# ALMA observations of the Extended Green Object G19.01–0.03:
II. A massive protostar with typical chemical abundances surrounded by four
low-mass prestellar core candidates

G. M. Williams,1 C. J. Cyganowski,2 C. L. Brogan,3 T. R. Hunter,3 P. Nazari4
and R. J. Smith5,2
1School of Physics & Astronomy, The University of Leeds, Leeds, LS2 9JT, UK
2Scottish Universities Physics Alliance (SUPA), School of Physics and
Astronomy, University of St Andrews, North Haugh, St Andrews, KY16 9SS, UK
3National Radio Astronomy Observatory (NRAO), 520 Edgemont Rd,
Charlottesville, VA 22903, USA
4Leiden Observatory, Leiden University, PO Box 9513, 2300 RA Leiden, NL
5Jodrell Bank Centre for Astrophysics, Department of Physics and Astronomy,
University of Manchester, Oxford Road, Manchester, M13 9PL, UK E-mail:
<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

We present a study of the physical and chemical properties of the Extended
Green Object (EGO) G19.01$-$0.03 using sub-arcsecond angular resolution
Atacama Large Millimeter/submillimeter Array (ALMA) $1.05$ mm and Karl G.
Jansky Very Large Array (VLA) $1.21$ cm data. G19.01$-$0.03 MM1, the
millimetre source associated with the central massive young stellar object
(MYSO), appeared isolated and potentially chemically young in previous
Submillimeter Array observations. In our $\sim 0.4\arcsec$-resolution ALMA
data, MM1 has four low-mass millimetre companions within $0.12$ pc, all
lacking maser or outflow emission, indicating they may be prestellar cores.
With a rich ALMA spectrum full of complex organic molecules, MM1 does not
appear chemically young, but has molecular abundances typical of high-mass hot
cores in the literature. At the $1.05$ mm continuum peak of MM1,
$\mathrm{N}(\mathrm{CH}_{3}\mathrm{OH})=(2.22\pm 0.01)\times 10^{18}$ cm-2 and
$T_{\mathrm{ex}}=162.7\begin{subarray}{c}+0.3\\\ -0.5\end{subarray}$ K based
on pixel-by-pixel Bayesian analysis of LTE synthetic methanol spectra across
MM1. Intriguingly, the peak CH3OH $T_{\mathrm{ex}}=165.5\pm 0.6$ K is offset
from MM1’s millimetre continuum peak by $0.22\arcsec\sim 880$ au, and a region
of elevated CH3OH $T_{\mathrm{ex}}$ coincides with free-free VLA $5.01$ cm
continuum, adding to the tentative evidence for a possible unresolved high-
mass binary in MM1. In our VLA $1.21$ cm data, we report the first NH3(3,3)
maser detections towards G19.01$-$0.03, along with candidate 25 GHz CH3OH
$5(2,3)-5(1,4)$ maser emission; both are spatially and kinematically
coincident with $44$ GHz Class I CH3OH masers in the MM1 outflow. We also
report the ALMA detection of candidate $278.3$ GHz Class I CH3OH maser
emission towards this outflow, strengthening the connection of these three
maser types to MYSO outflows.

###### keywords:

stars: individual: G19.01–0.03 – stars: formation – stars: massive – stars:
protostars – masers – techniques: interferometric

††pubyear: 2023††pagerange: ALMA observations of the Extended Green Object
G19.01–0.03:  II. A massive protostar with typical chemical abundances
surrounded by four low-mass prestellar core candidates–B

##  1 Introduction

High-mass stars (M${}_{\ast}>8$ M⊙) are influential in the dynamical and
chemical evolution of the interstellar medium (ISM), through their strong
feedback, outflows and jets, and through enrichment of the ISM with heavy
elements (e.g. Peters et al., 2017; Rosen & Krumholz, 2020; Mignon-Risse et
al., 2021; Grudić et al., 2022). Constraining exactly how high-mass stars form
however is hampered by their natal molecular clouds being significantly more
distant ($d>1$ kpc) and more clustered ($n_{\ast}>100$ pc-3) than those of
their low-mass counterparts (M${}_{\ast}<8$ M⊙). Furthermore, the short pre-
main sequence lifetimes ($<1$ Myrs; Mottram et al., 2011) of high-mass stars
ensures that the entirety of their formation is obscured in regions of high
extinction (e.g. Chevance et al., 2020; Kim et al., 2021). Now with the advent
of facilities capable of high-angular resolution observations in the
(sub)millimetre such as the Atacama Large Millimetre/(sub)millimetre Array
(ALMA), we have the ability to resolve and disentangle the thermal emission of
the early stages of high-mass star formation.

The core-fed theory of massive star formation (McKee & Tan, 2003; Tan et al.,
2014) describes the monolithic collapse of virialised high-mass prestellar
cores that have ceased accreting material from their surroundings, and are
supported against fragmentation into low-mass cores by magnetic and turbulent
pressures. In this picture, high-mass prestellar cores (e.g. Motte et al.,
2007) – starless and self-gravitating structures thought to form in infrared
dark clouds (IRDCs; Rathborne et al., 2006; Peretto & Fuller, 2009) – are the
earliest stage of massive star formation. Observationally, however, very few
candidates of truly quiescent high-mass prestellar cores exist (e.g.
Cyganowski et al., 2014; Cyganowski et al., 2022; Duarte-Cabral et al., 2014;
Wang et al., 2014; Kong et al., 2017; Nony et al., 2018; Barnes et al., 2023),
suggesting they may either be very short lived (e.g. Motte et al., 2007;
Kauffmann et al., 2013; Sanhueza et al., 2019), or not exist at all (e.g.
Motte et al., 2018).

High-mass protostars (or massive young stellar objects, MYSOs, e.g. Hoare et
al., 2005; Urquhart et al., 2008) are instead widely associated with active
star formation signatures such as 6.7 GHz CH3OH, 22 GHz H2O, and NH3(3,3)
masers (e.g. Pillai et al., 2006; Urquhart et al., 2011; Brogan et al., 2019;
Jones et al., 2020; Towner et al., 2021). Signatures of ongoing accretion are
also prevalent towards MYSOs, such as high-velocity bi-polar outflows (e.g.
Beuther et al., 2002; Duarte-Cabral et al., 2013; Yang et al., 2022) and less
commonly circumstellar accretion discs (e.g. Beltrán et al., 2014; Johnston et
al., 2015; Ilee et al., 2016; Cesaroni et al., 2017; Maud et al., 2018;
Dewangan et al., 2022). MYSOs are typically observed to have strong
(sub)millimetre continuum, and a sub-population of MYSO(s) – classed as
Extended Green Objects (EGOs; Cyganowski et al., 2008; Cyganowski et al.,
2009) – exhibit extended $4.5\mu$m emission attributed to shocks in outflows
driven by the central MYSO(s). Many of these EGOs only exhibit weak centimetre
continuum emission (Cyganowski et al., 2011b; Towner et al., 2021). The clump-
fed theory of massive star formation (Bonnell et al., 2001; Smith et al.,
2009) describes the competitive accretion of clusters of initially low-mass
star progenitors. High-mass star progenitors develop through the continued
accretion of material by protostellar sources that find themselves at the
centre of the gravitational potential of the cluster. This requirement for a
continually accreting protocluster in the clump-fed theory is a key factor
that distinguishes it from the monolithic core-fed theory, and is a crucial
observable for constraining models of early massive star formation (e.g.
Cyganowski et al., 2017; Issac et al., 2020; Law et al., 2022).

Complex organic molecules (COMs) are molecular species that contain at least
one carbon atom and a total of at least 6 atoms (Herbst & van Dishoeck, 2009).
They are thought to form on dust grain surfaces during both the cold collapse
and "warm-up" phases of massive star formation (e.g. Garrod & Herbst, 2006;
Garrod, 2013; Öberg, 2016; Garrod et al., 2022), with COMs forming earlier and
at lower temperatures in models including non-diffusive chemistry (Garrod et
al., 2022). In general, in gas-grain astrochemical models radiative heating
from MYSOs ultimately heats the grains enough to sublimate their ice mantles,
releasing both COMs and simpler molecules into the gas phase where they are
observable via (sub)millimetre wavelength spectral line emission (e.g. Öberg,
2016; Garrod et al., 2022, and references therein). While COMs can also be
produced by gas-phase mechanisms, production on grains dominates in recent
models (e.g. Garrod et al., 2022). MYSOs that are characterised by a forest of
molecular lines in the (sub)millimetre are classed as hot cores (e.g. Sánchez-
Monge et al., 2017; Sewiło et al., 2018; Liu et al., 2021), and typically have
temperatures that exceed 100 K (e.g. Öberg, 2016). Later in the evolution of
the system (though not necessarily independent of the hot core stage),
hypercompact (HC) Hii regions form, signposted by centimetre continuum due to
the ionization of the surrounding material by the high-mass protostar(s) (e.g.
Kurtz, 2005; Yang et al., 2019; Yang et al., 2021). With the exquisite
sensitivities achievable by ALMA in the (sub)millimetre and the Karl G. Jansky
Very Large Array (VLA) in the centimetre, combined studies of both the
chemical and physical properties of massive star progenitors and their
environments are now possible.

Figure 1: _Spitzer_ GLIMPSE three-colour image (RGB: 8.0, 4.5, 3.6$\mu$m). (a)
The dotted orange contour marks the 30 per cent response level of the ALMA
mosaic, designed to encompass the full extent of the ATLASGAL clump. ATLASGAL
870$\mu$m emission contours are shown in grey (at 12, 16, 20 and 24$\sigma$,
where $\sigma=0.08$ Jy beam-1; Schuller et al., 2009, 18$\arcsec$ resolution).
The zoomed inset of (b) is shown by the dashed black box. (b) High-velocity
blue- and red-shifted SMA 12CO(2–1) emission contours are shown in blue (7.2,
12.0, 15.6, 19.2, 22.8 Jy beam-1 km s-1) and red (4.8, 7.2, 9.6 Jy beam-1 km
s-1) respectively (Cyganowski et al., 2011a), and ALMA 1.05 mm continuum
contours are shown in black (0.00125, 0.004, 0.016, 0.200 Jy beam-1; Williams
et al., 2022). Magenta $+$’s mark the VLA positions of 44 GHz Class I CH3OH
masers (Cyganowski et al., 2009), and the cyan $\times$ marks the intensity-
weighted VLA 6.7 GHz Class II CH3OH maser position from Williams et al.
(2022). Herschel PACS $70\mu$m emission (Poglitsch et al., 2010) of the HI-GAL
clump (Elia et al., 2017) is plotted with white contours (2.5, 4.5, 6.5, 8.5,
10.5 Jy pixel-1). The PACS $70\mu$m beam is $6\arcsec\times 12\arcsec$ (FWHM
$\sim$ 8$\arcsec$; Herschel Explanatory Supplement Volume III, 2017). The ALMA
beam is plotted in the bottom left.

###  1.1 Target: EGO G19.01–0.03

The EGO G19.01–0.03 (hereafter G19.01) is an intriguing example of early
massive star formation. As in Williams et al. (2022), we adopt D=$4.0\pm 0.3$
kpc, the near kinematic distance estimated using the Galactic rotation curve
parameters of Reid et al. (2014) and the NH3 LSRK velocity from Cyganowski et
al. (2013). As shown in Figure 1, emission from the surrounding clump is
detected in the Hi-GAL and ATLASGAL surveys; the clump is also detected at
$1.1$ mm by the Bolocam Galactic Plane Survey (BGPS; Rosolowsky et al., 2010).
The clump mass is $\sim$1000 M⊙: the Hi-GAL catalogue of Elia et al. (2017)
reports a clump mass, radius and temperature of 974 M⊙, 0.16 pc and 17.9 K
respectively for D=3.658 kpc ($\sim$1165 M⊙ and 0.18 pc scaled to D=4.0 kpc),
while Schuller et al. (2009) report a clump mass of 1070 M⊙ for D=4.3 kpc
based on the $870\mu$m ATLASGAL data ($\sim$926 M⊙ scaled to D=4.0 kpc).
Towards the clump, a single millimetre continuum core (hereafter called MM1)
is seen in isolation with the Submillimeter Array (SMA) at 1.3 mm and with the
Combined Array for Research in Millimeter-wave Astronomy (CARMA) at 3.4 mm (at
2.4 and 5.4$\arcsec$ angular resolution respectively; Cyganowski et al.,
2011a). MM1 coincides with 6.7 GHz Class II CH3OH maser emission (Cyganowski
et al., 2009) placing it as a massive source (e.g. Urquhart et al., 2013;
Billington et al., 2020; Jones et al., 2020), and is seen to drive a highly
collimated, high-velocity bi-polar outflow observed with the SMA in 12CO(2–1)
and with CARMA in HCO+(1–0) and SiO(2–1) emission (Cyganowski et al., 2011a).
44 GHz Class I CH3OH maser emission is seen to trace the edges of the outflow
lobes (Cyganowski et al., 2009). At the sensitivity of the SMA and CARMA data,
MM1 exhibited some hot core emission but lacked chemical richness,
conspicuously lacking in emission from Oxygen-bearing COMs. Only two COMs were
detected (CH3OH and CH3CN) with only two lines with excitation temperature
$>$100 K (Cyganowski et al., 2011a). MM1 was not detected in deep, arcsecond-
resolution VLA observations at 3.6 and 1.3 cm to $4\sigma$ limits of 0.12 and
1.04 mJy beam-1 respectively (Cyganowski et al., 2011b), suggesting a very low
ionising luminosity capable of producing only a very small Hii region. Put
together, MM1 appeared as an isolated, high-mass millimetre continuum source,
in a state of ongoing accretion, without strong centimetre continuum or rich
hot-core line emission. As such, MM1 was until recently considered an
excellent candidate for a very early stage of evolution, with potential to
shed light on the core-fed theory of high-mass star formation.

The first paper of our ALMA Cycle 2 follow-up study (Williams et al., 2022,
hereafter Paper I) presented the highest angular resolution observations of
G19.01 to date, at $\sim 0.4\arcsec$ angular resolution in Band 7 at 1.05 mm.
With ALMA, MM1 was observed to exhibit a rich millimetre spectrum with a
variety of COMs, in contrast to the earlier lower resolution and sensitivity
SMA observations. Kinematic analysis of the strongest, most isolated ALMA-
detected molecules revealed the first direct evidence of a rotationally
supported accretion disc around MM1 traced by a velocity gradient
perpendicular to the bi-polar outflow direction, with an enclosed mass of
$40-70$ M⊙ within a 2000 AU radius. In conjunction with new VLA observations
at 5.01 and 1.21 cm, the centimetre-millimetre spectral energy distribution
(SED) was best described by a two-component model, with millimetre emission
dominated by thermal dust, and the $\sim 5$ cm continuum dominated by free-
free emission interpreted as a hypercompact Hii region, placing MM1 in a later
stage of evolution than that concluded with previous observations.
Furthermore, the ALMA 1.05 mm continuum revealed for the first time the
detection of four neighbouring millimetre sources in the vicinity of MM1,
hinting at the possibility of the early stages of protocluster formation.

In this paper (Paper II), we use our ALMA Cycle 2 1.05 mm and VLA 1.21 cm
observations to study the chemistry and protocluster environment of
G19.01–0.03 MM1 and the properties of the newly-detected millimetre sources.
In Section 2 we describe the observations, in Section 3 we present the ALMA
continuum and molecular line emission, as well as VLA ammonia and methanol
emission. In Section 4 we present our modelling of the COM emission, rotation
diagram analysis, continuum properties of the millimetre neighbours, and
discuss MM1’s chemistry in the context of sources from the literature. We
summarise our main conclusions in Section 5.

##  2 Observations

###  2.1 Atacama Large Millimetre/submillimetre Array (ALMA)

Our ALMA Cycle 2 observations were designed to search for low-mass cores
within the clump-scale gas reservoir associated with G19.01–0.03 MM1: the
extent of the ALMA mosaic ($\sim$40$\arcsec\approx$0.78 pc at D=4 kpc) is
shown in Figure 1 and observing parameters are summarised in Table 1 and
below. These observations are also described in detail in Paper I.

For our observations, the ALMA correlator was configured to cover seven
spectral windows (spws), including five narrow spws targeting particular
spectral lines and two wide spws. Details of the narrow spws are given in
Table 2. The wide spws, with central frequencies of $\sim$278.2 GHz and
$\sim$292.0 GHz, each have a bandwidth of 1.875 GHz, a Hanning-smoothed
spectral resolution of 1.13 MHz (1.156$\times$ the channel spacing of 0.977
MHz because of online channel averaging in the ALMA correlator).

Table 1: Observing parameters for ALMA 1.05 mm and VLA 1.21 cm data.

Parameter | ALMA 1.05 mm | VLA 1.21 cm
---|---|---
Observing date | 14 May 2015 | 11-12 Nov 2013
Project code (PI) | 2013.1.00812.S | 13B-359
| (C. Cyganowski) | (T. Hunter)
Gain calibrator | J1733–1304 | J1832–1035
Bandpass calibrator | J1733–1304 | J1924–2914
Flux calibrator |  Titana | J1331+3030
On-source integration time | 44 min | 139 min
Number of antennas | 37 | 25
Antenna configuration | C43-3/(4) | B
Phase Centre (J2000): |  |
R.A. (${}^{\mathrm{h\,m\,s}}$) |  18:25:44.61b | 18:25:44.80
Dec. (∘ ′ ′′) |  -12:22:44.00b | -12:22:46.00
Projected baseline lengths | 20–533 m | 0.14–9.98 km
|  19–508 $k\lambda$ |  12–825 $k\lambda$
Mean frequencyc | 285.12 GHz | 24.81 GHz
Mean wavelengthc | 1.05 mm | 1.21 cm
Number of pointings | 7 | 1
Field of viewc,d |  $\sim$40$\arcsec$ |  1.8′
Synthesised beamc | $0\aas@@fstack{\prime\prime}52\times 0\aas@@fstack{\prime\prime}35$ | $0\aas@@fstack{\prime\prime}33\times 0\aas@@fstack{\prime\prime}22$
Beam position anglec,e |  88.4∘ |  0.5∘
Maximum Recoverable Scalef | $4\aas@@fstack{\prime\prime}2$ | $4\aas@@fstack{\prime\prime}5$
Continuum rms noiseg |  0.25 mJy beam-1 |  6.0 $\mu$Jy beam-1

a Using Butler-JPL-Horizons 2012 models.

b For the central pointing of the mosaic.

c For the continuum image.

d ALMA: to 30% level of mosaic response. VLA: primary beam FWHP at mean
frequency.

e Measured East of North i.e. positive in the anti-clockwise direction.

f Calculated from the fifth percentile shortest baseline (as stated in the
ALMA Technical Handbook) and mean frequency, using au.estimateMRS from the
analysisUtils Python package.

g Estimated from emission-free regions within the 30% response level of the
ALMA mosaic.

As detailed in Paper I, the data were calibrated using the casa 4.2.2 version
of the ALMA calibration pipeline and line-free channels were identified using
the approach of Brogan et al. (2016); Cyganowski et al. (2017). These line-
free channels were used to construct a pseudo-continuum dataset and to perform
continuum subtraction in the _u,v_ -plane. For the narrow spw targeting C33S,
this process was problematic due to wide lines and possible absorption (see
also Paper I and Cyganowski et al., 2017) and we excluded this spw – which
overlaps one of the wide spws in our tuning – from the pseudo-continuum
dataset and our line analysis. The aggregate continuum bandwidth of the final
pseudo-continuum dataset is $\sim$1.6 GHz. The continuum data were iteratively
self-calibrated and were imaged using Briggs weighting with a robust parameter
of 0 and multi-frequency synthesis; the key parameters of the resulting image
are listed in Table 1. The synthesised beamsize of the continuum image
(0$\aas@@fstack{\prime\prime}$52$\times$0$\aas@@fstack{\prime\prime}$35)
corresponds to a physical scale of 2080$\times$1400 AU at 4 kpc. As our
observations included only the ALMA 12 m array, the maximum recoverable scale
is $4.2\arcsec{\sim}16,800$ au (at 4 kpc).

Table 2: Observing and image cube parameters for narrow spectral windows

Telescope | Targeted line |  Line Rest Frequencya |  E${}_{u}/k_{B}$a | Bandwidth |  $\Delta\nu$b |  $\Delta$vc | Synthesised beam |  rms noised
---|---|---|---|---|---|---|---|---
|  | (GHz) | (K) | (MHz) | (kHz) |  (km s-1) |  $\arcsec\times\arcsec$ [∘] |  (mJy beam-1) | (K)
ALMA |  N2H+(3–2) | 279.5117491 | 26.8 | 468.75 | 122.07 | 0.5 |  0.57 $\times$ 0.41 [83.2] | 6.1 | 0.41
ALMA | DCN (4–3) | 289.644917 | 34.8 | 117.1875 | 122.07 | 0.5 |  0.55 $\times$ 0.40 [82.0] | 6.3 | 0.42
ALMA |  34SO 67–56 | 290.562238 | 63.8 | 117.1875 | 122.07 | 0.5 |  0.55 $\times$ 0.40 [82.2] | 6.0 | 0.39
ALMA |  H2CO $4_{0,4}-3_{0,3}$ | 290.623405 | 34.9 | 117.1875 | 122.07 | 0.5 |  0.55 $\times$ 0.40 [82.3] | 6.0 | 0.39
ALMA |  C33S (6–5) | 291.485935 | 49.0 | 117.1875 | 122.07 | 0.5 |  0.55 $\times$ 0.40 [82.4] | … | …
VLA |  NH3 (1,1) | 23.694496 | 24.4 | 8.0 | 15.625 | 0.4 |  0.56 $\times$ 0.54 [28.0] | 1.24 | 8.93
VLA |  NH3 (2,2) | 23.722631 | 65.6 | 8.0 | 15.625 | 0.4 |  0.56 $\times$ 0.54 [28.5] | 1.18 | 8.47
VLA |  NH3 (3,3) | 23.870130 | 124.7 | 8.0 | 15.625 | 0.4 |  0.56 $\times$ 0.54 [30.5] | 1.13 | 8.01
VLA |  NH3 (5,5) | 24.532985 | 296.5 | 8.0 | 15.625 | 0.4 |  0.55 $\times$ 0.54 [38.4] | 1.10 | 7.52
VLA |  NH3 (6,6) | 25.056025 | 409.2 | 4.0 | 15.625 | 0.4 |  0.55 $\times$ 0.54 [41.1] | 0.95 | 6.23
VLA |  NH3 (7,7) | 25.715182 | 539.7 | 4.0 | 15.625 | 0.4 |  0.54 $\times$ 0.53 [47.8] | 1.06 | 6.84
VLAe |  CH3OH 3(2,1)-3(1,2) | 24.928707 | 36.2 | 2.0 | 15.625 | 0.4 |  0.56 $\times$ 0.55 [31.5] | 1.04 | 6.76
VLAe |  CH3OH 5(2,3)-5(1,4) | 24.9590789 | 57.1 | 2.0 | 15.625 | 0.4 |  0.55 $\times$ 0.54 [49.7] | 1.01 | 6.67
VLAe |  CH3OH 8(2,6)-8(1,7) | 25.2944165 | 105.8 | 2.0 | 15.625 | 0.4 |  0.57 $\times$ 0.55 [52.8] | 0.85 | 5.18
VLAe |  CH3OH 10(2,8)-10(1,9) | 25.8782661 | 150.0 | 2.0 | 15.625 | 0.4 |  0.56 $\times$ 0.54 [65.0] | 1.19 | 7.18

a From CDMS (Müller et al., 2001, 2005) for lines observed with ALMA and from
the TopModel line list for NH3 lines, both accessed via the NRAO spectral line
catalogue (Splatalogue; https://splatalogue.online/). CH3OH line frequencies
are from Müller et al. (2004).

b Channel spacing. For the ALMA data, the Hanning-smoothed spectral resolution
is 0.244 MHz for all of the narrow spws.

c Velocity channel width of image cubes; see §2.1 and §2.2.

d Median rms noise estimated from emission-free channels; the rms noise is up
to $\sim$1.5 times higher in channels with bright and/or complex emission. The
conversion to brightness temperature assumes the Rayleigh-Jeans approximation:
$T=1.222\times
10^{3}\frac{I}{\nu^{2}\theta_{\mathrm{maj}}\theta_{\mathrm{min}}}$, where $I$
is the rms in mJy beam-1, $\nu$ is the frequency in GHz, and
$\theta_{\mathrm{maj}}\times\theta_{\mathrm{min}}$ is the synthesised beam
size in arcseconds.

e As in Towner et al. (2017), these 25 GHz CH3OH lines will be referred to in
the main text by the following shorthand notation, given by the first two
values of the upper state quantum number of each transition respectively: 32,
52, 82, 102.

After applying the solutions from the continuum self-calibration, the
continuum-subtracted line data were imaged with Briggs weighting with a robust
parameter of 0.5. For the narrow spws, which were imaged with 0.5 km s-1
channels for better sensitivity to faint emission, the synthesised beamsizes
and rms noise levels of the image cubes are listed in Table 2. For the wide
spws, which were imaged with the native channel spacing, the synthesised
beamsizes are
0$\aas@@fstack{\prime\prime}$60$\times$0$\aas@@fstack{\prime\prime}$43 [P.A.
79.8∘] and
0$\aas@@fstack{\prime\prime}$57$\times$0$\aas@@fstack{\prime\prime}$41 [P.A.
78.9∘] for the spws centred at $\sim$278.2 GHz and $\sim$292.0 GHz,
respectively, and the rms noise is $\sim$3 mJy beam-1 in emission-free
channels (in channels with complex emission, the rms is up to $\sim$1.5 times
higher; see also Paper I). All measurements were made from images corrected
for the response of the primary beam.

###  2.2 Karl G. Jansky Very Large Array (VLA)

In this paper, we present our K-band VLA spectral line observations of
G19.01–0.03. Our VLA tuning included 10 narrow spws targeting NH3 and CH3OH
lines observed in other EGOs, including lines that exhibit maser activity in
some EGOs (e.g. Brogan et al., 2011; Towner et al., 2017, see also §3.2.3;
§3.2.4). The VLA observing parameters are summarised in Table 1 and details of
the narrow spws are given in Table 2. The VLA K-band tuning also included
16$\times$0.128 GHz spws for continuum, as detailed in Paper I which presented
the continuum results; for completeness, key parameters of the VLA 1.21 cm
continuum image are included in Table 1.

As explained in Paper I, the VLA data were calibrated using the casa 4.7.1
version of the VLA calibration pipeline. The data were Hanning smoothed, and
phase-only self-calibration was performed using the channel with the strongest
NH3(3,3) emission (after continuum subtraction in the _u,v_ -plane). These
solutions were then applied to all of the line data (as well as to the
continuum; see Paper I). For each narrowband spw, continuum subtraction was
performed in the _u,v_ -plane and the continuum-subtracted line data were
imaged with 0.4 km s-1 channels, Briggs weighting with a robust parameter of
0.5, and a _uv_ taper of 200 k$\lambda$ to improve the brightness temperature
sensitivity. The synthesised beamsizes and rms noise levels of the resulting
image cubes are presented in Table 2. Measurements were made from images
corrected for the primary beam response.

##  3 Results

###  3.1 ALMA 1.05 mm continuum emission

In Paper I, we presented the ALMA 1.05 mm continuum towards MM1, and noted the
detection of a further four continuum sources in the field for the first time
(see Figure 2), named MM2…MM5 in order of decreasing peak intensity. Their
observed properties (as well as the observed properties of MM1 as presented in
Paper I) are listed in Table 3, with their FWHM extents represented by pink
ellipses in Figure 2b. As detailed in Paper I, sources were extracted using
the astrodendro algorithm (Rosolowsky et al., 2008), with a minimum isocontour
value ($I_{\mathrm{min}}$) of $5\sigma_{\mathrm{rms}}$ (where
$\sigma_{\mathrm{rms}}=0.25$ mJy beam-1), minimum isocontour spacing ($\Delta
I_{\mathrm{min}}$) of $1\sigma_{\mathrm{rms}}$, and minimum size of a
structure ($n_{\mathrm{pix}}$) of $\approx n_{\mathrm{pix,beam}}/2$ (i.e. half
the beam size, where $n_{\mathrm{pix,beam}}\approx 50$). Two sources to the
south-west of MM4 are detected with peak emission $>5\sigma_{\mathrm{rms}}$
(Figure 2) but are only extracted by the dendrogram algorithm if the
parameters are dropped to $\Delta I_{\mathrm{min}}=0.9\sigma_{\mathrm{rms}}$
and $n_{\mathrm{pix}}=15$ pixels, meaning that they are only equivalent to a
third of a beam in size. We therefore do not consider these firm detections.
MM2…MM5 have angular separations from MM1 of 1.6, 5.5, 6.1 and 2.6$\arcsec$
respectively, ranging between $0.03-0.12$ pc at the 4 kpc distance, marking
the first detection of other millimetre sources within the parent clump of
MM1. MM4 appears non-gaussian in its emission morphology, whilst MM2 lies
within a common contour to MM1 with lower surface-brightness emission
connecting the two sources. This suggests that MM2 may be fragmenting out of
material that also feeds MM1.

Figure 2: (a) ALMA 1.05 mm continuum image, corrected for the primary beam
response, as shown in Paper I. The field shown is a sub-region of that
mosaiced, which contains all emission detected $\geq 5\sigma$. Black contours
are plotted at [5, 8, 16, 32, 64, 200, 400 and 800]$\times\sigma$, where
$\sigma=0.25$ mJy beam-1. The synthesised beam and scalebar are plotted in the
bottom-left and bottom-right respectively. (b) Same as (a), but with the
outlines of the extracted dendrogram structures in dashed black, and ellipses
representing the source sizes (Table 3) in pink. The peak positions of the
sources (Table 3) are marked by the black $\times$s. Table 3: Observed
properties of extracted 1.05 mm continuum sources.

Source | J2000.0 Coordinatesa | Peak | Integ. | Source size c | Source sizec
---|---|---|---|---|---
| $\alpha$ | $\delta$ | intensity b | flux b | Maj. $\times$ Min. [P.A.] |
| (${}^{\mathrm{h\,m\,s}}$) | (∘ ′ ′′) | (mJy beam-1) | (mJy) | (arcsec $\times$ arcsec [∘]) | (au)
MM1 | 18:25:44.782 | -12:22:45.92 | $266.3$ | $303.1$ | $1.15\times 0.84~{}[78.7]$ | $4600\times 3360$
MM2 | 18:25:44.888 | -12:22:45.68 | $4.7$ | $8.1$ | $0.83\times 0.74~{}[-82.1]$ | $3310\times 2970$
MM3 | 18:25:44.446 | -12:22:43.52 | $3.0$ | $7.1$ | $1.14\times 0.88~{}[-53.0]$ | $4560\times 3530$
MM4 | 18:25:44.880 | -12:22:51.86 | $2.3$ | $6.0$ | $1.56\times 0.65~{}[-47.7]$ | $6200\times 2610$
MM5 | 18:25:44.622 | -12:22:47.06 | $1.6$ | $-$ | $-$ | $-$

  * $a$

Peak position. The number of significant figures reflects a one pixel
uncertainty.

  * $b$

Evaluated within the intensity-weighted second moment size (not the total
dendrogram structure). For sources smaller than a beam, integrated fluxes are
marked “$-$”.

  * $c$

Deconvolved major and minor axes sizes; position angle is measured East of
North i.e. positive in the anti-clockwise direction. Sizes are the intensity-
weighted second moment, converted to FWHM (see Paper i). Sources smaller than
a beam are marked “$-$”.

###  3.2 Line emission

####  3.2.1 COMs towards MM1 with ALMA

Figure 3: ALMA spectra towards the 1.05 mm continuum peak of MM1 in the broad
spectral windows centred at 278.209 GHz and 292.021 GHz. Unlabelled lines
could not be confidently identified, mostly due to blending at the coarse
spectral resolution of our observations. Lines labelled in black were used for
the kinematic analysis presented in Paper I, while others are labelled in
grey. All marked lines are listed in Table 4. Table 4: Properties of spectral
lines identified in the broad ALMA spectral windows towards MM1, arranged by
decreasing $E_{u}$/$k_{B}$.

Speciesa | Transition | Frequency |  $E_{u}$/$k_{B}$ | $S_{ij}\mu^{2}$ | $g_{u}$ |  Catalogueb |  Paper I?c
---|---|---|---|---|---|---|---
|  | (GHz) | (K) |  (D2) |  |  |
CH3OH (vt = 0) |  234,19–225,18 | 278.96513 | 736.0 | 7.05321 | 47 | JPL | Y
CH3OH (vt = 0) |  21-2,20–20-3,18 | 278.47989 | 563.2 | 6.45393 | 43 | JPL | Y
HC3N (v7=1) | J = 32–31, I = 1f | 292.19837 | 551.4 | 443.01616 | 65 | CDMS | N
HC3N (v7=1) | J = 32–31, I = 1e | 291.78201 | 551.1 | 442.99035 | 65 | CDMS | N
CH3OH (vt = 0) |  176,12–184,13\- - | 291.90814 | 548.6 | 3.98258 | 35 | JPL | Y
CH3OH (vt = 0) |  185,13–194,16\- - | 278.72314 | 534.6 | 5.26251 | 37 | JPL | Y
CH3OH (vt = 0) |  185,14–194,15\+ + | 278.67303 | 534.6 | 5.26201 | 37 | JPL | N
CH3OH (vt = 1) |  101,10–90,9 | 292.51744 | 418.8 | 5.01720 | 21 | JPL | Y
CH3OCHO (v = 1) |  234,19–224,18A | 292.31711 | 365.4 | 59.03484 | 94 | JPL | N
t-CH3CH2OH |  276,22–275,23 | 278.84892 | 363.7 | 28.22445 | 55 | JPL | N
CH3OH (vt = 0) |  144,10–153,12 | 278.59908 | 339.6 | 4.14979 | 29 | JPL | Y
CH3OH (vt = 0) |  151,0–142,0 | 291.24057 | 295.3 | 5.70014 | 31 | JPL | N
t-CH3CH2OH |  236,17–235,18 | 278.68234 | 277.5 | 23.74792 | 47 | JPL | N
CH3CH2CN (v = 0) |  317,24–307,23 | 278.00758 | 267.8 | 435.64965 | 63 | JPL | Y
CH3CH2CN (v = 0) |  316,25–306,24 | 278.26670 | 253.4 | 441.91785 | 63 | JPL | Y
CH3CH2CN (v = 0) |  316,26–306,25 | 278.25123 | 253.4 | 441.86236 | 63 | JPL | N
CH3CH2CN (v = 0) |  315,26–305,25 | 278.86581 | 241.4 | 447.13304 | 63 | JPL | N
g-CH3CH2OH |  166,10–156,9 (vt=0-0) | 277.41431 | 213.8 | 21.96640 | 33 | JPL | Y
g-CH3CH2OH |  164,12–154,11 (vt=0-0) | 278.64299 | 189.7 | 23.96604 | 33 | JPL | Y
CH3OCHO (v = 0) |  242,22–232,21E | 279.05752 | 178.5 | 62.00988 | 98 | JPL | N
CH3OCHO (v = 0) |  242,22–232,21A | 279.06596 | 178.5 | 62.01870 | 98 | JPL | N
OCS (v = 0) | J = 24–23 | 291.83965 | 175.1 | 12.27714 | 49 | CDMS | N
CH3OCHO (v = 0) |  226,16–216,15E | 279.05067 | 175.0 | 54.25478 | 90 | JPL | N
CH3OCHO (v = 0) |  226,16–216,15A | 279.07471 | 175.0 | 54.26998 | 90 | JPL | N
CH3OCHO (v = 0) |  234,20–224,19A | 277.74543 | 173.4 | 58.74803 | 94 | JPL | N
H2CO |  43,2–33,1 | 291.38049 | 140.9 | 28.54373 | 27 | CDMS | N
t-CH3CH2OH |  171,16–162,15 | 277.38114 | 131.5 | 7.54333 | 35 | JPL | N
CH3OCH3 |  161,16–150,15 (EA) | 292.41225 | 120.3 | 70.75903 | 66 | JPL | Y
NH2CHO |  133,10–123,9 | 277.51403 | 119.8 | 482.76031 | 81 | JPL | Y
CH3OH (vt = 0)d |  9-1,9–80,8 | 278.30451 | 110.0 | 5.82054 | 19 | JPL | N
CH3CHO (vt = 0) |  15-1,15–14-1,14 (E) | 278.92443 | 109.8 | 188.77851 | 62 | JPL | N
CH3CHO (vt = 0) |  151,15–141,14 (A) | 278.93944 | 109.7 | 188.63110 | 62 | JPL | N
CH3OCH3 |  132,12–121,11 (EE) | 291.44307 | 88.0 | 116.69078 | 216 | JPL | Y
H2CO d |  42,2–32,1 | 291.94807 | 82.1 | 16.31005 | 9 | CDMS | Y
H2CO |  42,3–32,2 | 291.23777 | 82.1 | 16.30879 | 9 | CDMS | N
CH3OCH3 |  122,11–111,10 (EE) | 278.40706 | 76.3 | 103.86585 | 200 | JPL | Y
H2CS d |  81,7–71,6 | 278.88640 | 73.4 | 64.08219 | 51 | JPL | N
CH3OH (vt = 0) d |  61,5–51,4 \- - | 292.67291 | 63.7 | 4.72131 | 13 | JPL | Y
13CH3OH (vt = 0) |  32,2–41,3 | 291.53662 | 51.4 | 0.69489 | 7 | CDMS | Y
t-CH3CH2OH |  93,6–82,7 | 291.73815 | 49.2 | 7.83704 | 19 | JPL | N
13CS (v = 0) | 6–5 | 277.45540 | 46.6 | 46.00530 | 26 | CDMS | Y
CH3OCH3 |  73,5–62,4 (EE) | 277.64832 | 38.1 | 65.30773 | 120 | JPL | N
CH3OH (vt = 0) |  2-2,1–3-1,3 | 278.34222 | 32.9 | 0.32917 | 5 | JPL | Y

a Bold typeface indicates the lines used for the analysis presented in
Sections 4.1.2–4.1.4.

b CDMS (Müller et al., 2001) and JPL (Pickett et al., 1998), accessed via the
NRAO spectral line catalogue (Splatalogue; https://splatalogue.online/).

c Lines used for the kinematic analysis in Paper I are marked with a “Y”.
Lines that are newly presented here are marked with an “N”.

d Outflow tracing, see §3.2.2.

A forest of molecular lines is observed towards MM1 with ALMA, as seen in the
wideband spectra shown here in Figure 3. We follow the criteria presented by
Herbst & van Dishoeck (2009) in identifying molecular lines (also see Maret et
al., 2011). As outlined in Paper I, we attribute peaks in emission to
cataloged rest frequencies from the JPL (Pickett et al., 1998) and CDMS
(Müller et al., 2001) databases at the source systemic velocity ($59.9\pm 1.1$
km s-1; Cyganowski et al., 2011a). We further produced LTE synthetic spectra
of the molecular emission (using the Weeds extension of class; Maret et al.,
2011) – a molecular transition was positively identified when all lines for
that species predicted in the synthetic spectrum were present in the observed
spectrum, for typical model parameters expected of a hot core (e.g. $T>100$
K). The use of LTE synthetic spectra also allowed the identification of some
emission peaks that were consistent with multiple line rest frequencies within
our spectral resolution ($\sim$1 km s-1 in the wide spws), for example blended
lines and lines with shoulder features. Using this approach, we identify 43
line transitions from 11 different species in the wide ALMA spws at $\sim$278
GHz and $\sim$292 GHz (see Figure 3 and Table 4), including isolated lines,
blended lines, and lines with shoulder features. We note that many lines
remain unidentified despite their strong detection above the noise, generally
due to rest frequencies and/or synthetic line profiles that cannot be
confidently distinguished at our spectral resolution. Of particular note, we
report compact emission towards MM1 (see Figure 3 of Paper I) from a range of
complex organic molecules (COMs), including Oxygen and Nitrogen-bearing COMs
such as CH3OCH3, CH3CHO, CH3OH, NH2CHO and CH3OCHO. We also include in our
analysis of MM1’s line emission (§4.1 and §4.3) additional lines from three of
these species that are serendipitously included in the narrow, targeted ALMA
bands: CH3OCHO ($27_{1,27}-26_{0,26}$) with $\nu_{\mathrm{rest}}=289.62659$
GHz and $E_{u}/k_{B}=385.4$ K (blended with the CH3OH($6_{1,2}-5_{1,2}$) line
with $\nu_{\mathrm{rest}}=289.62430$ GHz and $E_{u}/k_{B}=731.4$ K), CH3OH
($11_{2,10}-10_{3,7}$) with $\nu_{\mathrm{rest}}=279.35193$ GHz and
$E_{u}/k_{B}=190.9$ K, and OCS ($23-22$) with $\nu_{\mathrm{rest}}=279.6853$
GHz and $E_{u}/k_{B}=161.1$ K. The majority of the lines targeted with narrow
spws (§2.1 and Table 2) exhibit extended emission and are discussed in Section
3.2.2. The exception is 34SO 67–56, which exhibits compact emission but which
we do not include in our analyses due to line blending in the narrowband cube.

In contrast, in the SMA data presented by Cyganowski et al. (2011a), MM1 was
relatively line-poor (see §1): the only detected COM emission in their 2 GHz-
wide bands, centred at $\sim$ 220 and $\sim$ 230 GHz, was from CH3OH and
CH3CN. With ALMA, we also detect molecular lines with up to four times higher
$E_{u}/k_{B}$ than detected with the SMA (e.g. CH3OH($23_{4,19}-22_{5,18}$) at
E${}_{u}/k_{B}=736$ K), and identify 19 lines with E${}_{u}/k_{B}>200$ K. The
relative dearth of molecular lines in the SMA spectra is likely attributable
to a combination of sensitivity and beam dilution; our $\sim$
0.4$\arcsec$-resolution ALMA observations improve on the
2.4$\arcsec$-resolution SMA observations by a factor of $\sim$30 in beam area,
and a factor of $\sim$1.8 in brightness temperature sensitivity. Comparing our
ALMA detections with those from the sensitive, 29$\arcsec$-resolution 1 mm
single-dish survey of He et al. (2012) confirms the importance of beam
dilution: their observations of G19.01$-$0.03 have a 1$\sigma$ rms of
$0.015-0.022$ K (compared to 0.25 K and 0.44 K for the ALMA and SMA
observations, respectively), but the only COM detected is CH3OH. In sum, our
ALMA results affirm the hot core classification of G19.01–0.03 MM1 and
illustrate the importance of sensitive, high-resolution observations for
studying the chemistry of MYSOs.

####  3.2.2 Extended emission with ALMA

Figure 4: ALMA peak intensity maps of (a) N2H+($3-2$), (b)
H2CO($4_{0,4}-3_{0,3}$), (c) H2CO($4_{2,3}-3_{2,1}$), (d) CH3OH
$\mathrm{v}_{t}=0$ ($6_{1,5}-5_{1,4}$), (e) CH3OH $\mathrm{v}_{t}=0$
($9_{-1,9}-8_{0,8}$), and (f) DCN v=0 ($4-3$), in units of Jy beam-1 shown on
a logarithmic scale (images shown are not primary beam corrected). Millimetre
sources are labelled in (f), where two zoomed insets are also shown centred on
MM1/MM2 and MM4. A zoomed inset of MM1/MM2 is also shown in (a). Overplotted
magenta contours in all panels show the 1.05 mm continuum emission at
5$\sigma$. These contours are also shown in the zoomed insets, but at 5, 8,
16, 64 and $600\sigma$ (where $\sigma=0.25$ mJy beam-1). In panel (d), white
contours show the VLA NH3(3,3) peak intensity map (in $10\sigma$ steps from
$5-145\sigma$, where $\sigma$ = 1.1 mJy beam-1, non-primary beam corrected),
black contours show the VLA NH3(6,6) peak intensity map (at the $5\sigma$
level, where $\sigma=0.9$ mJy beam-1, non-primary beam corrected), and the
orange cross indicates the peak position of the candidate VLA 25 GHz CH3OH
5(2,3)-5(1,4) maser emission (see §3.2.4). The NH3(3,3) maser groups from
§3.2.3 are numbered in panel (d) as in Table 5. Overplotted white crosses in
(e) mark the positions of 44 GHz Class I CH3OH masers from Cyganowski et al.
(2009). Each panel shows the ALMA synthesised beam in the bottom left (as well
as the VLA synthesised beam in panel (d)), a $1^{\prime\prime}$ scale bar in
the bottom right, and the molecule name, frequency and upper energy in the top
right.

Figure 4 presents peak intensity maps of a selection of lines that are
representative of the extended emission in our ALMA tuning. These include
three of the lines targeted in our narrow spectral windows (N2H+(3–2), DCN v=0
($4-3$) and H2CO($4_{0,4}-3_{0,3}$), with E${}_{u}/k_{B}=26.8$ K, 34.8 K and
34.9 K respectively), and three from our wide spectral windows
(H2CO($4_{2,3}-3_{2,1}$), CH3OH($6_{1,5}-5_{1,4}$) and
CH3OH($9_{-1,9}-8_{0,8}$), with E${}_{u}/k_{B}=82.1$ K, 63.7 K and 110.0 K
respectively). The H2CO($4_{2,3}-3_{2,1}$) and CH3OH($6_{1,5}-5_{1,4}$) lines
in Figure 4(c) and 4(d) were identified as having extended emission around MM1
in Paper I.

The four H2CO and CH3OH lines shown in Figure 4 appear to spatially trace the
same bi-polar outflow structure identified by Cyganowski et al. (2011a) in
12CO(2–1) with the SMA (see Figure 1) and in HCO+(1–0) and SiO(2–1) with
CARMA. Kinematically, however, the ALMA H2CO and CH3OH emission traces lower
velocity gas than the SMA 12CO or the CARMA HCO+ emission, with a median full
velocity extent of 25 km s-1 compared to $\sim$ 135 km s-1 for 12CO(2–1) and
$\sim$ 76 km s-1 for HCO+(1–0) (Cyganowski et al., 2011a). The H2CO and CH3OH
kinematics seen with ALMA do however exhibit a similar asymmetry as observed
with the SMA and CARMA, with the blue-shifted lobe extending to higher
velocities (up to $\sim$20 km s-1 from the systemic velocity, Vsys) than the
red-shifted lobe (up to $\sim$9 km s-1 from Vsys). Overplotted in Figure 4(e)
are the positions of 44 GHz Class I CH3OH masers from Cyganowski et al.
(2009), which appear to trace the outer edge of the CH3OH and H2CO outflow
lobes. Taken together, our results are consistent with the ALMA CH3OH and H2CO
emission tracing lower velocity, outflow–cloud interaction regions or outflow
cavity walls, as also seen in the EGO G11.92–0.61 by Cyganowski et al. (2017).

The CH3OH($9_{-1,9}-8_{0,8}$) line in Figure 4(e) (with
$\nu_{\mathrm{rest}}=278.30451$ GHz) is known to exhibit Class I maser
emission towards other sources in the literature (e.g. Voronkov et al., 2012;
Yanagida et al., 2014; Cyganowski et al., 2017). Comparing the emission of
this line with that of CH3OH($6_{1,5}-5_{1,4}$)
($\nu_{\mathrm{rest}}=292.67291$ GHz), shown in Figure 4(d), the two lines
have similar emission morphologies. As shown in Figure 5, these lines also
have similar fluxes towards the ALMA 1.05 mm continuum peak of MM1 (equivalent
to brightness temperatures of T${}_{b}=39$ and 33 K respectively). However,
towards the region of brightest CH3OH($9_{-1,9}-8_{0,8}$) emission in the
outflow (T${}_{b}=49$ K at 18h25m44$\aas@@fstack{s}$517
$-$12∘22′32$\aas@@fstack{\prime\prime}$652 (J2000)), the
CH3OH($6_{1,5}-5_{1,4}$) line is an order of magnitude weaker (Figure 5) with
T${}_{b}=4$ K. The CH3OH($9_{-1,9}-8_{0,8}$) line is also notably narrower at
the outflow position than towards the MM1 continuum peak ($\Delta$V${}_{\rm
FWHM}=2.01\pm 0.01$ km s-1 and $6.95\pm 0.08$ km s-1, respectively), and is
spatially and kinematically coincident with 44 GHz Class I CH3OH masers
(Cyganowski et al., 2009), candidate 229.759 GHz CH3OH maser emission
(Cyganowski et al., 2011a, see their Fig.12), and NH3(3,3) masers (see
§3.2.3). The coincidence with probable 229.759 GHz maser emission is notable
because the 229.759 and 278.305 GHz transitions are in the same Class I maser
series (the 36 GHz maser series; Voronkov et al., 2012), and Cyganowski et al.
(2018) found that all 229 GHz masers in the EGO G11.92–0.61 have probable 278
GHz maser counterparts. Though the angular resolution of our ALMA observations
is insufficient to rely on line brightness temperatures to distinguish between
thermal and maser behaviour (as was also the case for G11.92–0.61 in
Cyganowski et al., 2017), the line properties shown in Figure 5, and the
coincidence of the brightest CH3OH($9_{-1,9}-8_{0,8}$) emission in the outflow
with other shock-excited masers, strongly suggest thermalised emission towards
the MM1 hot core, and 278 GHz Class I CH3OH maser emission tracing
outflow–cloud interaction regions (e.g. Voronkov et al., 2012; Yanagida et
al., 2014; Cyganowski et al., 2017). Higher angular resolution observations of
the CH3OH($9_{-1,9}-8_{0,8}$) line would be required to confirm the presence
of maser emission based on brightness temperature.

Figure 5: Spectra of the CH3OH$(9_{-1,9}-8_{0,8})$ line at $E_{u}/k_{B}=110.0$
K (in blue) and the CH3OH$(6_{1,5}-5_{1,4})$ line at at $E_{u}/k_{B}=63.7$ K
(in orange) towards (a) the ALMA 1.05 mm continuum peak of MM1 (Table 3), and
(b) the position of peak CH3OH$(9_{-1,9}-8_{0,8})$ emission in the northern
outflow lobe (§3.2.2). The systemic velocity of MM1 ($59.9\pm 1.1$ km s-1;
Cyganowski et al., 2011a) is marked by the vertical dashed line.

As a dense gas tracer, N2H+ is known to trace infrared-dark clumps at both
quiescent and active evolutionary stages: Sanhueza et al. (2012), for example,
detect N2H+(1–0) towards the majority of their 92 infrared-dark clumps with
the Mopra telescope at 38′′ angular resolution (typically $\sim$0.8 pc across
their sample). Notably, while we observe extended N2H+(3–2) emission in
G19.01–0.03 (Figure 4a), its morphology does not resemble that of the 1.05 mm
dust emission or that of the outflow structure. Cyganowski et al. (2017) found
similar results in their ALMA N2H+(3–2) observations of the EGO G11.92–0.61
(their Fig. 3c), suggesting that in high-mass clumps the N2H+(3–2) and
millimetre continuum emission do not trace the same structures on small size
scales.111N2H+(3–2) is detected in absorption against the millimetre continuum
of MM1, see Appendix A. We also note that larger-scale N2H+(3–2) emission will
be affected by spatial filtering due to missing short-spacing information
(Table 1; §2.1). The morphology of the DCN(4–3) emission, on the other hand
(Figure 4(f)), more closely resembles the ALMA 1.05 mm continuum. DCN(4–3) is
detected towards MM1 to 38$\sigma$ (where $\sigma=6.3$ mJy beam-1): the
position of the DCN emission peak is offset from the ALMA 1.05 mm continuum
peak by 0.1$\arcsec$ $\sim$ 400 au (within a DCN beam; Table 2), likely due to
line opacity and self-absorption towards the continuum peak. The DCN emission
also traces some of the nearby millimetre companions to MM1, which is
discussed in Section 3.2.5.

####  3.2.3 Ammonia emission with the VLA

Of the ammonia lines in our VLA tuning, NH3(3,3) and NH3(6,6) are known to
exhibit maser behaviour thought to arise due to outflow-induced shocks (e.g.
Mangum & Wootten, 1994; Kraemer & Jackson, 1995; Brogan et al., 2011). Towards
G19.01$-$0.03, spatially compact regions of NH3(3,3) emission are detected at
the $\geq$5$\sigma$ level (shown as white contours in Figure 4(d)) coincident
both spatially and kinematically with the outflow-tracing ALMA H2CO and CH3OH
emission shown in Figure 4(b)-(e). To characterise the NH3(3,3) emission, we
identify as firm detections, and potential maser candidates, locations where
$\geq$5$\sigma$ emission in $\geq$2 consecutive velocity channels is spatially
contiguous within a VLA beam. Following the terminology of Towner et al.
(2021), we refer to emission in a single velocity channel (each 0.4 km s-1
wide; Table 2) as a “spot”, and emission from multiple spatially contiguous
spots as a “group”: a group thus contains at least two spots.

Table 5: Properties of NH3(3,3) emission groups.

Groupa | J2000.0 centroid positionb | Ipeakc | Angular spread d | Vmin, Vmax e | Vpeak f
---|---|---|---|---|---
ID |  $\alpha$ (${}^{\mathrm{h\,m\,s}}$) |  $dx$ (′′) |  $\delta$ (∘ ′ ′′) |  $dy$ (′′) | (mJy beam-1) |  $\alpha$, $\delta$ ($\arcsec$, $\arcsec$) | (km s-1) | (km s-1)
1 | 18:25:44.89 | 0.02 | -12:23:01.55 | 0.02 | 20.1 (1.2) | 0.012, 0.044 | 60.2, 61.4 | 60.6
2 | 18:25:44.79 | 0.02 | -12:23:01.02 | 0.02 | 18.9 (1.3) | 0.073, 0.091 | 59.8, 61.4 | 60.2
3 | 18:25:44.85 | 0.03 | -12:22:37.44 | 0.03 | 9.5 (1.2) | 0.016, 0.022 | 57.4, 57.8 | 57.4
4 | 18:25:44.49 | 0.01 | -12:22:32.60 | 0.01 | 216.0 (1.3) | 0.262, 0.059 | 51.4, 61.8 | 59.8
5 | 18:25:44.53 | 0.01 | -12:22:30.54 | 0.01 | 59.0 (1.3) | 0.092, 0.030 | 57.8, 59.4 | 59.0
6 | 18:25:45.11 | 0.02 | -12:22:29.89 | 0.02 | 26.2 (1.3) | 0.048, 0.019 | 59.0 , 59.4 | 59.4
7 | 18:25:45.69 | 0.03 | -12:22:29.73 | 0.03 | 12.4 (1.3) | 0.039, 0.067 | 58.6, 59.4 | 59.4
8 | 18:25:44.60 | 0.03 | -12:22:27.85 | 0.03 | 11.9 (1.3) | 0.041, 0.037 | 59.4, 59.8 | 59.8

a Labelled from south to north, see Figure 4(d).

b Intensity-weighted centroid position, and intensity-weighted mean
statistical error from the Gaussian fitting, of each emission group.

c Fitted peak intensity of the brightest emission spot in each group (i.e.
from a single channel). The statistical error from the Gaussian fitting in
listed in parentheses. $T_{B}$(K) $\approx 7186\times I$(Jy beam-1).

d The angular spread of the emission group, calculated as the standard
deviation of the difference between the position of each individual emission
spot (i.e. from a single channel) and the intensity- weighted centroid
position of the group.

e Velocity spread over which $\geq 5\sigma$ emission is present, see §3.2.3.

f Velocity of the channel at which the brightest emission in each group
appears.

To extract the position, peak intensity and velocity of each $\geq 5\sigma$
emission spot, we fit the observed emission with a 2D Gaussian using the casa
imfit task (as in Towner et al., 2021). The rms noise is measured in each
channel to ascertain the signal-to-noise of the detection as the noise is
higher in channels with bright emission due to dynamic range limitations. As
we expect the emission to be unresolved, we fit each emission spot as a point
source by fixing the size to that of the synthesised beam (e.g. Hunter et al.,
2018; Towner et al., 2021). In total, we identify 50 emission spots at the
$\geq 5\sigma$ level (where the mean $\sigma$ across all channels with
emission is $\sigma=1.14$ mJy beam-1) that reside in 8 emission groups. The
position of each emission spot is plotted in Figure 6, each maser group is
labelled in Figure 4(d), and Table 5 lists the properties of each group.
Unlabelled contours in Figure 4(d) correspond to emission that does not meet
our criteria of $\geq$5$\sigma$ emission in $\geq$2 consecutive velocity
channels within a VLA beam (generally because the emission centroids shift by
more than a beam in consecutive channels) and so is not included in our
analysis. The strongest emission group (ID 4 in Table 5) has a fitted peak
intensity ($216.0\pm 1.3$ mJy beam-1) corresponding to a brightness
temperature $\sim$1550 K, strongly suggestive of masing behaviour. Four
additional emission groups (i.e. IDs 1, 2, 5 and 6) also exhibit brightness
temperatures greater than the $E_{u}/k_{B}$ of the line. While the brightest
group (ID 4) spans both blue- and red-shifted velocities (Table 5), blue-
shifted and red-shifted NH3(3,3) emission groups generally reside towards the
blue- and red-shifted outflow lobes respectively. The same pattern was
observed in 44 GHz Class I CH3OH masers by Cyganowski et al. (2009) (their
Figure 5f), though the maser velocities are modest compared to the velocity
extent of the thermal molecular outflow lobes (Cyganowski et al., 2009,
2011a). All NH3(3,3) emission spots are coincident both spatially and
kinematically (within a beam and a channel, respectively) with 44 GHz Class I
CH3OH maser spots from Cyganowski et al. (2009) (see Figure 6), similar to the
results seen in the EGO G35.03$+$0.35 by Brogan et al. (2011). We note that in
G19.01$-$0.03, not all 44 GHz methanol masers are coincident with NH3(3,3)
masers.

In G19.01$-$0.03, all identified NH3(3,3) emission spots are located in the
outer portions of the bipolar outflow lobes, whilst the 44 GHz Class I
methanol maser spots are also found closer to the driving source. This
distinction is shown in Figure 6, which plots the normalised, cumulative
distribution of the angular separation ($\theta_{\mathrm{sep}}$) of each
NH3(3,3) and 44 GHz Class I CH3OH maser spot from the ALMA 1.05 mm dust
continuum peak of MM1. The two-sample Kolmogorov-Smirnov (K-S) and $k$-sample
Anderson-Darling (A-D) tests are both nonparametric null hypothesis tests
often used to compare two such distributions. The null hypothesis in this case
states that the two distributions may be drawn from the same underlying
distribution if the tests return $p$-values larger than 0.05. Running both the
K-S and A-D tests on the full range of angular separations
($\theta_{\mathrm{sep}}$) shown in Figure 6 returns vanishingly small
$p$-values, meaning the null hypothesis may be rejected and a conclusion drawn
that the NH3(3,3) and 44 GHz methanol maser distributions are significantly
different. However, running both tests for $\theta_{\mathrm{sep}}\geq
12.5\arcsec$ returns $p\gg 0.05$, meaning that the two distributions at large
angular separations are effectively indistinguishable and may be drawn from
the same underlying distribution.

Figure 6: Left panel: Fitted positions of each NH3(3,3) spot across all groups
(blue $\times$), and 44 GHz Class I CH3OH maser spot positions from Cyganowski
et al. (2009) (orange filled circles). ALMA 1.05 mm continuum contours are
plotted in black (at [5, 8, 16, 64, 600]$\times\sigma$, where $\sigma=0.25$
mJy beam-1). High-velocity, blue- and red-shifted 12CO(2–1) emission observed
with the SMA (Cyganowski et al., 2011a) are plotted in blue and red contours,
as in Figure 1. The VLA beam and a scalebar are plotted in the bottom left and
top left of the panel respectively. Right panel: Normalised, cumulative
histogram of the angular separation (in arcseconds) of each NH3(3,3) spot
(blue), and each 44 GHz Class I CH3OH maser spot (orange), from the ALMA 1.05
mm continuum peak of MM1 (Table 3). Angular separations were calculated using
the au.angularSeparation task of the analysisUtils python package.

We note that thermal NH3(1,1) and (2,2) emission are detected towards
G19.01–0.03 at 73$\arcsec$-resolution by Cyganowski et al. (2013) but are
undetected in our VLA observations (to $5\sigma$ limits of 6.2 and 5.9 mJy
beam-1 respectively). Our non-detection of these low-J transitions is
primarily attributable to the limited brightness temperature sensitivity of
our high-resolution observations (Table 2), with spatial filtering potentially
contributing to the non-detection of the NH3(2,2) line (the maximum
recoverable scale of our observations is only $\sim$4.5$\arcsec$ $\sim$0.08 pc
at D=4 kpc, see Table 1). Cyganowski et al. (2013) also detect NH3(3,3)
towards G19.01–0.03 that is likely dominated by thermal emission, presumably
arising near MM1. Extended NH3(3,3) emission is also seen towards the mm
continuum sources in the EGO G35.03$+$0.35 by Brogan et al. (2011) in
$3.7\arcsec\times 3.0\arcsec$ resolution VLA observations ($\sim$ $8,580\times
6,960$ au at the parallax distance of 2.32${}^{+0.24}_{-0.20}$ kpc from Wu et
al., 2014). In our data, NH3(3,3) emission at $\geq$5$\times$ the median rms
in Table 2 is present in a single channel within the area of MM1’s mm
continuum emission. This channel, however, has bright maser emission and so an
elevated rms, and the NH3(3,3) peak is $<$5$\times$ the channel rms. Since the
NH3(3,3) peak is also offset from the mm continuum peak and our tentative
NH3(6,6) and (7,7) detections (see below), we conclude that NH3(3,3) is
undetected towards MM1 in our data. NH3(5,5) is also undetected in our VLA
observations, to a $5\sigma$ limit of 5.5 mJy beam-1.

NH3(6,6) emission (shown in black contours in Figure 4(d)) is detected at the
$5.7\sigma$ level (I${}_{\rm peak}=$5.4 mJy beam-1) towards a position in the
outskirts of the northern, blue-shifted outflow lobe
(18h25m44$\aas@@fstack{s}$50 $-$12∘22′32$\aas@@fstack{\prime\prime}$52
(J2000)), spatially and kinematically coincident with an NH3(3,3) maser group
(ID 4 in Table 5), 44 GHz Class I CH3OH masers (Cyganowski et al., 2009), and
candidate ALMA 278.3 GHz Class I CH3OH maser emission (Section 3.2.2). Towards
MM1, NH3(6,6) is detected at the $5.6\sigma$ level (I${}_{\rm peak}=$5.3 mJy
beam-1), spatially and kinematically coincident with a tentative $4.4\sigma$
(4.6 mJy beam-1) detection of NH3(7,7). These tentative high-J detections
towards MM1 could represent tentative evidence of emission from hot gas
originating in the inner regions of the circumstellar accretion disc around
the central MYSO(s) (Paper I).

####  3.2.4 25 GHz methanol emission with the VLA

The four CH3OH lines in our VLA tuning (Table 2) can exhibit thermal and/or
Class I maser emission and are commonly detected towards EGOs: in their study
of 20 EGOs, Towner et al. (2017) detected thermal and/or maser emission
towards 16 (80%) of their sources. As Towner et al. (2017) adopt a 4$\sigma$
detection criterion, we consider $\geq$4$\sigma$ 25 GHz CH3OH emission so that
we can compare to their results; we also adopt their shorthand notation of 32,
52, 82 and 102 for these lines. Towards MM1, we tentatively detect all four
lines at the 4.8, 5.5, 4.5 and 4.5$\sigma$ levels, respectively (see Figure 7,
and Table 2 for $\sigma$ values), equivalent to brightness temperatures
($T_{B}$) of 32.6, 36.8, 23.2 and 32.4 K. Towner et al. (2017) detect all four
25 GHz CH3OH lines in half of their sample (i.e. $\sim$63% of those with a
detection of at least one line). The strongest emission towards MM1 is in the
52 line, a result also reported by Towner et al. (2017) towards the majority
of the sources in their sample. We also detect 7.8$\sigma$ 52 emission
($T_{B}=52.3$ K) offset by $6\arcsec\sim 24,000$ au to the north-west of MM1
(see Figures 4(d) and 7). This 52 emission is positionally and kinematically
coincident with outflow-tracing H2CO and CH3OH emission observed with ALMA
(§3.2.2 and Figure 4) and with 44 GHz Class I CH3OH maser emission (Cyganowski
et al., 2009).

Figure 7: VLA peak intensity maps of the four $\sim$25 GHz CH3OH lines in our
VLA tuning (Table 2), here referred to by the short hand notation 32, 52, 82
and 102. Each panel shows the peak intensity map of the labeled transition in
colourscale and black contours, with contour levels of 4$\sigma$ for the 32,
82 and 102 lines (where $\sigma=1.04$, 0.85 and 1.19 mJy beam-1 respectively),
and [4, 5, 6, 7]$\times\sigma$ for the 52 line (where $\sigma=1.01$ mJy
beam-1). ALMA 1.05 mm continuum contours are overplotted in white at 5, 8, 16,
64 and 600$\sigma$ (where $\sigma=0.25$ mJy beam-1). Magenta crosses mark the
positions of 44 GHz Class I CH3OH masers from Cyganowski et al. (2009). Each
synthesised beam, and a $1\arcsec$ scale bar, are plotted in the bottom left
and bottom right of each panel respectively. Figure 8: ALMA spectra towards
the 1.05 mm continuum peaks of MM2, MM3, MM4 and MM5 (Table 3) for the two
broad spectral windows. The break in the frequency axis between the two broad
spectral windows is marked by the bold black lines. The $5\sigma$ level (where
$\sigma=3$ mJy beam-1, §2.1) is marked by the horizontal dashed black line.
Lines with emission $>5\sigma$ are labelled.

To distinguish between thermal and maser 25 GHz CH3OH emission, we consider
the two approaches detailed by Towner et al. (2017): (i) following criteria on
the spatial and spectral extent of the line emission, and (ii) comparing
observed line intensities to predictions for optically thin LTE emission. For
the first approach, spatially and spectrally broad emission (i.e. resolved
$\geq 4\sigma$ emission in $\geq 5$ channels $\sim 2$ km s-1) is considered
thermal, whilst unresolved/point-like spectrally narrow emission (i.e. $\geq
4\sigma$ emission in $\leq 4$ channels $\sim 1.6$ km s-1) is considered a
candidate maser. The strongest 52 emission towards the MM1 outflow appears
both unresolved (i.e. with spatial extent less than a beam) and spectrally
narrow with a fitted full-width half-maximum velocity width ($\Delta V_{\rm
FWHM}$) of $0.6\pm 0.1$ km s-1. We therefore consider this a candidate maser.
Towards MM1, the emission from all four 25 GHz CH3OH lines appears spatially
unresolved. Both broad and narrow Gaussian components are needed to describe
the 52 emission towards MM1, with velocity widths of $3.5\pm 0.1$ and $0.7\pm
0.1$ km s-1 respectively, while the spectral profiles of the 32, 82 and 102
lines are best fit by single Gaussians with velocity widths of $3.4\pm 0.1$,
$6.0\pm 0.2$ and $6.7\pm 0.2$ km s-1 respectively. All of the fitted Gaussian
components have centroid velocities consistent with the systemic velocity of
MM1 (59.9$\pm$1.1 km s-1; Cyganowski et al., 2011a). The 82 and 102 lines also
have velocity widths of a similar order to the thermal CH3OH emission seen
with ALMA towards MM1 (see left panel of Figure 5). Line models under LTE and
optically thin conditions (produced using the Weeds package of the class
software) reasonably describe the broad emission components of all four lines
towards MM1. As a whole, our data suggest that both thermal (broad components)
and weak maser (narrow 52 component) may be present towards MM1, though we
emphasise that this is a tentative result due to the low S/N of the 25 GHz
CH3OH emission towards MM1.

####  3.2.5 MM2…MM5 with ALMA and the VLA

In Figure 8, we show spectra for the broad ALMA spectral windows towards the
1.05 mm continuum peaks of MM2, MM3, MM4 and MM5. The lines shown in Figure
4(c)-(e) lie within these broad bands. MM3 and MM5 appear devoid of molecular
line emission in these bands, whilst MM2 and MM4 both have $>5\sigma$
detections of a handful of lines (labelled in Figure 8). All of the line
emission detected towards MM2 and MM4 in these broad bands appears spatially
and kinematically coincident with emission attributed to the outflow driven by
MM1 (including the H2CO and CH3OH lines shown in Figures 4(c), (d) and (e);
see §3.2.2). As such, it is likely that this line emission is associated with
the MM1 outflow, rather than being physically associated with MM2 and MM4.
N2H+(3–2) emission (see Figure 4(a)) is coincident with all four new
millimetre sources, however given the emission does not morphologically
resemble the ALMA 1.05 mm continuum, it is unlikely to arise from the same
volumes as MM2…MM5. As shown in Figure 4(f), DCN($4-3$) does exhibit an
emission morphology similar to the ALMA 1.05 mm continuum. MM2 and MM4 are
detected in emission at the 10$\sigma$ and 9$\sigma$ levels respectively
(where $\sigma=6.3$ mJy beam-1), and do not exhibit signs of self-absorption
(unlike MM1, see §3.2.2), whilst MM3 and MM5 are undetected to a $5\sigma$
limit of 31.5 mJy beam-1. The positions of peak DCN emission towards MM2 and
MM4 are offset from their corresponding ALMA 1.05 mm continuum peaks by 0.09
and 0.40$\arcsec$ respectively (equivalent to 360 and 1600 au), within a DCN
beam ($\sim$0.47$\arcsec$; see Table 2). Systemic velocities for MM2 and MM4
are found from Gaussian fitting of the DCN peak emission to be $59.88\pm 0.06$
km s-1 and $59.72\pm 0.05$ km s-1 respectively, consistent with the $59.9\pm
1.1$ km s-1 systemic velocity of MM1 (Cyganowski et al., 2011a), kinematically
linking MM2 and MM4 to MM1 and the G19.01$-$0.03 clump.

A handful of Class I 44 GHz CH3OH maser spots lie $0.60\arcsec\sim 2400$ au to
the east of the ALMA 1.05 mm continuum peak of MM4 (see Figure 4(e)),
coincident spatially and kinematically with H2CO and CH3OH emission attributed
to the MM1 outflow (see §3.2.2). No NH3($J$ = $K$ = 1, 2, 3, 5, 6, 7) emission
is detected towards, or in the immediate vicinity of, MM2…MM5 in our VLA data,
to $5\sigma$ limits of 6.2, 5.9, 5.7, 5.5, 4.8, 5.3 mJy beam-1 respectively.
Similarly, MM2…MM5 are undetected in the 25 GHz CH3OH lines in our VLA
observations ($5\sigma$ limits 5.2, 5.1, 4.3 and 6.0 mJy beam-1 for the 32,
52, 82 and 102 lines). The CH3OH 52 emission spot detected towards the outflow
driven by MM1 (see §3.2.4) is $3.6\arcsec$ $\sim$ $14,200$ au away from the
ALMA 1.05 mm continuum peak of MM3. In all, with no clear evidence for
protostellar activity in MM2…MM5 in our ALMA and VLA data, the current
evidence suggests that these sources may be starless condensations, prestellar
cores, millimetre knots in MM1’s outflow lobes, or a combination of the three
(see also Section 4.2).

##  4 Discussion

###  4.1 Physical properties of MM1 from CH3OH emission

In the following section, we estimate the column density and temperature of
MM1 from the observed CH3OH emission using two complementary methods: (i) a
rotation diagram analysis (§4.1.1), and (ii) a set of synthetic spectra that
represent the observed emission (§4.1.2).

####  4.1.1 Rotation diagram analysis of CH3OH

Assuming local thermodynamic equilibrium and optically thin emission, a single
temperature may describe the emission of all observed lines. This temperature
is defined by the Boltzmann distribution, written as:

$\log\left(\frac{N_{u}}{g_{u}}\right)=\log\left(\frac{N}{Q(T_{\mathrm{rot}})}\right)-\frac{E_{u}}{k_{B}T_{\mathrm{{rot}}}}\,,$
(1)

where $N_{u}$ is the upper energy level column density, $g_{u}$ is the upper
energy state degeneracy, $N$ is the total column density,
$Q(T_{\mathrm{rot}})$ is the partition function, $E_{u}$ is the upper level
energy, $k_{B}$ is the Boltzmann constant, and $T_{\mathrm{rot}}$ is the
rotational temperature (e.g. Goldsmith & Langer, 1999). The left-hand side of
equation 1 may be re-written as:

$\log\left(\frac{N_{u}}{g_{u}}\right)=\log\left(\frac{3k_{B}\int
T_{\mathrm{mb}}\mathrm{dV}}{8\pi^{3}\nu S_{{}_{ij}}\mu^{2}}\right)\,,$ (2)

where $\int T_{\mathrm{mb}}\mathrm{dV}$ is the integrated line strength over
velocity with units of $\mathrm{K}\times\mathrm{km\,s^{-1}}$, $\nu$ is the
rest frequency of the transition, $\mu$ is the permanent dipole moment with
units of Debye, and $S_{{}_{ij}}$ is the unit-less intrinsic line strength
(e.g. Cummins et al., 1986; Herbst & van Dishoeck, 2009). When multiple lines
of a species are detected, assuming that their emission arises from the same
physical structure, and assuming that all lines are optically thin, one may
construct a rotation diagram by plotting $\log(N_{u}/g_{u})$ against
$E_{u}/k_{B}$. Fitting a straight line will yield $T_{\mathrm{rot}}$ from the
slope, and $N/Q(T_{\mathrm{rot}})$ from the y-intercept. If a straight line
reasonably fits the rotation diagram, the assumption is made that all
transitions are thermalised, and thus $T_{\mathrm{rot}}$ is expected to equal
the excitation temperature, $T_{\mathrm{ex}}$ (e.g. Goldsmith & Langer, 1999).
Taking this $T_{\mathrm{rot}}$, the total column density ($N$) may be found
with the partition function, which following Townes & Schawlow (1955) and
Purcell et al. (2009), for example, may be written for CH3OH as:

$Q(T_{\mathrm{rot}})=1.2327\,T_{\mathrm{rot}}^{1.5}\,\,.$ (3)

The most detected species towards MM1 in our spectral tuning at $>$5$\sigma$
is CH3OH, with thirteen transitions identified across both the wide and narrow
bands (see §3.2.1) with $E_{u}/k_{B}$ ranging $32-736$ K. We produce a
rotation diagram towards the ALMA 1.05 mm continuum peak of MM1 (Table 3)
using eleven CH3OH transitions, excluding the CH3OH($15_{1,0}-14_{2,0}$) line
from the wide bands, and CH3OH($6_{1,2}-5_{1,2}$) from the narrow bands, as
they are significantly blended by the H2CO(42,3–32,2) and
CH3OCHO($27_{1,27}-26_{0,26}$) lines respectively. We include both A and E
transitions in this analysis since we do not detect enough transitions to to
fit them separately. The integrated intensity of each line is evaluated from
Gaussian profiles fitted to the CH3OH spectra at the ALMA 1.05ṁm continuum
peak position (see Table 3). Errors on the Gaussian fits are propagated
through to the rotation diagram, and the best fit is found following a
weighted least-squares fit. This rotation diagram is shown in Figure 9(a),
where it can be seen that two data points (marked white, belonging to the
CH3OH($6_{1,5}-5_{1,4}$) and CH3OH($9_{-1,9}-8_{0,8}$) lines with
E${}_{u}/k_{B}=63.7$ and 110.0 K respectively) deviate from the general trend
of the other nine lines. These lines are seen in Figure 4 to be outflow
tracing, and their generally low $\log\left(N_{u}/g_{u}\right)$ is an
indication that they may also be optically thick. The derived rotational
temperature is $186\pm 8$ K.

As previously stated, a rotation diagram is constructed under the assumption
that all lines are optically thin. Optically thick lines can inflate derived
rotational temperatures due to an artificially shallow slope (e.g. Goldsmith &
Langer, 1999). We counter this by applying an opacity correction,
$C_{\tau_{i}}=\tau_{i}/(1-e^{\tau_{i}})$, where the subscript $i$ refers to
the $i^{\mathrm{th}}$ CH3OH transition. We follow Brogan et al. (2007); Brogan
et al. (2009) by iteratively solving for the $\tau_{i}$ and $T_{\mathrm{rot}}$
values that produce the best fit, minimising the $\chi^{2}$ value. Figure 9(b)
shows the best fitting opacity-corrected rotation diagram. The $\tau_{i}$
values for this best fit range between $\sim 0.4-11.7$, with indeed the
CH3OH($6_{1,5}-5_{1,4}$) and CH3OH($9_{-1,9}-8_{0,8}$) lines being the most
optically thick (with $\tau=11.7$ and 11.3 respectively). The highest-J,
CH3OH($23_{4,19}-22_{5,18}$) line with E${}_{u}/k_{B}=736$ K is unsurprisingly
the most optically thin. Of the lines with compact emission morphologies (i.e.
that trace MM1 and not the outflow) in the 278 GHz wideband spectral window,
the CH3OH($14_{4,10}-15_{3,12}$) line is the most optically thick, with
$\tau=1.9$, implying an emitting region of
$\sim$0$\aas@@fstack{\prime\prime}$25$\sim$ 1000 AU (following Section 4.4 of
Brogan et al., 2009). The derived opacity-corrected rotational temperature is
$166\pm 9$ K, and the column density is $(2.0\pm 0.4)\times 10^{18}$ cm-2.

Figure 9: Rotation diagram of CH3OH at the ALMA 1.05 mm dust continuum peak of
MM1, (a) before opacity correction, and (b) following opacity correction. The
best fit is shown by the dashed black line, with the statistical uncertainty
on the slope and y-intercept represented by the shaded blue region. The
CH3OH($6_{1,5}-5_{1,4}$) and CH3OH($9_{-1,9}-8_{0,8}$) lines (with
$E_{u}/k_{B}=63.7$ and $110.0$ K respectively) are coloured in white as they
are outflow tracers (see Figure 4) and optically thick.

####  4.1.2 Synthetic CH3OH spectra

We make a second estimate of the column density and temperature of MM1 by
producing synthetic spectra that are representative of the observed data using
the weeds extension (Maret et al., 2011) of the class software. weeds solves
the radiative transfer equation assuming a state of local thermodynamic
equilibrium (LTE), and takes into account the background continuum emission.
Using values for the frequency, Einstein A coefficients, partition function,
and upper level degeneracy and energy from the JPL and CDMS catalogues
(Pickett et al., 1998; Müller et al., 2001), Weeds calculates the opacity of
each line, meaning that corrections such as line broadening are applied to
optically thick lines. The free parameters of the synthetic spectra are the
column density, excitation temperature, systemic velocity, FWHM velocity
width, and source size. The projected diameter of the telescope is also
required to calculate the beam dilution/filling factor. We refer the reader to
Maret et al. (2011) for a more detailed description of the full procedure.

We explore the posterior distribution of the free parameters of the LTE line
models using the Python package emcee (Foreman-Mackey et al., 2013), an
affine-invariant Monte Carlo Markov Chain (MCMC) sampler (Goodman & Weare,
2010). A posterior distribution of model parameters consistent with the data
is evaluated by multiple “walkers”, allowed to walk for a number of
iterations. A “burn-in” period of $n_{\mathrm{burn}}$ iterations allows the
walkers to converge on the posterior region. These “burn-in” solution chains
are then discarded before beginning a “production” run with walkers
initialised around the posterior region identified during the burn-in. Walkers
are allowed to walk through the posterior for
$n_{\mathrm{prod}}>n_{\mathrm{burn}}$ iterations, to allow probing of the
covariance of the model parameters. The model parameters that are most
representative of the data are found following the maximisation of the log-
likelihood function:

$\mathcal{L}(T_{\mathrm{mb}}|\theta)=-0.5\sum_{i}\left(\frac{T_{\mathrm{mb},i}^{\mathrm{mod}}-T_{\mathrm{mb},i}^{\mathrm{obs}}}{\sigma^{\mathrm{obs}}_{\mathrm{rms}}}\right)^{2}\,,$
(4)

where $\theta$ are the model parameters, $T_{\mathrm{mb},i}^{\mathrm{mod}}$
and $T_{\mathrm{mb},i}^{\mathrm{obs}}$ are the model and observed main-beam
brightness temperature of the $i^{\mathrm{th}}$ channel, respectively, and
$\sigma^{\mathrm{obs}}_{\mathrm{rms}}$ is the observed rms noise of the
spectrum (see §2.1). The statistical uncertainty on each model parameter is
evaluated from the 0.16 and 0.84 quantiles of the posterior (i.e. the
1$\sigma$ level). Our scripts are publicly
available222https://github.com/gwen-williams/WeedsPy_MCMC, and are packaged
together in a reusable form we call weedspy_mcmc where they are generalised to
take in user-defined parameters such as the priors, number of walkers, initial
walker positions within the parameter space, and the number of burn-in and
production iterations, from a text file. A class-readable text file is created
containing the $\theta$ model parameters currently walked to and the name of
the molecule to be modelled, and then a class script is called that runs the
LTE model for those $\theta$. Plots are generated of the highest likelihood
synthetic spectrum, the corner plot of the 1D and 2D posterior distributions
of the free parameters (as in Figure 10(a)), and the trace plot of the emcee
walkers through the parameter space. To efficiently sample the probability
distributions for the free parameters of the synthetic spectra generated with
weeds, spectral windows are treated separately in our scripts.

####  4.1.3 Pixel-by-pixel analysis of CH3OH synthetic spectra

Applying this procedure to our ALMA data, we find the highest likelihood model
on a pixel-by-pixel basis across MM1, allowing investigation of any spatial
variation in the column density, temperature, centroid velocity and velocity
width. As in §4.1.1, this analysis focuses on CH3OH emission, as it is the
species with the highest number of identified transitions in our tuning. Of
the thirteen total CH3OH lines in the tuning, seven are located in the
$\sim$278 GHz band, four are in the $\sim$292 GHz band (see Figure 3 and Table
4), one lies in the narrow spectral window that targets N2H+($3-2$), and one
lies in the narrow spectral window that targets DCN($4-3$). As the
CH3OH($6_{1,5}-5_{1,4}$) and CH3OH($9_{-1,9}-8_{0,8}$) lines (with
$\nu_{\mathrm{rest}}=292.67291$ and 278.30451 GHz respectively) are shown in
Figure 4 to be outflow tracing, we exclude them from this analysis since an
assumption made here is that the emitting region is the same for all modelled
transitions. We also exclude the CH3OH($15_{1,0}-14_{2,0}$) line
($\nu_{\mathrm{rest}}=291.24057$ GHz) as it is heavily blended by a close-by
H2CO line (see also §4.1.1), the CH3OH($6_{1,2}-5_{1,2}$) line
($\nu_{\mathrm{rest}}=289.62430$ GHz) as it is blended with a CH3OCHO line
(see §3.2.1 and §4.1.1), and the CH3OH($10_{1,10}-9_{0,9}$) line
($\nu_{\mathrm{rest}}=292.51744$ GHz) as it is the only identified line in the
first torsionally excited state (i.e. $v_{t}=1$). This leaves 6 lines in the
$\sim$278 GHz band, 1 line in the $\sim$292 GHz band, and 1 line in a narrow
band. Since modelling one line is not sufficient to constrain four free
parameters, we conduct our MCMC analysis on the six remaining CH3OH lines in
the $\sim$278 GHz band (indicated in bold in Table 4). We then use the best-
fit parameters to generate synthetic spectra for the other spectral windows
and check the resulting spectra for goodness of fit (see also Section 4.3.1).

Figure 10: (a) One-dimensional and two-dimensional histograms of the posterior
distributions for the free parameters in the LTE synthetic spectral modelling
(N(CH3OH), Tex, $\mathrm{V}_{\mathrm{cen}}$ and $\Delta\mathrm{V}$) towards
the ALMA 1.05 mm continuum peak. Black contours on the 2D histograms are
placed at 2, 1.5 and 1 $\sigma$, whilst the dashed black lines on the 1D
histograms are placed at the 0.16, 0.5 and 0.84 quantiles. (b) Observed
spectrum towards the ALMA 1.05 mm continuum peak (black) overplotted with the
corresponding highest likelihood model (red) i.e. model parameters defined by
the mode of the posterior, printed in panel (a). The six CH3OH lines that are
modelled and included in the maximisation of the log-likelihood function are
marked by vertical grey lines. The excluded CH3OH($9_{-1,9}-8_{0,8}$) line
with $E_{u}/k_{B}=110.0$ K is labelled.

For our pixel-by-pixel modelling, we fix the source size to the geometric mean
of the ${\sim}$278 GHz synthesised beam (${\sim}0.5\arcsec$).333The image cube
is converted to brightness temperature using the Rayleigh-Jeans approximation
and beamsize with the tt.brightnessImage function of toddTools. We also limit
the pixel area being modelled to within the continuum source size of MM1 from
Table 3 (see white contour in Figures 11(a) and (b)), shown in Paper I to
match well the region of compact molecular line emission (see their Fig.3).
The projected diameter of the ALMA array was calculated to be 558
metres444Calculated using the au.getBaselineLengths function from the
analysisUtils Python package, and we estimate the background continuum on a
pixel-by-pixel basis (McGuire et al., 2018), with the ALMA 1.05 mm continuum
converted to brightness temperature555Conversion done under the Rayleigh-Jeans
approximation using the tt.brightnessImage function of toddTools.. We
initialise 60 walkers, and allow them to run for $n_{\mathrm{burn}}=200$ and
$n_{\mathrm{prod}}=500$ iterations (see Appendix B for a discussion of these
choices). The remaining free parameters are the column density, excitation
temperature, systemic velocity and FWHM velocity width. We assume uniform,
uninformative priors, and initialise the walkers in a four-dimensional sphere
about an initial value that lies within those priors. Our priors are estimated
from initial “dummy” runs of the LTE synthetic spectra outside of
weedspy_mcmc, and are set to be $(0.05-3.0)\times 10^{18}$ cm-2 for the column
density, $80-300$ K for the temperature, $56.0-64.0$ km s-1 for the centroid
velocity, and $2.0-8.0$ km s-1 for the velocity width.

Figure 10(a) shows the 1D and 2D posterior distributions of column density,
temperature, centroid velocity and velocity width, evaluated for the spectrum
at the ALMA 1.05 mm continuum peak of MM1. The narrow, normally distributed 1D
posteriors for column density, temperature and centroid velocity show that
these parameters are reasonably constrained. The slope in the 2D column
density–temperature posterior distribution indicates an unsurprising, but only
slight, degeneracy between these two parameters, meaning the statistical
errors are small despite this. There is no obvious covariance of the centroid
velocity with either column density or temperature. The velocity width is less
well constrained, with two slight peaks around the mode of the posterior
distribution (within $<0.03$ km s-1). Though significantly smaller than the
spectral resolution of our data ($\sim$1 km s-1), this uncertainty is
attributed to line blending rather than under-sampling as the double peak
persists in tests where $n_{\mathrm{prod}}$ is increased to $>$3000\. The
double peak also persists when walkers are initialised with a different random
number seed, so is independent of initial walker position. The highest
likelihood model spectrum at the continuum peak position of MM1 is shown in
Figure 10(b), with parameters N$($CH3OH$)=(2.22\pm 0.01)\times 10^{18}$ cm-2,
T${}_{\mathrm{ex}}=162.7\begin{subarray}{c}+0.3\\\ -0.5\end{subarray}$ K,
V${}_{\mathrm{cen}}=59.73\pm 0.01$ km s-1 and $\Delta
V=6.38\begin{subarray}{c}+0.01\\\ -0.03\end{subarray}$ km s-1. These remain
unchanged (within the errors) when walkers are initiated with a different
random number seed, and when $n_{\mathrm{prod}}$ is increased to $>3000$.
These column density and temperature values are remarkably consistent with
those derived from the opacity-corrected rotation diagram of $(2.0\pm
0.4)\times 10^{18}$ cm-2 and $166\pm 9$ K respectively (see §4.1.1), and the
centroid velocity is also consistent with the known systemic velocity of
$59.9\pm 1.1$ km s-1 (Cyganowski et al., 2011a). The line optical depths
returned by class are lower than those estimated in the rotation diagram
analysis with an average $\tau=0.3$ compared to $\tau=1.0$ (excluding
CH3OH($9_{-1,9}-8_{0,8}$) with $E_{u}/k_{B}=110.0$ K). We attribute the higher
line opacities estimated from the rotation diagram to the source size not
being fixed in the rotation diagram approach (see §4.1.1). Figure 10(b) shows
the synthetic spectrum well reproduces the observed CH3OH emission, other than
significantly over-estimating the line strength of the optically thick,
outflow tracing CH3OH($9_{-1,9}-8_{0,8}$) line (see Figure 4) that was
excluded from the maximisation of the log-likelihood. This vindicates its
exclusion from the analysis, further suggesting that the emission from the
CH3OH($9_{-1,9}-8_{0,8}$) line may not occur in the same volume as the other
modelled, optically thin lines.

Figure 11: Pixel-by-pixel highest likelihood model parameters from the Weeds
CH3OH synthetic spectra towards MM1: (a) centroid velocity in km s-1, (b)
CH3OH column density in cm-2, (c) excitation temperature in Kelvin, and (d)
FWHM velocity width in km s-1. ALMA 1.05 mm continuum contours are plotted in
panels (a) and (b) at 8, 16, 64 and 200$\sigma$ (where $\sigma=0.25$ mJy
beam-1). The white contour (at $64\sigma$) represents the level above which
pixels were fed into weedspy_mcmc, and matches the source size of MM1 (see
Table 3). Contours of VLA 5.01 cm emission (from Paper I) are plotted in
panels (c) and (d) at the 4 and 5$\sigma$ levels (where $\sigma=5.0$ $\mu$Jy
beam-1) in panel. A scale bar and the ALMA beam are plotted in panel (a), and
the VLA beam is plotted in panel (c).

####  4.1.4 Images of CH3OH physical parameters

Maps of the pixel-by-pixel, highest likelihood CH3OH model parameters towards
MM1 are shown in Figure 11. The highest likelihood centroid velocity map
(Figure 11a) shows a clear gradient in velocity across MM1 in agreement with
the gradient reported in Paper I from a simpler moment analysis of individual
lines. The CH3OH column density (Figure 11b) peaks coincident with, and
follows the same morphology as, the ALMA 1.05 mm continuum emission. The
position of peak CH3OH temperature (18h25m44$\aas@@fstack{s}$763
$-$12∘22′46$\aas@@fstack{\prime\prime}$00 (J2000); see Figure 11c) is however
offset from the ALMA 1.05 mm dust continuum and CH3OH column density peaks by
$0.22\arcsec\sim 880$ au. Whilst the column density and temperature are
somewhat degenerate (see Figure 10a), the difference between the offset peak
temperature ($165.5\pm 0.6$ K) and the temperature at the column density peak
($162.7\pm 0.4$ K) is significantly greater than its propagated statistical
error (i.e. $2.8\pm 0.7$ K). In Paper I, a centimetre counterpart to MM1
(called CM1) was detected for the first time with the VLA at 5.01 and 1.21 cm.
The VLA 1.21 cm emission peak was coincident with the ALMA 1.05 mm continuum
peak of MM1, whilst the VLA 5.01 cm emission peak of CM1
(18h25m44$\aas@@fstack{s}$773 $-$12∘22′46$\aas@@fstack{\prime\prime}$00
(J2000); Paper I) was found to be offset from MM1 by 0.16$\arcsec\sim 640$ au
(Paper I). Indeed, the entire morphology of the region of elevated CH3OH
temperature matches that of the VLA 5.01 cm emission, with both being
elongated along a south-easterly direction relative to MM1 (as shown in Figure
11c). The velocity width map in Figure 11(d) also shows a similar elongation
to the VLA 5.01 cm continuum, however it does not peak coincident with either
the CH3OH column density, the CH3OH temperature, nor the VLA 5.01 cm continuum
position. Following analysis of the combined centimetre-millimetre spectral
energy distribution of CM1/MM1 in Paper I, it was revealed that a free-free
emission component was required to describe the VLA 5.01 cm emission. This
component was interpreted as a small, 66 au diameter gravitationally trapped
hypercompact (HC) Hii region, which directly implies the presence of a source
causing ionisation and heating of its surroundings. Given that the morphology
of the elevated CH3OH temperature aligns with that of the VLA 5.01 cm
continuum, and that the position of peak CH3OH temperature is only offset from
the VLA 5.01 cm position of CM1 by $0.07\arcsec\sim 280$ au (equivalent to the
absolute positional uncertainty of the VLA 5.01 cm data of 0.07$\arcsec$;
Paper I), it is reasonable to surmise that both are attributable to the same
underlying source/mechanism. We further speculated in Paper I that the VLA
5.01 cm emission being misaligned with respect to the direction of the bipolar
outflow emanating from MM1 was suggestive of the ionisation being driven by a
second object, perhaps indicative of an unresolved high-mass binary system. In
that case, a second possible origin for the offset elevated CH3OH temperature
could be the presence of an ionised jet driven by the unresolved high-mass
binary companion. Outflows and jet-like outflows are noted in the literature
to cause heating of molecular gas (e.g. Zhang et al., 2007; Wang et al.,
2012).

###  4.2 Physical properties of MM1…MM5 from dust emission

As mentioned in §1 and §4.1.4, it was shown in Paper I from an analysis of the
centimetre-millimetre spectral energy distribution of MM1 that its ALMA 1.05
mm flux was 99.99 per cent dominated by thermal dust emission. MM2…MM5 are
also likely dominated by thermal dust emission, given their non-detection in
the centimetre at 1.21 and 5.01 cm with the VLA to 5$\sigma$ limits of 30 and
25$\mu$Jy beam-1 respectively (Paper I). As such, the masses
($\mathrm{M_{\mathrm{gas}}}$) of MM1…MM5 may be calculated from their
integrated flux densities at 1.05 mm ($S_{1.05\,\mathrm{mm}}$) assuming
isothermal dust emission following Hildebrand (1983):

$\mathrm{M_{\mathrm{gas}}}=\frac{d^{2}\,R\,S_{1.05\,\mathrm{mm}}\,C_{\tau_{\mathrm{dust}}}}{\kappa_{1.05\,\mathrm{mm}}\,B_{\nu}(T_{\mathrm{dust}})}\,,$
(5)

where $d$ is the distance, $R$ is the gas-to-dust mass ratio (here assumed to
be 100), $\kappa_{1.05\,\mathrm{mm}}$ is the dust opacity at
$1.05\,\mathrm{mm}$, $B_{\nu}(T_{\mathrm{dust}})$ is the Planck function, and
$T_{\mathrm{dust}}$ is the dust temperature. We include a correction for dust
optical depth, $C_{\tau_{\mathrm{dust}}}$ as:

$C_{\tau_{\mathrm{dust}}}=\tau_{\mathrm{dust}}/(1-e^{-\tau_{\mathrm{dust}}})\,,$
(6)

where $\tau_{\mathrm{dust}}$ may be estimated as:

$\tau_{\mathrm{dust}}=-\ln\left(1-\frac{T_{b}}{T_{\mathrm{dust}}}\right)\,.$
(7)

We estimate $T_{b}$ as the mean brightness temperature across the size of each
source (as listed in Table 3), though this may underestimate
$\tau_{\mathrm{dust}}$ for small-scale structures. From Ossenkopf & Henning
(1994), we set $\kappa_{1.05\,\mathrm{mm}}=1.45$ cm2g-1 for dust grains with
thick ice mantles in high density gas (as in Paper I; Cyganowski et al.,
2017). We also assume that the dust is well-coupled to the gas such that
$T_{\mathrm{dust}}=T_{\mathrm{gas}}$. The temperatures we assume for each of
the sources is discussed below in §4.2.1 and §4.2.2. Assuming spherical
geometry, a mean molecular weight per hydrogen molecule
($\mu_{\mathrm{H}_{2}}$) of 2.8 (e.g. Kauffmann et al., 2008), and a radius
equal to half the geometric mean of the source sizes presented in Table 3, we
further calculate the average H2 column density and H2 volume density. All
derived source properties are listed in Table 6.

####  4.2.1 Temperature, mass and stability of MM1

In Paper I, we assumed a temperature range for MM1 of $100-130$ K based on the
CH3CN(J$=$12–11) line fitting of Cyganowski et al. (2011a) from 2.3$\arcsec$
angular resolution SMA data. Our new ALMA data (with 0.4$\arcsec$ angular
resolution) improves on that by a factor of $\sim$30 in beam area. We revise
our temperature assumption to 163 K, as calculated at the ALMA 1.05 mm
continuum peak of MM1 in §4.1.3 from the CH3OH synthetic spectra. This is an
increase of 33–66 K on the SMA-derived temperature. Observations of the EGO
G11.92–0.61 (hereafter G11.92) by both Cyganowski et al. (2011a) and Ilee et
al. (2016) provide us with a point of comparison on temperature. With
$2.4\arcsec$ angular resolution SMA observations, Cyganowski et al. (2011a)
derived 77 and 166 K temperatures for G11.92 from a two-component fit to CH3CN
emission. Taking the same approach but with $0.5\arcsec$ angular resolution
SMA observations (a factor of $\sim$27 improvement in beam area), Ilee et al.
(2016) find the two G11.92 temperature components to be 150 and 230 K. This is
an increase of 60–70 K in temperature, attributed to the difference in probed
spatial scales between the two data sets. The same resolution-dependent
behaviour is observed here for the temperature of G19.01 MM1.

Table 6: Derived properties for mm continuum sources.

Source | Tb | Tdust | $\tau_{\mathrm{dust}}$ | Mgas | N(H2) | n(H2)
---|---|---|---|---|---|---
|  |  |  |  | $\times 10^{23}$ | $\times 10^{6}$
| (K) | (K) |  | (M⊙) | (cm-2) | (cm-3)
MM1 | 10.9 | 163 | 0.07 | 4.3 | 6.7 | 17.0
MM2 | 3.5 | 24 | 0.16 | 1.3 | 3.2 | 10.3
|  | 15 | 0.27 | 2.7 | 6.6 | 21.1
MM3 | 3.0 | 24 | 0.13 | 1.0 | 1.5 | 3.9
|  | 15 | 0.22 | 2.1 | 3.1 | 7.8
MM4 | 2.9 | 24 | 0.13 | 0.9 | 1.4 | 3.4
|  | 15 | 0.22 | 1.8 | 2.7 | 6.8
MM5 | 2.8 | 24 | 0.13 | 0.1 | – | –
|  | 15 | 0.21 | 0.2 | – | –

Column 1: Source name. Col. 2: Planck brightness temperature averaged over the
source size listed in Table 3. Col. 3: Assumed dust temperature, derived from
the LTE synthetic spectra for MM1 (see §4.1.3), and a temperature range of
$15-24$ K assumed for MM2…MM5 based on large scale clump observations
(Schuller et al., 2009; Cyganowski et al., 2013; Elia et al., 2017). Col. 4:
Dust optical depth, calculated from the average $T_{b}$ in column 2. Col. 5:
Opacity-corrected gas mass, in solar masses. Col. 6: H2 column density, and
Col. 7: H2 volume density, both assuming spherical symmetry.

The mass of MM1 is here revised, from M${}_{\mathrm{gas}}=5.4-7.2$ M⊙ at
130–100 K derived in Paper I, to M${}_{\mathrm{gas}}=4.3$ M⊙ at 163 K (Table
6). Attributing this gas mass to the now known circumstellar disc around MM1,
the stability of the disc can be assessed by estimation of the disc-to-star
mass ratio (as done in Paper I), with unstable discs exhibiting typical values
of $M_{\mathrm{gas}}/M_{\ast}>0.1$ (e.g. Kratter & Lodato, 2016). Following
Paper I, the stellar mass ($M_{\ast}$) can be calculated as
$M_{\ast}=M_{\mathrm{enc}}-M_{\mathrm{gas}}$ where $M_{\mathrm{enc}}$ is the
enclosed mass within the disc outer radius (estimated as $\sim 40-70$ M⊙ from
kinematic modelling in Paper I). The disc-to-star mass ratio was found in
Paper I to be $\sim 0.08-0.22$ hinting that the disc could be unstable and be
undergoing fragmentation into as-of-yet undetected low-mass stellar
companions. It was noted however that should the temperature of MM1 increase
on the size-scales probed by the ALMA data (as it did for the SMA observations
of G11.92; Cyganowski et al., 2011a; Ilee et al., 2016), then the resulting
lower $M_{\mathrm{gas}}$ could push the disc towards stability. Our revised
$M_{\mathrm{gas}}/M_{\ast}$ is 0.07–0.12, indeed indicative of the disc being
potentially stable against fragmentation. This approach does however risk
underestimating the disc mass, as it does not account for variations in the
temperature, dust opacity and dust optical depth in the disc (Johnston et al.,
2015; Forgan et al., 2016).

####  4.2.2 Low mass protocluster members

As there are insufficient molecular lines detected with ALMA towards MM2…MM5
for the estimation of temperature, we use clump-scale observations of G19.01
to inform our temperature assumptions. Elia et al. (2017) derive
$T_{\mathrm{dust}}=17.9$ K from SED fitting of HI-GAL observations, whilst
Schuller et al. (2009) and Wienen et al. (2012) derive $T_{\mathrm{gas}}=19.5$
K from NH3 hyperfine structure fitting of $\sim$40$\arcsec$ angular resolution
Effelsberg 100 m data. Cyganowski et al. (2013) also estimated temperature for
the G19.01 clump using NH3 observations, taken with the Nobeyama 45 m
telescope. Their single-temperature fit to the NH3 spectra yielded a kinetic
temperature of $23.8\pm 0.4$ K, but they found that a two-component fit with
Tkin,cool = $14.7\pm 0.5$ K and Tkin,warm = $50.2\pm 4.4$ K better represented
the observed data (Cyganowski et al., 2013). As the warm $\sim$50 K component
is likely attributable to the central MYSO (MM1), we exclude it and use the
remaining temperature estimates to inform an assumed temperature range for
MM2…MM5 of $15-24$ K. Table 6 lists derived properties for
$T_{\mathrm{dust}}=$ 15 and 24 K. We note that H2 column and volume densities
are not included for MM5 as its size is smaller than a beam.

The gas masses of MM2…MM5 range between $0.1-1.3$ M⊙ and $0.2-2.7$ M⊙ at 24
and 15 K respectively (see Table 6), and their median radius is $\sim$2000 au
(excluding the unresolved MM5; see Table 3). As such, the physical properties
of these millimetre sources appear consistent with typical low-mass
pre/protostellar cores observed in infrared dark clouds (with masses and radii
of a few solar masses and a few thousand au respectively, e.g. Sanhueza et
al., 2019; Morii et al., 2021; Redaelli et al., 2022) rather than discs around
low-mass YSOs (with masses and radii $<$0.1 M⊙ and a few hundred au
respectively, e.g. Huang et al., 2018; Dullemond et al., 2018). Furthermore,
their lack of COM emission or signs of outflow activity indicate that they may
be candidate low-mass pre-stellar rather than protostellar cores.
Interestingly, this result contrasts with the situation in the EGO
G11.92$-$0.61, where several of the low-mass cores identified by Cyganowski et
al. (2017) drive molecular outflows, clearly identifying them as protostellar
cores.

The number and spatial distribution of the low-mass cores in G19.01$-$0.03
also contrasts with G11.92$-$0.61, where Cyganowski et al. (2017)’s 1.05 mm
ALMA observations of the ATLASGAL clump detected 16 new compact mm continuum
sources (for a total of 19, including 3 previously-known massive sources).
Compared to G11.92$-$0.61, the low-mass sources detected in G19.01$-$0.03 are
located closer to the most massive member of the protocluster (median
projected separation $\sim$0.08 pc compared to $\sim$0.17 pc) and comprise a
lower fraction of the clump mass ($<$1% compared to 3-5%; Cyganowski et al.,
2017). We consider our results in the context of expectations for thermal
Jeans fragmentation. For the simple assumption of a uniform density sphere,
the minimum mass necessary for a fragment to be bound, and so collapse, is

$M_{J}=\left(\frac{5\,c_{s}^{2}}{2\,G}\right)^{3/2}\left(\frac{3}{4\pi\rho}\right)^{1/2}\,,$
(8)

where $\rho$ is the density, and $c_{s}$ is the sound speed defined as:

$c_{s}=\left(\frac{k_{B}\,T}{\mu_{H_{2}}\,m_{H}}\right)^{1/2}\,,$ (9)

where $\mu_{H_{2}}$ is the mean molecular weight per hydrogen molecule (equal
to 2.8; see also §4.2), and $m_{H}$ is the mass of a Hydrogen atom. The length
scale for a fragment to be bound, under the same assumptions, is the Jeans
radius

$R_{J}=\left(\frac{15\,c_{s}^{2}}{8\pi\,G\rho}\right)^{1/2}\,,$ (10)

where $R_{J}$ is the radius of a sphere of mass $M_{J}$ and the Jeans length
($\lambda_{J}$) is $\lambda_{J}\approx 2R_{J}$. Using the Hi-GAL clump
properties (1165 M⊙ and R$=$0.175 pc when scaled to D$=$4 kpc, T$=$17.9 K; §1,
Elia et al., 2017), we calculate the thermal Jeans radius and thermal Jeans
mass of the G19.01$-$0.03 clump to be $R_{J}\approx 0.012$ pc and
$M_{J}\approx 0.36$ M⊙ respectively. Using the ATLASGAL clump properties (926
M⊙ and R$=$0.358 pc for D$=$4.0kpc, T$=$19.5 K, §1, Schuller et al., 2009),
$R_{J}\approx 0.041$ pc and $M_{J}\approx 1.36$ M⊙. The linear resolution of
our ALMA data is $\sim$0.008 pc, and the $5\sigma$ detection limit (converted
to mass) is $0.1-0.3$ M⊙ at $24-15$ K, so we are sensitive to the relevant
fragmentation scales. With median masses of $1.0-2.0$ M⊙ at $24-15$ K,
respectively, MM2…MM5 are of order the thermal Jeans mass, consistent with the
results of Palau et al. (2015) from their study of mm fragments within the
inner 0.1 pc of $\sim$20 massive dense cores. The number of millimetre sources
detected in G19.01 is however notably lower than expected for thermal Jeans
fragmentation. For the 13% core formation efficiency (CFE) found by Palau et
al. (2015), $N_{\rm Jeans}=({M_{\rm clump}CFE)}/{M_{J}}$ $\sim$420 and
$\sim$90 for the Hi-GAL and ATLASGAL-based estimates, respectively. Strong
magnetic fields are expected to suppress fragmentation, and recent
observational work found a tentative correlation between the number of mm
fragments and the mass-to-flux ratio for a sample of 18 massive dense cores
with polarization data (Palau et al., 2021, and references therein).
Observations of polarized dust emission in G19.01$-$0.03 are needed to assess
whether a dynamically important magnetic field contributes to the low level of
observed fragmentation.

####  4.2.3 Virial analysis of MM4

The virial parameter ($\alpha_{vir}$) is a commonly used diagnostic in
determining the boundedness of structures. To qualify as a stellar progenitor,
a source must be in a gravitationally bound state. This is satisfied by
$\alpha_{vir}<2$ in the absence of pressure terms such as magnetic energy
density. Furthermore, a core is said to be in hydrostatic equilibrium when
$\alpha_{vir}\sim 1$, or to be gravitationally unstable when $\alpha_{vir}<1$.
The virial parameter is often expressed as follows:

$\alpha_{vir}=\frac{M_{vir}}{M_{\mathrm{core}}}\,,$ (11)

where $M_{vir}$ is the virial mass, and $M_{\mathrm{core}}$ is the observed
core mass. The virial mass for a spherical core may be calculated following
MacLaren et al. (1988):

$M_{vir}=3\left(\frac{5-2n}{3-n}\right)\frac{\sigma_{NT}^{2}R}{\mathrm{G}}\,,$
(12)

where $n$ is the index of the density profile of the source (where
$\rho\propto r^{-n}$), $R$ is the radius, and G is the gravitational constant.
The non-thermal velocity dispersion, $\sigma_{NT}$, may be calculated
following:

$\sigma_{NT}^{2}=\sigma_{\mathrm{obs}}^{2}-\sigma_{\mathrm{th}}^{2}=\sigma_{\mathrm{obs}}^{2}-\frac{k_{B}\,T_{\mathrm{dust}}}{\mu\,m_{H}}\,,$
(13)

where $\sigma_{\mathrm{obs}}$ and $\sigma_{\mathrm{th}}$ are the observed and
thermal velocity dispersions respectively (Bertoldi & McKee, 1992),
$T_{\mathrm{dust}}$ is the dust temperature (assumed to equal the gas
temperature), $\mu$ is the molecular weight of the molecular tracer, and
$m_{H}$ is the mass of a Hydrogen atom.

As discussed in §3.2.2 and §3.2.5, of the cold and dense gas tracers we detect
it is only the emission from DCN (with $\mu$ = 28) that appears
morphologically similar to the ALMA 1.05 mm dust continuum emission, with
$\geq 5\sigma$ detections towards MM1, MM2 and MM4. As we cannot safely
disentangle the DCN emission of MM2 from the surrounding envelope emission
(see Figure 4), we focus our virial analysis on MM4 as it appears isolated. As
we aim to test whether MM4 could be a prestellar core, we assume a flat
density profile where $n=0$ (e.g. Ward-Thompson et al., 1994). DCN exhibits
nuclear quadrupole hyperfine structure, with seven lines in the $J=4-3$
transition with velocity offsets from the main component between -1.6 to 2.1
km s-1 (Müller et al., 2001) with a mean velocity offset of 0.1 km s-1 (within
our 1 km s-1 spectral resolution). The effect of these components is most
important in the lower energy DCN transitions such as $J=1-0$ or $J=2-1$ (e.g.
Parise et al., 2009), with the strongest central component dominating in
transitions such as $J=4-3$. We fit the hyperfine structure lines on a pixel-
by-pixel basis across MM4 using the hfs fitting routine of the class software.
We set $\sigma_{\mathrm{obs}}$ to the mean velocity dispersion across the FWHM
size of MM4 of $0.58\pm 0.03$ km s-1. Taking the radius as half the geometric
mean of the deconvolved size of MM4 listed in Table 3 (equal to $\sim$2000
AU), the virial mass of MM4 is calculated to be $3.8\pm 0.4$ M⊙. With
$M_{\mathrm{core}}$ calculated from the 1.05 mm dust emission (1.84 M⊙ at 15
K, or 0.91 M⊙ at 24 K; see Table 6), this yields $\alpha_{vir}=2.0-4.1(\pm
0.4)$. This indicates that MM4 could be an unbound structure, perhaps a knot
in the red-shifted outflow lobe driven by MM1, likely to disperse back into
the ISM if internal pressures continue to dominate over gravity. On the other
hand, MM4 could be on the cusp of boundedness at the lower end of our
temperature range, possibly indicative of the early stages of a prestellar
core.

###  4.3 Chemistry of MM1 in context

Figure 12: Observed spectra towards the ALMA 1.05 mm continuum peak (black) of
MM1, overplotted with the combined highest-likelihood LTE synthetic spectra
(red) for all COM species listed in Table 7. Vertical grey lines mark the
transitions identified in §3.2.1, with dotted lines indicating transitions
mentioned in §4.3.1 that have overestimated brightness temperatures in the
synthetic spectrum.

As discussed in Sections 1 and 3.2.1, the detection of a line forest towards
MM1 with ALMA (see Figure 3) reveals a chemical richness that was lacking in
the SMA observations of Cyganowski et al. (2011a). We here estimate the
abundances of these newly detected molecular species towards MM1, to see where
MM1 sits in relation to other sources in the literature.

####  4.3.1 Estimation of likely column densities and abundances

In Sections 4.1.2 and 4.1.3, we produced LTE synthetic spectra to evaluate the
highest likelihood column density, temperature, centroid velocity and velocity
width of CH3OH towards MM1. However, the other detected species listed in
Table 4 do not have enough identified transitions for the constraint of four
free parameters, nor for the construction of a rotation diagram as in Section
4.1.1. We instead evaluate order of magnitude estimates for the column
densities of these species (CH3OCH3, 13CH3OH, CH3CH2OH, CH3OCHO, OCS, H2CO,
CH3CHO, H2CS, CH3CH2CN, NH2CHO, HC3N and 13CS) at the ALMA 1.05 mm continuum
peak of MM1 by keeping only the column density as a free parameter (e.g.
Csengeri et al., 2019) in the emcee sampler with weedspy_mcmc. The synthetic
spectra for each molecular species are evaluated individually. As in Section
4.1.3, the source size is fixed to the geometric mean of the ${\sim}278$ GHz
synthesised beam (${\sim}0.5\arcsec$). The temperature, centroid velocity and
velocity width are set to the highest likelihood values evaluated at the ALMA
1.05 mm continuum peak from the CH3OH spectra (i.e. 162.7 K, 59.7 km s-1 and
6.4 km s-1 respectively, see Figure 10). As discussed in Paper I, for
molecular lines associated with MM1 (i.e. not outflow-tracing) the extent of
the line emission is generally consistent across species and with the extent
of the 1.05 mm continuum emission (e.g. Fig. 3 of Williams et al., 2022). For
molecular species with transitions in multiple spectral windows, the synthetic
spectra were evaluated in the spectral window with the most transitions. The
resulting highest likelihood column density was then used to generate
synthetic spectra for the other spectral windows.

The highest likelihood column density of each species is listed in Table 7,
and Figure 12 shows the combined LTE synthetic spectrum for the two wide-band
spectral windows for all identified molecular species. On the whole the
combined model reproduces the observed emission well. Some molecular lines
appear in the synthetic spectrum that were not listed in Table 4 (or labelled
in Figure 3 or Figure 12), which is attributed to overlapping lines from the
same and/or multiple species making their unequivocal identification in
Section 3.2.1 difficult. Some of the identified transitions have their
emission over-estimated. These include the CH3OH($6_{1,5}-5_{1,4}$) and
CH3OH($10_{1,10}-9_{0,9}$) lines (with $\nu_{\mathrm{rest}}=292.67291$ and
$292.51744$ GHz respectively), the former being of an outflow-tracing nature
(see Figure 4) with an apparently low integrated intensity in the rotation
diagram of Figure 9(a), and the latter being the only transition identified in
the tuning to be in the first torsionally excited state. The heavily blended
H2CO($4_{2,3}-3_{2,2}$) and CH3OH($15_{1,0}-14_{2,0}$) lines (with
$\nu_{\mathrm{rest}}=291.23777$ and $291.24057$ GHz respectively) also have
their intensity over-estimated, which we attribute to the line blending and
the limitations of linearly combining independent models of different species.
The H2CO($4_{2,2}-3_{2,1}$) line ($\nu_{\mathrm{rest}}=291.94807$ GHz) also
has its intensity over-estimated, attributed to this line also being heavily
blended. Finally, the CH3OCH3($16_{1,16}-15_{0,15}$) line
($\nu_{\mathrm{rest}}=292.41225$ GHz) is over-estimated despite all other
CH3OCH3 lines in the tuning having their intensities well reproduced, and it
being shown in Paper I to exhibit compact emission consistent with the other
identified lines.

Fractional abundances of each species with respect to CH3OH and H2 are
evaluated from the ratio of their respective column densities. We take the
CH3OH column density at the ALMA 1.05 mm continuum peak from the synthetic
spectra in Section 4.1.3, and the H2 column density found in Section 4.2 from
the ALMA 1.05 mm continuum. As listed in Table 7, we report fractional
abundances of the identified species ranging from $10^{-1}-10^{-4}$ with
respect to CH3OH, and $10^{-6}-10^{-9}$ with respect to H2.

Table 7: Highest likelihood column densities of molecular species identified
towards the ALMA 1.05 mm continuum peak of MM1, and their abundances with
respect to CH3OH and H2, arranged by decreasing column density.

Molecule | N(X) | Fractional abundance
---|---|---
(X) | (cm-2) | N(X) / N(CH3OH) | N(X) / N(H2)
CH3OH | $2.2\times 10^{18}$ | $10^{0}$ | $3.3\times 10^{-6}$
CH3OCH3 | $4.3\times 10^{17}$ | $2.0\times 10^{-1}$ | $6.5\times 10^{-7}$
13CH3OH | $3.9\times 10^{17}$ | $1.8\times 10^{-1}$ | $5.9\times 10^{-7}$
CH3OCHO | $1.3\times 10^{17}$ | $5.9\times 10^{-2}$ | $2.0\times 10^{-7}$
CH3CH2OH | $6.0\times 10^{16}$ | $2.7\times 10^{-2}$ | $9.1\times 10^{-8}$
OCS | $4.4\times 10^{16}$ | $2.0\times 10^{-2}$ | $6.7\times 10^{-8}$
H2CO | $3.0\times 10^{16}$ | $1.4\times 10^{-2}$ | $4.6\times 10^{-8}$
CH3CHO | $1.7\times 10^{16}$ | $7.7\times 10^{-3}$ | $2.6\times 10^{-8}$
H2CS | $1.6\times 10^{16}$ | $7.2\times 10^{-3}$ | $2.4\times 10^{-8}$
CH3CH2CN | $1.0\times 10^{16}$ | $4.6\times 10^{-3}$ | $1.5\times 10^{-8}$
NH2CHO | $5.0\times 10^{15}$ | $2.3\times 10^{-3}$ | $7.6\times 10^{-9}$
HC3N | $4.4\times 10^{15}$ | $2.0\times 10^{-3}$ | $6.6\times 10^{-9}$
13CS | $1.5\times 10^{15}$ | $6.8\times 10^{-4}$ | $2.3\times 10^{-9}$

Column 1: Molecule name. Col. 2: Column density of each identified molecule at
the ALMA 1.05 mm continuum peak of MM1 (Table 3). Col. 3: Abundance of each
molecule with respect to the 12CH3OH column density at the ALMA 1.05 mm
continuum peak of MM1 ($2.2\times 10^{18}$ cm-2; Section 4.1.3). Col. 4: As
for column 3, but for H2 ($6.6\times 10^{23}$ cm-2; Table 6).

####  4.3.2 Comparison to other YSOs observed at high resolution

To put the methanol column density and these fractional abundances into
context, we compare them to values reported in the literature. Collated in
Figure 13 are ALMA observations for both low-mass YSOs and high-mass hot cores
that have reported observations of CH3OH, CH3OCH3 and CH3CHO, and that have
rotational temperatures and CH3OH column densities derived from the main CH3OH
isotopologue. Our MM1 data are overplotted for comparison. Compared to the
other high-mass YSOs shown in Figure 13, the CH3OH column density, temperature
and CH3CHO fractional abundance of MM1 appear consistent, whilst MM1 may be
one of the most abundant sources in CH3OCH3 (Figure 13b).

It is also interesting to consider the low- and high-mass sources in Figure 13
in relation to each other. Of the high-mass sources shown, G10.6$-$0.4 hot
core 1 (Law et al., 2021) has the lowest CH3OCH3 and CH3CHO fractional
abundances (Figure 13b, c), a result of its comparatively low CH3OCH3 and
CH3CHO column densities. Of the low-mass sources, IRAS 16293-2422 A (Manigand
et al., 2020) has the lowest CH3CHO fractional abundance (Figure 13c), and is
noted to have more extended CH3CHO emission morphology compared to the other
COMs, whilst SVS13$-$A VLA4B (Bianchi et al., 2022) exhibits the lowest
CH3OCH3 abundance (Figure 13b). However despite these outliers, the general
parameter spaces occupied by the low- and high-mass sources in Figure 13
appear remarkably similar to each other. Regarding the CH3OH column density
and temperature in particular, single-dish observations in the literature have
previously reported a distinct separation in these parameters across mass
scales. Öberg et al. (2014) for instance with their literature sample show low
CH3OH column density and temperature towards low-mass COM-emitting sources,
and high CH3OH column density and temperature towards high-mass COM-emitting
sources. With beam dilution a potential contributing factor to this, Öberg et
al. (2014) also collate a handful of resolved SMA observations, noting the
distribution to be more continuous rather than separated across mass when
including resolved observations. With the more recent ALMA observations
collated here, this behaviour is even more pronounced, with the chemistry of
low- and high-mass sources appearing similar in this set of interferometric
observations. A similar trend has recently been noted by Chen et al. (2023),
who compare the abundance ratios of Oxygen-bearing COMs towards their ALMA
sample of 14 high-mass sources with interferometric observations of 5 low-mass
sources from the literature. The Chen et al. (2023) sample includes
G19.01–0.03 MM1, and their values for the CH3OCH3, CH3OCHO and CH3CHO column
densities towards MM1 – derived using LTE synthetic spectra – are within a
factor of $2-4$ of the order of magnitude estimates in Table 7. We note that,
for consistency, we do not include datapoints from Chen et al. (2023) in
Figure 13 because they derive CH3OH column densities from modelling of the
CH${}_{3}^{18}$OH isotopologue and an assumed 16O$/^{18}$O isotope ratio, and
do not derive T${}_{\rm CH_{3}OH}$ values for most (13/14) of the sources in
their sample.

It is worth emphasising that comparisons of column densities are subject to a
number of varying effects including beam sizes, source distance and beam
dilution. The sources collated here, despite all being observed with ALMA,
appear to trace different spatial scales across mass. The low-mass data probe
physical scales of $13-140$ au (physical scale of the synthesised beam; median
and standard deviation of 60 and 40 au respectively), a factor of up to two
orders of magnitude smaller than the $400-2600$ au spatial scales (physical
scale of the synthesised beam; median and standard deviation of 990 and 870 au
respectively) probed by the high-mass data. It is suggested by van Gelder et
al. (2022) that methanol production could be less efficient towards MYSOs due
to their warmer prestellar phase and/or shorter prestellar lifetimes. Nazari
et al. (2023) model the methanol emission towards different configuration
high-mass sources. Comparing their results to a similar study of theirs
towards low-mass sources (Nazari et al., 2022), they find that optically thick
millimetre dust, as well as disc shadowing causing a reduction in environment
temperature, are more effective in lowering the methanol emission towards low-
mass YSOs than they are towards MYSOs with high luminosity ($10^{4}-10^{5}$
L⊙). It is therefore further suggested by Nazari et al. (2023) that other
factors such as the presence of Hii regions, could contribute to the lowering
of methanol emission towards some MYSOs. It is possible a combination of these
factors could be contributing to the methanol column densities, temperature
and fractional abundances appearing so similar across mass scales in Figure
13. It may also be explained by considering the conditions under which these
COMs are formed. Recent works have suggested that the chemical similarity
between various sources may be explained by the COMs being formed under
similar conditions (e.g. Quénard et al., 2018). Across mass scales this would
correspond to the ices of the pre-stellar phase (e.g. Coletta et al., 2020;
Nazari et al., 2022; Chen et al., 2023).

Figure 13: (a) CH3OH column density, (b) abundance of CH3OCH3 with respect to
CH3OH, and (c) abundance of CH3CHO with respect to CH3OH, all plotted against
CH3OH rotational temperature, of both low-mass (green) and high-mass (blue)
sources from the literature (inspired by Figs.7&8 of Öberg et al., 2014). Our
G19.01 data is plotted in yellow. All sources were observed with ALMA: the
synthesised beams of the collated low- and high-mass literature observations
range from $13-140$ au and $400-2600$ au respectively. High-mass sources were
collated from Csengeri et al. (2019), Molet et al. (2019), Law et al. (2021)
and Baek et al. (2022), and low-mass sources were collated from Lee et al.
(2019), Manigand et al. (2020), Bianchi et al. (2022), Chahine et al. (2022)
and Hsu et al. (2022). The G19.01 errors bars are smaller than the marker
size, and not all literature sources have error bars.

In addition to CH3OCH3 and CH3CHO, Csengeri et al. (2019) also report the
fractional abundances of CH3OCHO, CH3CH2OH and CH3CH2CN with respect to CH3OH
towards three positions around the high-mass hot core G328.2551–0.5321. In
every instance, the corresponding molecular abundance reported here towards
MM1 is of the same order as those towards G328.2551–0.5321. Cyganowski et al.
(2011a) find a CH3CN column density towards MM1 of $1.7\times 10^{16}$ cm-2
with their SMA 1.3 mm observations (with $32\times$ larger beam area than our
ALMA data). They also find that the CH3CN spectrum is best-fit by a model with
a CH3CN-emitting region of size 0.6″(equivalent to 2400 AU at D $=$ 4 kpc),
intermediate between our ALMA beam (Table 1) and the size of the MM1 1.05 mm
dust continuum from Table 3. Nazari et al. (2022) present CH3CN column
densities towards 37 high-luminosity, potentially high-mass, protostars; 13
have column densities of a similar order to MM1 (Cyganowski et al., 2011a),
whilst Law et al. (2021) reports a CH3CN column density for G10.6 HC 1 an
order of magnitude larger than MM1.

The comparatively weak, sparse COM emission towards MM1 in the SMA
observations of Cyganowski et al. (2011a) led to the suggestion that MM1 was a
relatively young source. Now, with a rich line forest detected with ALMA, and
fractional abundances of COMs similar to other sources in the literature, it
is clear that MM1 is not especially chemically young and in fact appears to be
a typical hot core.

####  4.3.3 ${}^{12}\mathrm{C}/^{13}\mathrm{C}$ isotope ratio

With the identification of both the 12C and 13C isotopologues of CH3OH, we can
estimate the ${}^{12}\mathrm{C}/^{13}\mathrm{C}$ isotope ratio from the ratio
of their column densities (see Table 7). This is valid if the molecular
transitions are optically thin, described by the same rotational temperature,
and exhibit similar spatial emission extents (e.g. Wirström et al., 2011).
Towards the ALMA 1.05 mm continuum peak of MM1, we estimate that
${}^{12}\mathrm{C}/^{13}\mathrm{C}\simeq 6$. This is a factor of $\sim$7 lower
than expected from the relations of Wilson & Rood (1994) and Milam et al.
(2005) for galactic molecular clouds (where
${}^{12}\mathrm{C}/^{13}\mathrm{C}\simeq 39-46$ at the MM1 galactocentric
distance of 4.4 kpc).

Though low-J transitions of CH3OH are commonly optically thick (e.g. Ginsburg
et al., 2017), 5 out of the 6 lines included in our LTE line modelling (see
§4.1.3) are relatively high-J transitions (see Table 4). The estimated
opacities of the 6 lines are $\tau~{}{\sim}~{}0.1$ to 0.6 based on the class
modelling (with an average of 0.3; see §4.1.3), and $\tau~{}{\sim}~{}0.4$ to
2.1 from the rotation diagram analysis (with an average of 1.0; see §4.1.1).
The opacity of the 13CH3OH($3_{2,2}-4_{1,3}$) line reported by class is
$\tau=1.0$. It is worth noting that the LTE modelling reproduces the observed
emission of the highest-J and lowest-$\tau$ 12CH3OH line very well (i.e.
CH3OH($23_{4,19}-22_{5,18}$) with $\tau=0.1$ and $0.4$ from class and the
rotation diagram respectively), and that the spatial extent of emission from
this line is consistent with that of 13CH3OH($3_{2,2}-4_{1,3}$) (see Fig.3 of
Paper I). To test for dependence on the assumed source size, we reran our
weedspy_mcmc analysis at the 1.05 mm continuum peak for the case of unresolved
emission (source size 0$\aas@@fstack{\prime\prime}$2, §4.1.1) and re-estimated
the 13CH3OH column density as described in Section 4.3.1 using the resulting
parameters. We emphasise that the best-fitting model with this smaller source
size is a poorer representation of the observed 12CH3OH emission than the
model described in Sections 4.1.3-4.1.4, and we use it only to check the
effect of the assumed source size on the ${}^{12}\mathrm{C}/^{13}\mathrm{C}$
ratio. From the models with a 0$\aas@@fstack{\prime\prime}$2 source size, we
derive a ${}^{12}\mathrm{C}/^{13}\mathrm{C}$ ratio of $\sim$7.5, indicating
that the result of low ${}^{12}\mathrm{C}/^{13}\mathrm{C}$ is robust to the
assumption of resolved or unresolved emission.

There are suggestions in the literature that
${}^{12}\mathrm{C}/^{13}\mathrm{C}$ may be generally lower on the smaller
scales of high-mass hot cores (e.g. IRAS 20126+4104, G31.41$-$0.31, AFGL 4176,
W3 IRS4 and G10.6$-$0.4; Palau et al., 2017; Beltrán et al., 2018; Bøgelund et
al., 2019; Mottram et al., 2020; Law et al., 2021, respectively) and low-mass
hot corinos (Hsu et al., 2022) than on the larger scale of molecular clouds.
Proposed chemical explanations include cold temperature enhancement of 13C in
H2CO formation on dust grains (Mottram et al., 2020) and the destruction of
HC3N in chemical reactions (Bøgelund et al., 2019). The possible contribution
of optical depth effects has also been considered (Beltrán et al., 2018;
Bøgelund et al., 2019; Law et al., 2021). Beltrán et al. (2018) and Bøgelund
et al. (2019) conduct synthetic analyses of CH3CN and HC3N emission
respectively (taking similar approaches to that detailed here towards G19.01),
and comment that optically thick emission could artificially lower their
observed ${}^{12}\mathrm{C}/^{13}\mathrm{C}$ ratios. Law et al. (2021)
conclude that optical depth effects are unlikely to drive their results as
$\tau~{}{\sim}~{}0.1-0.4$ for their CH3OH lines (derived using a rotation
diagram analysis), but note they cannot definitively discount optical depth
effects in unresolved structures contributing to their results.

Similarly to Law et al. (2021), our derived line opacities are not suggestive
of optical depth being the primary explanation for the low
${}^{12}\mathrm{C}/^{13}\mathrm{C}$ ratio observed in G19.01, but we cannot
rule out optical depth effects contributing to this result. We emphasise that
our observations provide only a tentative addition to the emerging picture of
low ${}^{12}\mathrm{C}/^{13}\mathrm{C}$ isotope ratios towards hot cores.
Observations of additional and optically thin transitions of the 12CH3OH and
13CH3OH isotopologues, as well as the isotopologues of other molecular
species, would be required for confirmation.

##  5 Conclusions

In this paper (Paper II), we have presented a study of the physical properties
and chemistry of the high-mass (proto)star G19.01–0.03 MM1 and its
environment, using sub-arcsec-resolution ALMA 1.05 mm and VLA 1.21 cm data.
Our main findings are as follows:

(i) A (sub)millimetre forest of molecular line emission is observed towards
MM1. We analyse 47 lines from 11 different species (43 of which were
identified in our ALMA 1.875 GHz wide-bands alone), including COMs such as
CH3OCHO, CH3CH2CN, CH3CH2OH, CH3OCH3, NH2CHO and CH3OH.

(ii) A Bayesian analysis (using our publicly available weedspy_mcmc scripts)
of CH3OH LTE synthetic spectra towards the MM1 dust peak returns highest
likelihood column density and temperature of $(2.22\pm 0.01)\times 10^{18}$
cm-2 and $162.7\begin{subarray}{c}+0.3\\\ -0.5\end{subarray}$ K respectively.
These are consistent with high-mass hot core sources from the literature, and
with values found from our opacity-corrected rotation diagram analysis of
$(2.0\pm 0.4)\times 10^{18}$ cm-2 and $166\pm 9$ K respectively.

(iii) The peak CH3OH temperature ($165.5\pm 0.6$ K) is offset from the MM1
dust peak by 0.22$\arcsec\sim 880$ au. The morphology of the region of
elevated temperature is aligned with the VLA 5.01 cm continuum emission
presented in Paper I.

(iv) We report abundances of all identified molecular species (CH3OCH3,
13CH3OH, CH3CH2OH, CH3OCHO, OCS, H2CO, CH3CHO, H2CS, CH3CH2CN, NH2CHO, HC3N
and 13CS) of $10^{-1}-10^{-4}$ and $10^{-6}-10^{-9}$ with respect to CH3OH and
H2 respectively, consistent with high-mass hot core sources from the
literature.

(v) The known bipolar molecular outflow driven by MM1 is traced by thermal
CH3OH and H2CO emission and by newly-identified NH3(3,3) and 278.3 GHz Class I
CH3OH maser candidates, strengthening the connection of these types of masers
with outflows driven by MYSOs. We identify a total of 50 $>$5$\sigma$ NH3(3,3)
emission spots across 8 compact emission groups, which are in the outer lobes
of the outflow and spatially and kinematically coincident with 44 GHz Class I
CH3OH masers. Candidate 25 GHz CH3OH 5(2,3)-5(1,4) maser emission is also
detected towards the outflow, offset from MM1 by 6$\arcsec$ $\sim$ 24,000 au.

(vi) Four new millimetre continuum sources (MM2…MM5) are detected between
$0.03-0.12$ pc from MM1, but all are undetected in NH3($J$=$K$=1,2,3,5,6,7)
and 25 GHz CH3OH emission. Two of these sources, MM2 and MM4, are associated
with $>5\sigma$ DCN(4–3) emission, a tracer of cold and dense gas.

(vii) The median mass of MM2..MM5 (assuming isothermal dust emission) is
1.0$-$2.0 M⊙ for a temperature range of $24-15$ K, and their median radius is
$\sim$ 2000 au. Since none of these sources are associated with any masers or
have any sign of outflows, they are candidate low-mass pre-stellar companions
to MM1.

In all, our results place G19.01–0.03 MM1 as a typical high-mass hot core
source, with four low-mass, potentially pre-stellar, cores within a projected
linear distance of 0.12 pc. Both these findings are in contrast with the
results of previous, lower-resolution SMA observations (Cyganowski et al.,
2011a), which identified G19.01–0.03 MM1 as a candidate for isolated high-mass
star formation, potentially at a very early stage due to its lack of chemical
complexity. Our new results show that G19.01–0.03 MM1 is in fact not
especially chemically young, and that its physical and chemical properties are
typical of other high-mass hot core sources in the literature. Intriguingly,
our analysis of the ALMA CH3OH lines reveals a region of elevated CH3OH
temperature that is aligned with VLA 5.01 cm continuum emission attributed to
free-free emission, which is offset from the thermal dust emission peak (as
discussed in detail in Paper I). This positional coincidence is tentative
support for the possibility, suggested in Paper I, that the 5.01 cm continuum
might be associated with an unresolved high-mass binary companion to MM1. Our
ALMA Cycle 6 follow-up study, tuned to probe $\sim$0.09″$\sim$360 au scales,
will be a next step in testing this possibility with higher angular resolution
observations.

## Acknowledgements

We thank the anonymous referee for their constructive report and comments that
helped improve the quality of this paper. GMW and CJC thank Ian Bonnell for
helpful discussions. GMW acknowledges support from the UK’s Science and
Technology Facilities Council (STFC) under ST$/$W00125X$/$1\. CJC acknowledges
support from the UK’s STFC under ST$/$M001296$/$1\. PN acknowledges support by
grant 618.000.001 from the Dutch Research Council (NWO) and support by the
Danish National Research Foundation through the Center of Excellence
"InterCat" (Grant agreement no.: DNRF150). This research has made use of:
NASA’s Astrophysics Data System Bibliographic Services, gildas
(https://www.iram.fr/IRAMFR/GILDAS) and python packages astropy (Astropy
Collaboration et al., 2013), astrodendro (Rosolowsky et al., 2008), aplpy
(http://aplpy.github.com), cmocean (Thyng et al., 2016), corner (Foreman-
Mackey, 2016), emcee (Foreman-Mackey et al., 2013), matplotlib (Hunter, 2007),
numpy (Harris et al., 2020), pandas (McKinney, 2010), scipy (Virtanen et al.,
2020), scitkit-image (Van der Walt et al., 2014), and analysisUtils (Hunter et
al., 2023). This paper makes use of the following ALMA data:
ADS$/$JAO.ALMA$\\#$2013.1.00812.S. ALMA is a partnership of ESO (representing
its member states), NSF (USA) and NINS (Japan), together with NRC (Canada),
NSC and ASIAA (Taiwan), and KASI (Republic of Korea), in cooperation with the
Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and
NAOJ. The National Radio Astronomy Observatory is a facility of the National
Science Foundation operated under cooperative agreement by Associated
Universities, Inc. For the purposes of open access, the author has applied a
Creative Commons Attribution (CC BY) licence to any Author Accepted Manuscript
version arising.

## Data Availability

The ALMA 1.05 mm dust continuum, VLA 1.21 and 5.01 cm continuum images, and
the line cubes for the two broad (1.875 GHz bandwidth) ALMA spectral windows
are available at doi:10.5281/zenodo.8059531. These data underlie both this
article and Williams et al. (2022). The scripts used to conduct the Bayesian
analysis of LTE synthetic spectra are packaged together in a reusable form we
call weedspy_mcmc, and are made freely available at https://github.com/gwen-
williams/WeedsPy_MCMC.

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