³, Section 3.3.3), (iv) the association with a luminous protostar (see below), and (v) the presence of chemically rich molecular gas. Those properties suggest that the source is associated with a hot molecular core.

Figure 7 shows an SED of SMM1, where the data are collected from available databases and literature (Brand & Wouterloot 2007; Wright et al. 2010; Yamamura et al. 2010). The bolometric luminosity of the source is estimated to be 8.4 × 10³ L_☉ based on the SED fitting with the model of Robitaille et al. (2007). This luminosity is equivalent to a stellar mass of about 10 M_☉ according to the mass–luminosity relationship of zero-age main-sequence (ZAMS) stars (Zinnecker & Yorke 2007).

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Note that far-infrared data, which is important for the luminosity determination of embedded sources, is insufficient for SMM1. Only upper limits are provided due to the low angular resolution of available AKARI FIS all-sky survey data. Thus, the derived luminosity (and therefore mass) may be lower than the current estimate. Future high-spatial-resolution infrared observations in those missing wavelengths are highly required.

Alternatively, we can estimate the luminosity of SMM1 by scaling the luminosity of a low-metallicity LMC hot core, ST16, whose SED is well determined based on a comprehensive infrared data set from 1 to 1200 μm (Shimonishi et al. 2020). This LMC hot core has a total luminosity of 3.1 × 10⁵ L_☉ and a K_s -band magnitude ([K_s ]) of 13.4 mag at 50 kpc, while SMM1 has [K_s ] = 14.7 mag at 10.7 kpc. Scaling the luminosity of ST16 with the distance and K_s -band magnitude, we obtain 4.3 × 10³ L_☉ for SMM1, which is a factor of 2 lower than the estimate by the SED fitting. In either case, present estimates suggest that the luminosity of SMM1 would correspond to the lower end of high-mass ZAMS or upper end of intermediate-mass ZAMS.

The observed emission lines and continuum show different spatial distributions depending on species. Those distributions have important clues to understand their origins. A schematic illustration of the temperature structure and molecular gas distribution in WB 89–789 SMM1 is shown in Figure 8 based on the discussion in this section.

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We have estimated the spatial extent of the observed emission by fitting a two-dimensional Gaussian to the continuum center (Table 3). Compact distributions (FWHM =05–06, 0.026–0.031 pc), comparable with the beam size, are seen in HDO, COMs, CH₃CN, HNCO, OCS, and high-excitation SO₂ lines. HC₃N is slightly extended (FWHM = 065). They are concentrated at the hot-core position, suggesting that they originate from a warm region where ice mantles are sublimated.

SO, ³⁴SO, ³³SO, and low-excitation SO₂ show relatively compact distributions (FWHM = 05–07, 0.026–0.036 pc) at the hot-core position, but also show a secondary peak at the south side of the hot core. This secondary peak coincides with the peak of the NO emission. Other sulfur-bearing species such as C³⁴S, C³³S, and H₂CS show compact distributions (FWHM = 06–00.7, 0.031–0.052 pc) centered at the hot core.

A characteristic distribution that is symmetric to the hot-core position is seen in SiO. It shows a compact emission (FWHM = 065) at the hot-core center, but also shows other peaks at the northeast and southwest sides of the hot core. Those secondary peaks are slightly elongated. SiO is a well-known shock tracer. The observed structure would have originated from the shocked gas produced by bipolar protostellar outflows. A driving source of the outflows would be a protostar embedded in a hot core as the distribution of SiO is symmetric with the hot-core position.

Even extended distributions (FWHM > 10) are seen in CN, CCH, H¹³CO⁺, HC¹⁸O⁺, H¹³CN, HC¹⁵N, NO, CS, H₂CO, and HDCO, D₂CO, and low-excitation CH₃OH. Gas-phase reactions and nonthermal desorption of icy species would have a nonnegligible contribution to the formation of those species because they are widely distributed beyond the hot core. We note that dust continuum, H¹³CN, HC¹⁵N have a moderately sharp peak (FWHM < 10) at the hot-core position in addition to the extended component. c-C₃H₂ shows a patchy distribution, whose secondary peak at the southwest of the hot core does not coincide with those of other species.

Molecular radicals (CN, CCH, and NO) do not have their emission peak at the hot-core position. This would suggest that the chemistry outside the hot-core region largely contributes to their production. CN and CCH are known to be abundant in photodissociation regions (PDRs) because atomic carbon is efficiently provided by the photodissociation of CO under moderate UV fields (e.g., Fuente et al. 1993; Jansen et al. 1995; Sternberg & Dalgarno 1995; Rodriguez-Franco et al. 1998; Pety et al. 2017). In the present source, their emission shows a similar spatial distribution. A similar distribution between CN and CCH has been also observed in an LMC hot core by Shimonishi et al. (2020); they argue that CN and CCH would trace PDR-like outflow cavity structures that are irradiated by the UV light from a protostar associated with a hot core. We speculate that this is also the case for WB 89–789 SMM1.

Figure 9 shows the velocity maps (moment 1) of CN and CCH lines. CN and CCH emission are elongated in the southwest direction from the hot core (see also Figure 4). The figure also shows a possible direction of protostellar outflows expected from the spatial distribution of SiO. The elongated directions of CN and CCH coincide with the inferred direction of outflows. In addition, the elongated southwest parts of CN and CCH are blueshifted by ∼1–2 km s⁻¹ compared to the hot-core position. This may be due to outflow gas motion, although CN and CCH would trace an outflow cavity wall rather than the outflow gas itself. Actually, the observed velocity shift is smaller than a typical value of high-velocity wing components in massive protostellar outflows (≥5 km s⁻¹; e.g., Beuther et al.2002; Maud et al. 2015). We note that a clear velocity structure is not seen in the SiO velocity map, except for the position of another embedded protostar discussed in Section 4.5. Future observations of optically thick outflow tracers such as CO are necessary to confirm the presence of high-velocity gas associated with protostellar outflows.

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Figure 10 shows a comparison of molecular abundances between WB 89–789 SMM1 and other known Galactic hot cores. The data for an intermediate-mass hot core, NGC 7192 FIRS2, are adopted from Fuente et al. (2014). The abundances are based on the 220 GHz region observations for a 0.009 pc diameter area centered at the hot core. The luminosity of NGC 7192 FIRS2 (∼500 L_☉) corresponds to that of a 5 M_☉ ZAMS. The data for a high-mass source, the Orion hot core, are adopted from Sutton et al. (1995), based on the 340 GHz region observations for a 0.027 pc diameter area at the hot core. The abundance of HNCO is taken from Schutte & Greenberg (1997).

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The molecular abundances in WB 89–789 SMM1 are generally lower than those of the inner Galactic counterparts. The degree of the abundance decrease is roughly consistent with the lower metallicity of the WB 89–789 region as indicated by the scale bar in Figure 10. Particularly, SMM1 and the intermediate-mass hot core NGC 7192 FIRS2 show similar molecular abundances after taking into account the four times lower metallicity of the former source. For the comparison with Orion, it seems that HC₃N, C₂H₅CN, and SO₂ are significantly less abundant in SMM1 even taking into account the lower metallicity, while CH₃OH is overabundant in SMM1 despite the low metallicity.

To further focus on chemical complexity at low metallicity, Figure 11 shows a comparison of fractional abundances of COMs normalized by the CH₃OH column density for WB 89–789 SMM1 and NGC 7192 FIRS2. Such a comparison is useful for investigating the chemistry of organic molecules in warm and dense gas around protostars (Herbst & van Dishoeck 2009; Drozdovskaya et al. 2019) because CH₃OH is believed to be a parental molecule for the formation of even larger COMs (e.g., Nomura & Millar 2004; Garrod & Herbst 2006). In addition, CH₃OH is a product of grain surface reaction, thus warm CH₃OH gas mainly arises from a high-temperature region, where ices are sublimated and characteristic hot-core chemistry proceeds. Furthermore, the normalization by CH₃OH can cancel the metallicity effect in the abundance comparison.

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The N(X)/N(CH₃OH) ratios are remarkably similar between WB 89–789 SMM1 and NGC 7192 FIRS2 as shown in Figure 11(a). The ratios of SMM1 coincide with those of NGC 7192 FIRS2 within a factor of 2 for the most molecular species. The correlation coefficient is calculated to be 0.94, while it is 0.96 if CH₃CN is excluded. It seems that CH₃CN deviates from the overall trend, although the uncertainty is large due to the opacity effect (see 3.3.4). C₂H₅OH also slightly deviates from the trend. The reason for their behavior is still unclear, but it may be related to the formation pathway of those molecules.

The above two comparisons suggest that chemical compositions of the hot core in the extreme outer Galaxy scale with the metallicity. In the WB 89–789 region, the metallicity is expected to be four times lower compared to the solar neighborhood. The observed abundances of COMs in the SMM1 hot core are lower than the other Galactic hot cores, but the decrease is proportional to this metallicity. Furthermore, similar N(COMs)/N(CH₃OH) ratios suggest that CH₃OH is an important parental species for the formation of larger COMs in a hot core, as suggested by the aforementioned theoretical studies.

CH₃OH ice is believed to form on grain surfaces, and several formation processes are proposed by laboratory experiments; i.e., hydrogenation of CO, ultraviolet photolysis, and radiolysis of ice mixtures (e.g., Hudson & Moore 1999; Watanabe et al. 2007). It is known that CH₃OH is already formed in quiescent prestellar cores before star formation occurs (Boogert et al. 2011). Solid CH₃OH will chemically evolve to larger COMs through a combination of photolysis, radiolysis, and grain heating during the warm-up phase that leads to the formation of a hot core (Garrod & Herbst 2006). High-temperature gas-phase chemistry of sublimated CH₃OH would also contribute to the COM formation (Nomura & Millar 2004; Taquet et al. 2016). The present results suggest that various COMs can form even in a low-metallicity environment, if their parental molecule, CH₃OH, is efficiently produced in a star-forming core. The detection of a chemically rich star-forming core in the extreme outer Galaxy has an impact on the understanding of the occurrence of the chemical complexity in a primordial environment of the early phase of the Galaxy formation. We here note that observations of ice mantle compositions have not been reported for the outer Galaxy so far. Future infrared observations of ice absorption bands toward embedded sources in the outer Galaxy are important.

It is still unknown if the observed simply metallicity-scaled chemistry of COMs in the WB 89–789 SMM1 hot core is common in other hot-core sources in the outer Galaxy. A comparison of the present data with those of hot cores in the LMC would provide a hint for understanding the universality of low-metallicity hot-core chemistry. The metallicity of the LMC is reported to be lower than the solar value by a factor of 2 to 3 (e.g., Dufour et al. 1982; Westerlund 1990; Russell & Dopita 1992; Choudhury et al. 2016), which is in common with the outer Galaxy.

Figure 12 shows a comparison of molecular abundances between WB 89–789 SMM1 and three LMC hot cores. The plotted molecular column densities for LMC hot cores are adopted from Shimonishi et al. (2016a) for ST11, Shimonishi et al. (2020) for ST16, and Sewiło et al. (2018) from N113 A1. Another LMC hot core in Sewiło et al. (2018), N113 B3, has similar molecular abundances to N113 A1. The ${N}_{{{\rm{H}}}_{2}}$ value of ST11 and N113 A1 is reestimated using the same dust opacity data and dust temperature (T_d = 60 K) as in this work; we obtained ${N}_{{{\rm{H}}}_{2}}=1.2\times {10}^{24}$ cm⁻² for ST11 and ${N}_{{{\rm{H}}}_{2}}\,=9.2\times {10}^{23}$ cm⁻² for N113 A1. The dust temperature assumed in ST16 is 60 K as described in Section 3.3.3. Molecular column densities are estimated for circular/elliptical regions of 0.12 × 0.12 pc, 0.10 × 0.10 pc, and 0.21 × 0.13 pc for ST11, ST16, and N113 A1, respectively. For a fair comparison, we have recalculated ${N}_{{{\rm{H}}}_{2}}$ and molecular column densities of SMM1 for a 0.1 pc (193) diameter region. Those abundances are plotted in Figure 12 and summarized in Table 4.

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The chemical composition of the outer Galaxy hot core does not resemble those of LMC hot cores as seen in Figure 12. The dissimilarity is also seen in the N(X)/N(CH₃OH) comparison between SMM1 and ST16 as shown in Figure 11(b), where the correlation coefficient is calculated to be 0.69.

Shimonishi et al. (2020) argue that SO₂ will be a good tracer of low-metallicity hot-core chemistry because (i) it is commonly detected in LMC hot cores with similar abundances, and (ii) it originated from a compact hot-core region. SO also shows similar abundances within LMC hot cores. In WB 89–789 SMM1, however, the abundances of SO₂ and SO relative to H₂ are lower by a factor of 28 and 5 compared with LMC hot cores. The measured rotation temperatures of SO₂ are similar between those hot cores, i.e., 166 K (SO₂) for SMM1, 232 K (SO₂) and 86 K (³⁴SO₂) for ST16, 190 K (SO₂) and 95 K (³⁴SO₂) for ST11. The SO₂ column densities for ST16 and ST11 are estimated from ³⁴SO₂, while that for SMM1 is from SO₂. However, the SO₂ column density of SMM1 increases only by a factor of up to 3 when it is estimated from ³⁴SO₂ (see Section 3.3.4). Thus, the low SO₂ abundance in the outer Galactic hot core would not be due to the optical thickness.

In contrast to the S–O bond-bearing species, the C–S bond-bearing species such as CS, H₂CS, and OCS do not show a significant abundance decrease in WB 89–789 SMM1. Thus, it is not straightforward to attribute the low abundance of SO₂ (and perhaps SO) to the low elemental abundance ratio of sulfur in the outer Galaxy. Hot-core chemistry models suggest that SO₂ is mainly produced in high-temperature gas-phase reactions in warm gas, using H₂S sublimated from ice mantles (Charnley 1997; Nomura & Millar 2004). This also applies to the SO₂ formation in low-metallicity sources as shown in astrochemical simulations for LMC hot cores (Shimonishi et al.2020). We speculate that the different behavior of SO₂ in the outer Galaxy and LMC hot cores may be related to differences in the evolutionary stage of hot cores. A different luminosity of host protostars may also contribute to the different sulfur chemistry, i.e., ∼8 × 10³ L_☉ for WB 89–789 SMM1, while several ×10⁵ L_☉ for LMC hot cores. A different cosmic-ray ionization rate between the outer Galaxy and the LMC may also affect the chemical evolution, although the rate is not known for the outer Galaxy.

Among nitrogen-bearing molecules, NO shows interesting behavior in LMC hot cores. After correcting for the metallicity, NO is overabundant in LMC hot cores compared with Galactic counterparts despite the low elemental abundance of nitrogen in the LMC (Shimonishi et al. 2020). Only NO shows such behavior among the nitrogen-bearing molecules observed in LMC hot cores. In WB 89–789 SMM1, however, such an overabundance of NO is not observed. The NO abundance of SMM1 is 1.6 × 10⁻⁹ for a 0.1 pc region data. This is a factor of 5 lower than a typical NO abundance in Galactic high-mass hot cores (8×10⁻⁹, Ziurys et al. 1991), which is consistent with a factor of 4 lower metallicity in WB 89–789. The present high-spatial-resolution data have revealed that NO does not mainly arise from a hot-core region, as shown in Figure 4. It has an intensity peak at the south part of the hot core, where low-excitation lines of SO and SO₂ also have a secondary peak (Section 4.2). Thus, shock chemistry or photochemistry, rather than high-temperature chemistry, would contribute to the production of NO in low-metallicity protostellar cores. In that case, a lower luminosity of SMM1 than those of LMC hot cores may contribute to the different behavior of NO.

For other nitrogen-bearing molecules, HNCO and CH₃CN, a clear difference is not identified between outer Galactic and LMC hot cores, although the number of data points is limited and the abundance uncertainty is large. The reason for the unusually low abundance of SiO in SMM1 is unknown. It may be related to different shock conditions or grain compositions because dust sputtering by shock is mainly responsible for the production of SiO gas.

The formation of COMs is one of the important standpoints for low-metallicity hot-core chemistry. It is reported that CH₃OH shows a large abundance variation in LMC hot cores (Shimonishi et al. 2020). There are organic-poor hot cores such as ST11 and ST16, while N113 A1 and B3 are organic rich. The CH₃OH abundance of WB 89–789 SMM1 is higher than those of any known LMC hot cores. The abundances of HCOOCH₃ and CH₃OCH₃ in SMM1 are comparable with those of an organic-rich LMC hot core, N113 A1. The detection of many other COMs in SMM1 suggests the source has experienced rich organic chemistry despite its low-metallicity nature.

Astrochemical simulations for LMC hot cores suggest that dust temperature at the initial ice-forming stage has a major effect on the abundance of CH₃OH gas in the subsequent hot-core stage (Acharyya & Herbst 2018; Shimonishi et al. 2020). Simulations of grain surface chemistry dedicated to the LMC environment also suggest that dust temperature is one of the key parameters for the formation of CH₃OH in dense cores (Acharyya & Herbst 2015; Pauly & Garrod 2018). This is because (i) CH₃OH is mainly formed by the grain surface reaction, and (ii) the hydrogenation reaction of CO, which is a dominant pathway for the CH₃OH formation, is sensitive to the dust temperature due to the high volatility of atomic hydrogen. For this reason, it is inferred that organic-rich hot cores had experienced a cold stage (≲10K) that is sufficient for CH₃OH formation before the hot-core stage, while organic-poor ones might have missed such a condition for some reason. Alternatively, the slight difference in the hot core's evolutionary stage may contribute to the CH₃OH abundance variation, because the high-temperature gas-phase chemistry is rapid, and it can decrease CH₃OH gas at a late stage (e.g., Nomura & Millar 2004; Garrod & Herbst 2006; Vasyunin & Herbst 2013; Balucani et al. 2015).

Low-metallicity hot-core chemistry simulations in Shimonishi et al. (2020) argue that the maximum achievable abundances of CH₃OH gas in a hot-core stage significantly decrease as the visual extinction of the initial ice-forming stage decreases. On the other hand, the simulations show that the CH₃OH gas abundance is simply metallicity scaled if the initial ice-forming stage is sufficiently shielded. In a well-shielded initial condition, the grain surface is cold enough to trigger CO hydrogenation, and the resultant CH₃OH abundance is roughly regulated by the elemental abundances. The observed metallicity-scaled chemistry of COMs in WB 89–789 SMM1 implies that the source had experienced such an initial condition before the hot-core stage.

Deuterium chemistry is widely used in interpreting the chemical and physical history of interstellar molecules (e.g., Caselli & Ceccarelli 2012; Ceccarelli et al. 2014). The measured CH₂DOH/CH₃OH ratio in WB 89–789 SMM1 is 1.1% ± 0.2%, which is comparable to the higher end of those ratios observed in high-mass protostars and the lower end of those in low-mass protostars (e.g., see Figure 2 in Drozdovskaya et al. 2021). The ratio is orders of magnitude higher than the deuterium-to-hydrogen ratio in the solar neighborhood (2×10⁻⁵; Linsky et al. 2006; Prodanović et al. 2010) and that in the big bang nucleosynthesis (3 × 10⁻⁵; Burles 2002 and references therein). This suggests that the efficient deuterium fractionation occurred upon the formation of CH₃OH in SMM1. The D₂CO/HDCO ratio is 45% ± 10%, which is comparable to those observed in low-mass and high-mass protostars (e.g., Zahorecz et al. 2021). This would suggest that physical conditions for deuterium fractionation could be similar between WB 89–789 SMM1 and inner Galactic protostars. Note that higher-spatial-resolution observations and detailed multiline analyses would affect the measured abundance of deuterated species as reported in Persson et al. (2018) for the case of a nearby low-mass protostar. The H₂CO column density derived in this work may be a lower limit because the line is often optically thick, thus we do not discuss the abundance ratio relative to H₂CO.

It is known that the deuterium fractionation efficiently proceeds at low temperature (e.g., Roberts et al. 2003; Caselli & Ceccarelli 2012; Taquet et al. 2014; Furuya et al. 2016). This is because the key reaction for the trigger of deuterium fractionation, ${{\rm{H}}}_{3}^{+}$ + HD → H₂D⁺ + H₂ + 232 K, is exothermic and its backward reaction cannot efficiently proceed below 20 K. In addition, gaseous neutral species such as CO and O efficiently destruct H₂D⁺, thus their depletion at low temperature further enhances the deuterium fractionation (e.g., Caselli & Ceccarelli 2012). A sign of high deuterium fractionation observed in WB 89–789 SMM1 suggests that the source had experienced such a cold environment during its formation. This picture is consistent with the implication obtained from the metallicity-scaled chemistry of COMs, which also suggests the occurrence of a cold and well-shielded initial condition as discussed above.

Although the low metallicity is common between the outer Galaxy and the LMC, their star-forming environments would be different; the LMC has more harsh environments as inferred from active massive star formation over the whole galaxy, while that for the outer Galaxy might be quiescent due to its low star formation activity. Those environmental differences need to be taken into account for further understanding of the chemical evolution of star-forming regions at low metallicity. A future extensive survey of protostellar objects toward the outer Galaxy is thus vitally important for further discussion. Astrochemical simulations dedicated to the environment of the outer Galaxy, and the application to lower-mass protostars, are also important.

We have also detected a compact source associated with high-velocity SiO gas at the east side of WB 89–789 SMM1. Hereafter, we refer to this source as WB 89–789 SMM2. According to the SiO emission, the source is located at R.A. = 06^h17^m24246 and decl. = 14°54'4325 (ICRS), which is 27 (0.14 pc) away from SMM2. Figure 13(a) shows the SiO(6–5) spectrum extracted from a 06 diameter region centered at the above position. The SiO line is largely shifted to the blue and red sides relative to the systemic velocity in a symmetric fashion. The peaks of the shifted emission are located at V_sys ± 25 km s⁻¹.

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Figure 13(b) shows a velocity map and the integrated intensity distribution of SiO(6–5). In the figure, to focus on SiO in WB 89–789 SMM2, the intensity is integrated over a much wider velocity range (0–60 km s⁻¹) compared with that adopted in Figure 4 (31–38 km s⁻¹). The velocity map clearly indicates that the velocity structure of SiO in SMM2 is spatially symmetric to the SiO center. At this position, a local peak is seen in the 1200 μm continuum as shown in the figure, suggesting the presence of an embedded source. SMM2 does not show any emission lines of COMs, and no alternative molecular lines are identified at the frequencies of Doppler-shifted SiO emission. Also taking into account the clear spectral and spatial symmetry, the observed lines must be attributed to high-velocity SiO gas.

The spectral characteristics of the observed high-velocity SiO resemble those of extremely high-velocity (EHV) outflows observed in Class 0 protostars (Bachiller et al. 1991; Tafalla et al. 2010, 2015; Tychoniec et al. 2019). The EHV flows are known to appear as a discrete high-velocity (V ≳ 30 km s⁻¹) peak and are observed in the youngest stage of star formation (Bachiller 1996; Matsushita et al. 2019 and references therein). The EHV flows extend up to several thousands of astronomical units from the central protostar in SiO and usually have collimated bipolar structures (e.g., Bachiller et al. 1991; Hirano et al. 2010; Matsushita et al. 2019; Tychoniec et al. 2019). The beam size of the present data is about 5000 au, thus such structures will not be fully spatially resolved. Actually, a symmetric spatial distribution of blue-/redshifted SiO is only marginally resolved into two beam size regions (Figure 13(b)). The spatial extent of SiO emission is about 1'' (0.052 pc). Assuming an outflow velocity of 25 km s⁻¹, we estimate a dynamical timescale of EHV flows to be at least 2000 yr. This is roughly consistent with the dynamical timescales of other EHV sources, which range from a few hundred to a few thousand years (Bachiller 1996 and references therein).

A 1200 μm continuum flux in a 06 diameter region centered at SMM2 is 0.60 ± 0.05 mJy beam⁻¹. Assuming T_d = 20 K, we obtain ${N}_{{{\rm{H}}}_{2}}=3.2\times {10}^{23}$ cm⁻². This is equivalent to a gas number density of ${n}_{{{\rm{H}}}_{2}}=4.9\times {10}^{6}$ cm⁻³. If we assume a higher T_d , i.e., 40 K, then the derived column density is 2.5 times lower than the 20 K case. In either case, the continuum data suggest the presence of high-density gas at this position. A column density and fractional abundance of SiO gas at the above position is estimated to be N(SiO) ∼ 2 × 10¹³ cm⁻² and N(SiO)/${N}_{{{\rm{H}}}_{2}}\sim 6\times {10}^{-11}$, where we assume optically thin emission in the LTE and the gas/dust temperature of 20 K. The fractional abundance will be two times higher if we assume the gas/dust temperature of 10 K or 40 K. The SiO abundance in SMM2 is at least 30 times higher than that observed in SMM1. The observed enhancement of SiO in SMM2 would be related to shock chemistry triggered by EHV outflows.

Previous single-dish observations of CO detected extended (∼20'') molecular outflows in the WB 89–789 region (Brand & Wouterloot 1994, 2007). The center of the outflow gas coincides with the position of the IRAS source (IRAS 06145+1455; 06^h17^m242, 14°54'42'', J2000). This position is consistent with those of SMM1 or SMM2, given the large beam size of CO(3–2) observations (14'') in Brand & Wouterloot (2007). The observed CO outflow gas has an extended blueshifted component (20 < V_LSR < 31 km s⁻¹) toward the southeast direction from the center, while a redshifted component (37 < V_LSR < 55 km s⁻¹) is extended toward the northwest direction (see Figure 9 in Brand & Wouterloot 2007). This outflow direction coincides with that of the high-velocity SiO outflows observed in this work. The SiO outflows from SMM2 may have a common origin with the large-scale CO outflows.

In summary, it is likely that a compact, high-density, and embedded object is located at the position of WB 89–789 SMM2. Presumably, a protostar associated with SMM2 is driving the observed high-velocity SiO gas flows. Its short dynamical timescale and similarity to EHV flows suggest that the object is at the youngest stage of star formation (Class 0/I). Nondetection of warm gas emission also supports its young nature. We note that the detailed structure of high-velocity SiO gas is not fully spatially resolved, and CO lines, which often trace high-velocity outflows, are not covered in the present data. Future high-spatial-resolution observations of CO and other outflow tracers are key to further clarify the nature of WB 89–789 SMM2.