11institutetext: 1 Institute of Astronomy, Faculty of Physics, Astronomy and
Informatics, Nicolaus Copernicus University, Grudziadzka 5, 87-100 Toruń,
Poland
2 Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121, Bonn,
Germany
3 National Centre for Nuclear Research, ul. Pasteura 7, 02-093, Warszawa,
Poland
4 Exoplanets and Stellar Astrophysics Laboratory, NASA Goddard Space Flight
Center, Greenbelt, MD 20771, USA
5 Center for Research and Exploration in Space Science and Technology, NASA
Goddard Space Flight Center, Greenbelt, MD 20771
6 Department of Astronomy, University of Maryland, College Park, MD 20742, USA
7 Astronomical Observatory of the Jagiellonian University, Orla 171, 30-244,
Kraków, Poland
8 Deutsches SOFIA Institut, University of Stuttgart, Pfaffenwaldring 29,
70569, Stuttgart
9 SOFIA/USRA, NASA Ames Research Center, P.O. Box 1, MS 232-12, Moffett Field,
CA 94035, USA
10 Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD
21218, USA
11 Niels Bohr Institute, Centre for Star and Planet Formation, University of
Copenhagen, Øster Voldgade 5-7, 1350 Copenhagen, Denmark

# Far-infrared line emission from the outer Galaxy cluster Gy 3–7 with
SOFIA/FIFI-LS: Physical conditions and UV fields††thanks: Tables 2, 4, 5, 6,
7, 8, and 9 are available in electronic form at the CDS via anonymous ftp to
cdsarc.cds.unistra.fr (130.79.128.5) or via https://cdsarc.cds.unistra.fr/cgi-
bin/qcat?J/A+A/

Ngân Lê 11 A. Karska Corresponding author: Agata Karska
1122<EMAIL_ADDRESS>M. Figueira 3311 M. Sewiło 445566 A. Mirocha 77 Ch.
Fischer 88 M. Kaźmierczak-Barthel 88 R. Klein 99 M. Gawroński 11 M. Koprowski
11 K. Kowalczyk 11 W. J. Fischer 1010 K. M. Menten 22 F. Wyrowski 22 C. König
22 L. E. Kristensen 1111

(Received February 14, 2023; accepted April 10, 2023)

###### Abstract

Context. Far-infrared (FIR) line emission provides key information about the
gas cooling and heating due to shocks and UV radiation associated with the
early stages of star formation. Gas cooling via FIR lines might, however,
depend on metallicity.

Aims. We aim to quantify the FIR line emission and determine the spatial
distribution of the CO rotational temperature, ultraviolet (UV) radiation
field, and H2 number density toward the embedded cluster Gy 3–7 in the
CMa–$l224$ star-forming region, whose metallicity is expected to be
intermediate between that of the Large Magellanic Cloud and the Solar
neighborhood. By comparing the total luminosities of CO and [O i] toward Gy
3–7 with values found for low- and high-mass protostars extending over a broad
range of metallicities, we also aim to identify the possible effects of
metallicity on the FIR line cooling within our Galaxy.

Methods. We studied SOFIA/FIFI-LS spectra of Gy 3–7, covering several CO
transitions from $J=14-13$ to $31-30$, the OH doublet at 79 $\mu$m, the [O i]
63.2 and 145.5 $\mu$m, and the [C ii] 158 $\mu$m lines. The field of view
covers a $2^{\prime}\times 1^{\prime}$ region with a resolution of
$\sim$7$\arcsec$-18$\arcsec$.

Results. The spatial extent of CO high$-J$ ($J_{\mathrm{up}}\geq$14) emission
resembles that of the elongated 160 $\mu$m continuum emission detected with
Herschel, but its peaks are offset from the positions of the dense cores. The
[O i] lines at 63.2 $\mu$m and 145.5 $\mu$m follow a similar pattern, but
their peaks are found closer to the positions of the cores. The CO transitions
from $J=14-13$ to $J=16-15$ are detected throughout the cluster and show a
median rotational temperature of 170$\pm$30 K on Boltzmann diagrams.
Comparisons to other protostars observed with Herschel show a good agreement
with intermediate-mass sources in the inner Galaxy. Assuming an origin of the
[O i] and high$-J$ CO emission in UV-irradiated $C-$shocks, we obtained pre-
shock H2 number densities of 104-105 cm-3 and UV radiation field strengths of
0.1-10 Habing fields ($G_{\mathrm{0}}$).

Conclusions. Far-IR line observations reveal ongoing star formation in Gy 3–7,
dominated by intermediate-mass Class 0/I young stellar objects. The ratio of
molecular-to-atomic far-IR line emission shows a decreasing trend with
bolometric luminosities of the protostars. However, it does not indicate that
the low-metallicity has an impact on the line cooling in Gy 3–7.

###### Key Words.:

stars:formation – stars: protostars – ISM: jets and outflows – ISM: molecules

##  1 Introduction

During the earliest stages of star formation, the gravitational collapse of
dense cores is accompanied by an ejection of bipolar jets originating from the
resulting protostars, which may alter the physical conditions and chemistry of
their environment, even on clump scales (Arce et al., 2007; Frank et al.,
2014). Non-dissociative shock waves develop as the jets (and winds) interact
with the surrounding medium (Kaufman & Neufeld, 1996a; Flower & Pineau des
Forêts, 2012) and heat up the gas up to typically $\sim$300 K (Karska et al.,
2018; Yang et al., 2018). Additionally, ultraviolet (UV) photons contribute to
the gas heating and influence the chemical composition of the low- to high-
mass protostars’ envelopes (Bruderer et al., 2009; Visser et al., 2012).
Similarly to some pc-scale outflows, UV photons may operate over a significant
fraction of low-mass star-forming clumps and clusters (Mirocha et al., 2021).
The cooling of the gas, which in the case of embedded objects is dominated by
line emission in the far-infrared (FIR) and (sub)millimeter domains, provides
important constraints on the heating mechanisms and observations of these
cooling lines allow us to constrain gas temperatures, densities, and UV fields
(Goldsmith & Langer, 1978; Hollenbach & McKee, 1989).

Recent observations with the Herschel Space Observatory (Pilbratt et al.,
2010)111Herschel was an ESA space observatory with science instruments
provided by European-led Principal Investigator consortia and with important
participation from NASA. targeted the main gas cooling lines toward a
significant sample of protostars spanning a broad range of masses (van
Dishoeck et al., 2021). In particular, the Photodetector Array Camera and
Spectrometer (PACS; Poglitsch et al., 2010) provided detections of high$-J$ CO
($J_{\mathrm{up}}\geq$14), H2O, and OH lines, as well as forbidden transitions
of [O i] and [C ii], all of them being important diagnostic tools in molecular
clouds. Among the key findings with PACS are the following: (i) the presence
of ubiquitous ”warm” gas ($\sim$300 K) associated with low- to high-mass
protostars (Green et al., 2013; Manoj et al., 2013; Karska et al., 2013,
2014a; Matuszak et al., 2015); (ii) the detection of a plethora of high$-J$ CO
(up to $49-48$; $E_{\mathrm{u}}$ $\sim$6700 K) and H2O (up to $E_{\mathrm{u}}$
$\sim$1500 K) lines tracing the ‘” hot” gas component (Herczeg et al., 2012;
Goicoechea et al., 2012), with a median temperature of $\sim$720 K (Karska et
al., 2018); (iii) the identification of the origin of the FIR line emission in
UV-irradiated non-dissociative shocks extending along the outflows (Karska et
al., 2014b; Benz et al., 2016; Kristensen et al., 2017b; Karska et al., 2018);
and (iv) the recognition of the dominant role of CO and H2O in the gas cooling
budget of low-mass (LM) protostars (Karska et al., 2013, 2018) as well as CO
and [O i] as coolants of high-mass (HM) protostars (Karska et al., 2014a). The
above observations have provided a surprisingly uniform picture of the FIR
line emission from deeply-embedded protostars, but targeted only on relatively
nearby regions ($d$¡450 pc in case of LM objects van Dishoeck et al., 2021).

Figure 1: Digital Sky Survey (DSS) R (top left), 2MASS composite image using
the J, H, and Ks filters (top right), Spitzer/GLIMPSE360 4.5 $\mu$m (Sewiło et
al. 2019; bottom left), and Herschel/Hi-GAL composite image at 70, 160, and
250 $\mu$m (Elia et al. 2013; bottom right) of the Gy 3–7 cluster. Circles in
the top panels show the positions of YSO candidates from Tapia et al. (1997).
Yellow diamond and red cross in the bottom-left panel show the position of the
H2O maser (Urquhart et al., 2011) and the IRAS source at the south-west side
(IRAS 07069–1045), respectively. The blue $"\times"$ symbols in the bottom
right panel show the positions of dense cores as traced by the H2 column
density (Elia et al., 2013). Dense core in the west corresponds to IRAS
07069–1045.

Observations of the Large and Small Magellanic Clouds (LMC and SMC) with
Herschel have shed some light on the FIR line cooling from protostars in a
significantly different, low-metallicity environment ($Z$ of 0.2–0.5
$Z_{\odot}$; Russell & Dopita 1992). The Spectral and Photometric Imaging
Receiver (SPIRE, Griffin et al., 2010) and PACS provided detections of CO (up
to $J$=14-13), H2O, OH, and [O i], and [C ii] lines, facilitating comparisons
with Galactic sources (Oliveira et al., 2019). The CO emission toward 15
sources in the LMC showed two relatively cool gas components, with
temperatures of $\sim$40 and $\sim$120 K, consistent with Galactic
measurements using SPIRE (White et al., 2010; Jiménez-Donaire et al., 2017;
Yang et al., 2018). The line cooling budget of protostars in the Magellanic
Clouds is dominated by [O i] and [C ii] emission, with an increasing
contribution of CO as the metallicity increases from the SMC to LMC and to the
Galactic young stellar objects (YSOs, Karska et al., 2014a; Oliveira et al.,
2019). The low fraction of CO line cooling is interpreted as a metallicity
effect: the combined result of a reduced carbon abundance and higher grain
temperatures due to a lower shielding from UV photons (Oliveira et al., 2019).

The outer parts of our Galaxy provide an alternative site for testing the
impact of metallicity on the FIR gas properties of protostars. Due to the
negative-metallicity gradient, the abundances of dust and molecules decrease
in the outer Galaxy (Sodroski et al., 1997). The metallicity affects the gas
and dust cooling budget of molecular clouds and results in lower CO rotational
temperatures, $T\rm{{}_{rot}}$ (Roman-Duval et al., 2010). Despite the overall
decreasing trend of the mass surface density of molecular clouds in the outer
Galaxy (for a review, see Heyer & Dame, 2015), some star-forming regions show
a significant star-formation activity. For example, the CMa–$l224$ star-
forming region at a Galactocentric radius, $R_{\mathrm{GC}}$, of 9.1 kpc
consists of $\sim$290 Class I/II YSOs, as identified by Sewiło et al. (2019)
using data from GLIMPSE360: Completing the Spitzer Galactic Plane Survey (PI:
B. Whitney) and the Herschel infrared Galactic Plane Survey (Hi-GAL; Molinari
et al. 2010, see also Elia et al. 2013). The expected metallicity of this
region is $\sim$0.55-0.73 $Z_{\odot}$, depending on the adopted O/H
Galactocentric radial gradient (Balser et al., 2011; Fernández-Martín et al.,
2017; Esteban & García-Rojas, 2018).

Table 1: Catalog of lines observed with FIFI-LS toward Gy 3–7

Line | Transition | $E_{\mathrm{u}}/k_{\mathrm{B}}$ | $A_{\rm{u}}$ | $g_{\rm{u}}$ | $\lambda_{\mathrm{lab}}$ | Map size | Beam
---|---|---|---|---|---|---|---
|  | (K) | (s-1) |  | ($\mu$m) | ($\arcsec\times\arcsec$) | ($\arcsec$)
CO | $14-13$ | 580.5 | 2.7(-4) | 29 | 185.99 | 60$\times$60 | 18.3
CO | $16-15$ | 751.7 | 4.1(-4) | 33 | 162.81 | 60$\times$60 | 16.1
CO | $17-16$ | 845.6 | 4.8(-4) | 35 | 153.27 | 60$\times$60 | 15.2
CO | $22-21$ | 1397.4 | 1.0(-3) | 45 | 118.58 | 30$\times$30 | 11.9
CO | $30-29$ | 2564.9 | 2.3(-3) | 61 | 87.19 | 30$\times$30 | 8.6
COaa$a$CO 31-30 and OH line at 84.42 $\mu$m are blended. | 31-30 | 2735.3 | 2.5(-3) | 63 | 84.41 | 30$\times$30 | 8.3
OHbb$b$OH line at 163.12 $\mu$m is not detected. |  $\frac{1}{2}$,$\frac{1}{2}-\frac{3}{2}$,$\frac{1}{2}$ | 270.1 | 6.5(-2) | 4 | 163.12 | 60$\times$60 | 16.1
OHcc$c$The 79.18 $\mu$m is detected at the band edge, and its line fluxes cannot be securely measured due to the lack of a baseline. |  $\frac{1}{2}$,$\frac{1}{2}-\frac{3}{2}$,$\frac{3}{2}$ | 181.7 | 2.9(-2) | 3 | 79.18 | 30$\times$30 | 7.9
OHaa$a$CO 31-30 and OH line at 84.42 $\mu$m are blended. |  $\frac{7}{2}$,$\frac{3}{2}-\frac{5}{2}$,$\frac{3}{2}$ | 290.5 | 5.1(-1) | 9 | 84.42 | 30$\times$30 | 8.3
[O i] |  3P0–3P2 | 227.7 | 8.9(-5) | 3 | 63.18 | 60$\times$60 | 7.1
[O i] |  3P1–3P2 | 326.6 | 1.8(-5) | 1 | 145.52 | 60$\times$60 | 14.4
[C ii] |  2P3/2–2P1/2 | 91.2 | 2.4(-6) | 5 | 157.74 | 90$\times$90 | 15.6
222Molecular data adopted from the Leiden Atomic and Molecular Database
(LAMDA, Schöier et al. 2005) and the JPL database (Pickett et al., 1998).

Figure 2: Distribution of dust continuum emission toward Gy 3–7. Left:
Continuum map of Gy 3–7 at 70 $\mu$m observed with Herschel/PACS. Red $"+"$
and blue $"\times"$ symbols refer to the 70 $\mu$m continuum peaks, adopted as
the positions for the dense cores A and B in the subsequent analysis. Circles
show the positions of HIGALBM224.6079–1.0065 and HIGALBM224.6128–1.0013 cores
from the Herschel/Hi-GAL catalog (Elia et al., 2021). Right: Map of the H2
column density ($N_{\mathrm{H_{2}}}$). Black circles with the beam size of
20$\arcsec$ indicate the extract regions of the SOFIA FIFI-LS spectra toward
the two dense cores A and B (see more in Section 3.1). White contours in each
map show the continuum at 70 $\mu$m, with contour levels at 5, 10, 40, 80,
400, and 800 mJy/arcsec2.

Gy 3–7 is a deeply-embedded cluster with exceptionally bright FIR continuum
emission, located at the second-most massive filament in the CMa–$l224$ region
(Sewiło et al., 2019) at a distance of $\sim$1 kpc (e.g., Lombardi et al.,
2011). It is associated with IRAS 07069–1045, which was recognized early on as
a star-forming region driving a CO outflow, with a bolometric luminosity
($L_{\rm{bol}}$) of 980 $L_{\rm{\odot}}$ (assuming a distance of 1.4 kpc;
Wouterloot & Brand, 1989). The source was considered as a candidate HM star-
forming region based on its position in the IRAS color-color diagrams and the
presence of dense gas traced by CS $2-1$ (Wood & Churchwell, 1989; Bronfman et
al., 1996). However, several attempts have failed to detect the CH3OH maser
emission, which is a common signature of HM protostars (Menten et al., 1992;
Szymczak et al., 2000). Instead, H2O maser and thermal NH3 emission was
detected as part of the Red MSX Survey (RMS; Urquhart et al., 2011, see Figure
1).

Figure 3: SOFIA FIFI-LS continuum-subtracted spectra toward the two dense
cores in Gy 3–7: HIGALBM224.6079–1.0065 (source “A”) and
HIGALBM224.6128–1.0013 (source “B”). The emission is extracted within a beam
size of 20$\arcsec$ indicated by black circles in the right panel of Figure 2.
Vertical lines show the laboratory wavelengths of the detected lines. Spectra
of the [O i] line at 63 $\mu$m are multiplied by a factor of 0.2 and those of
the CO $14-13$ line by a factor of 0.6 to better illustrate the line
detections.

Recent FIR observations with Herschel/Hi-GAL spatially resolved two dense
cores in Gy 3–7 with $L_{\mathrm{bol}}$ of 75.9 and 324.2 $L_{\rm{\odot}}$,
the latter corresponding to IRAS 07069–1045 (Elia et al., 2013, 2021, see
Figure 1). Near-IR observations revealed an extended H2 emission, which may
arise from the jets from protostars (Navarete et al., 2015). Gy 3–7 contains
several more evolved YSOs, with spectral types ranging from B1 to B5
(Gyulbudaghian, 2012; Tapia et al., 1997) and it is cataloged as a young
stellar cluster (Soares & Bica, 2002, 2003; Bica et al., 2003).

In this paper, we investigate the FIR line emission toward Gy 3–7 obtained
using the Stratospheric Observatory for Infrared Astronomy (SOFIA)
observations with the Field-Imaging Far-Infrared Line Spectrometer (FIFI-LS;
Klein et al., 2014; Fischer et al., 2018). We also consider the gas densities
and UV radiation fields in Gy 3–7, as well as the stellar content of the
cluster. We examine whether the rotational temperatures and the ratio of CO
and [O i] line luminosities in Gy 3–7 are consistent with the picture of star
formation in the inner Galaxy and/or the low-metallicity environment of the
Magellanic Clouds. Finally, we also present the results of our search for
water masers with the Toruń 32-m radio telescope (RT4).

The paper is organized as follows. Section 2 describes the observations with
SOFIA and RT4. In Section 3, we present line and continuum maps, as well as
the line profiles at selected positions of interest. Section 4 shows the
analysis of the results and Section 5 provides the discussion of the results
in the context of previous studies. We provide our summary and conclusions in
Section 6.

##  2 Observations

###  2.1 SOFIA FIFI-LS

SOFIA/FIFI-LS observations were collected in November 2019 as part of the
SOFIA Cycle 7 (Project ID $07_{-}0157$, PI: M. Kaźmierczak-Barthel). FIFI-LS
is an integral field unit consisting of two grating spectrometers with a
spectral coverage ranging from 51 to 120 $\mu$m (blue) and from 115 to 200
$\mu$m (red), facilitating simultaneous observations of selected wavelength
intervals (0.3-0.9 $\mu$m) in both channels (SOFIA Observer’s Handbook for
Cycle 10333https://www-sofia.atlassian.net/wiki/spaces/OHFC1/overview). The
spectral resolution, $R$, ranges from $\sim$500 to 2000 and increases with
wavelength for a given grating order. The corresponding velocity resolution of
$\sim$150 to 600 km s-1 provides unresolved spectral profiles of all the FIR
lines, including H2O (Kristensen et al., 2012; Mottram et al., 2017), CO
(Kristensen et al., 2017b), and [O i] (Kristensen et al., 2017a; Yang et al.,
2022a).

Figure 4: FIFI-LS contour maps of the [O i] lines at 63.2 and 145.5 $\mu$m,
the [C ii] line at 157.7 $\mu$m, the CO lines with $J_{\mathrm{up}}$ = 14, 16,
17 at 186, 163, and 153 $\mu$m respectively (white contours) on top of the
continuum emission at 160 $\mu$m (at 70 $\mu$m for the [O i] line at 63.2
$\mu$m) from Herschel/PACS. The white contours show line emission at 25%, 50%,
75%, and 95% of the corresponding line emission peak. The $"+"$ and $"\times"$
signs show the positions of the dense cores A and B.

Figure 5: FIFI-LS contour maps of the [O i] lines at 63.2 and 145.5 $\mu$m,
the [C ii] line at 157.7 $\mu$m, and the CO $14-13$ line at 186 $\mu$m (white
contours) on top of the continuum emission at 4.5 $\mu$m from Spitzer/IRAC.
The contours show line emission at 25%, 50%, 75%, and 95% of the corresponding
line emission peak. The ”$+$” and ”$\times$” signs show the positions of the
dense cores A and B, respectively. The orange squares show the positions of
two YSO candidates with envelopes, and the yellow circles show the positions
of the remaining YSOs (Section 4.5).

Figure 6: Integrated intensity maps of selected pairs of FIR lines from FIFI-
LS with the top line shown in colors, and the bottom line in white contours.
The contour levels are 25%, 50%, 75%, and 95% of the corresponding line
emission peak. The ”+” and ”x” signs show the positions of the dense cores A
and B, respectively. The spaxel size in the emission map of the [O i] line at
63 $\mu$m is $6\arcsec\times 6\arcsec$ (blue channel) and in the map of other
lines (CO $14-13$, [O i] at 145 $\mu$m, and [C ii] at 157.7 $\mu$m) is
$12\arcsec\times 12\arcsec$ (red channel). Gray circles show the beam size for
each color map.

The FIFI-LS detector is composed of 5$\times$5 spatial pixels (hereafter
spaxels) with the centers offset by 10$\arcsec$, similar to the PACS
spectrometer on Herschel. The spaxel size is $6\arcsec\times 6\arcsec$ in the
blue channel (field of view, FOV, of 30$\arcsec\times 30\arcsec$) and
$12\arcsec\times 12\arcsec$ in the red channel (FOV of $1\arcmin\times
1\arcmin$), providing an improvement over PACS by matching the actual
wavelength-dependent beam sizes. The FIR sky background was subtracted by
symmetric chopping around the telescope’s optical axis with a matched
telescope nod. The 300$\arcsec$ throw in the east-west direction was chosen to
avoid contaminated regions.

Table 1 shows the full catalog of lines targeted with FIFI-LS. The maps of the
[O i] line at 63.2 $\mu$m and the [C ii] line at 157.7 $\mu$m were obtained
simultaneously as $2\times 2$ mosaics, with a FOV of $60\arcsec\times
60\arcsec$ in the blue channel and $90\arcsec\times 90\arcsec$ in the red
channel, during the wall-clock time of 51 minutes. The maps of the remaining
lines were collected as a single FOV with dithering. The total observing time
was 239 minutes ($\sim$4 hours).

The data were reduced with the SOFIA FIFI-LS pipeline (Vacca et al., 2020),
which contains all the necessary calibrations and flat-field corrections. For
the telluric correction in the pipeline reduction, we used water vapor values
obtained with the method described in Fischer et al. (2021) and Iserlohe et
al. (2021). Then, the IDL-based software FLUXER
v.2.78444http://ciserlohe.de/fluxer/fluxer.html was used to produce the
continuum and line emission maps. The continuum was fitted as a 0th order
polynomial in spectral areas free of line emission and the spectral line was
fitted with a Gaussian. The selection of channels for the baseline
subtraction, the observed wavelength of the targeted line and its width were
obtained over the area with strong line detections, separately for each
species. Subsequently, these values were adopted for the entire datacube to
obtain the integrated line fluxes for the entire map. Further processing of
the maps was performed with Python.

The overall calibration accuracy can be assumed to be within 20%, including a
10% calibration accuracy of the instrument, and the additional uncertainty due
to telluric effect, which is assumed to be well under 10%. For the [O i] line
at 63.2 $\mu$m, located close to the water absorption line, the water vapor
overburden was determined between 3.5 and 3.7 $\mu$m. With an error of 10% on
this water vapor range, this results in a transmission at the line location of
78-83%.

###  2.2 Herschel/PACS

We used the Herschel/PACS 70 $\mu$m continuum map to verify the positions of
the dense cores in Gy 3–7 listed in the Herschel/Hi-GAL compact source
catalogue555https://vialactea.iaps.inaf.it/vialactea/public/HiGAL_360_clump_catalogue.tar.gz
(Elia et al., 2021). Figure 2 shows the 70 $\mu$m continuum map and the H2
column density, $N$(H2), map obtained using the $\tt ppmap$ tool with the
Herschel/Hi-GAL survey (Marsh et al., 2017). Overall, there is a good
agreement between the catalog positions and the peaks in the continuum at 70
$\mu$m and $N$(H2). In the subsequent analysis, we adopt the coordinates of
the 70 $\mu$m peaks as the dense core coordinates. We refer to the core at
(RA, Dec) = ($7^{h}09^{m}20\aas@@fstack{s}4,-10\degr 50\arcmin
28\aas@@fstack{\prime\prime}4$) as A and to that at
($7^{h}09^{m}21\aas@@fstack{s}9,-10\degr 50\arcmin 35\arcsec$) as B; they
correspond to the Herschel/Hi-GAL catalog sources HIGALBM224.6079–1.0065 and
HIGALBM224.6128–1.0013, respectively (Elia et al., 2021).

###  2.3 RT4

We conducted a survey at 22 GHz using the Toruń 32-m radio telescope (RT4)
toward the entire CMa–l224 region. The full-beam width at half maximum of the
antenna at 22 GHz is $\sim$106$\arcsec$, with a pointing error of
$\lesssim$12$\arcsec$ (Lew, 2018).

Two series of observations were performed from 2019 April to 2020 January and
from 2020 January to May 2020 in which a total of 205 positions were observed.
We used the correlator with 2$\times$4096 channels and 8 MHz bandwidth
operating in the frequency–switching mode, which provided a local standard of
rest velocity coverage from $-11.9$ to $41.9$ km s-1. The spectral resolution
of the observations is 0.03 km s-1. The observations were calibrated by the
chopper wheel method and corrected for the gain elevation effect. The system
temperature varied from $\sim$120 K during winter to $\sim$200 K in summer.
Overall, the $3\sigma$ detection limit was between $\sim$3.8 and $\sim$7.5 Jy
per 0.03 km s-1 channel. The telescope pointing was checked every $\sim$2 hrs
by observing a nearby bright, point source. We successfully detected a
variable water maser emission at (RA, Dec) =
($7^{h}09^{m}21\aas@@fstack{s}05,-10\degr 50\arcmin
05\aas@@fstack{\prime\prime}4$), in direct vicinity of Gy 3-7, and at (RA,
Dec) = ($7^{h}09^{m}22\aas@@fstack{s}66,-10\degr 30\arcmin
45\aas@@fstack{\prime\prime}6$), associated with IRAS 07069-1026 (see Appendix
A).

##  3 Results

Near-IR images of Gy 3–7 reveal an extended nebulosity associated with the two
Hi-GAL dense cores and several YSO candidates (see Figure 1). Spatially-
resolved FIR line emission data obtained with FIFI-LS allows us to study key
gas cooling lines at $\sim$10 000 au scales and identify regions where
processes responsible for the gas heating are at play.

###  3.1 Line detections

Figure 3 shows the FIR line emission toward the two Hi-GAL dense cores in Gy
3–7 (see Section 2.2; for the full list of targeted lines see Table 1). All
lines are spectrally-unresolved with FIFI-LS and can be represented by single
Gaussian profiles (Section 2.1).

The spectra show strong line emission in the [O i] lines at 63 $\mu$m and 145
$\mu$m, as well as in the high$-J$ CO lines: the CO $14-13$ line at 186.0
$\mu$m, the CO $16-15$ line at 162.8 $\mu$m, and the CO $17-16$ line at 152.3
$\mu$m. Core A shows also a detection of the CO $22-21$ line at 118.6 $\mu$m
($E_{\mathrm{u}}$ of $\sim$1400 K) and a tentative detection of CO $30-29$
line at 87.2 $\mu$m ($E_{\mathrm{u}}$ of $\sim$2600 K). The CO $31-30$ line at
84.41 $\mu$m is blended with the OH line at 84.42 $\mu$m, and due to the lack
of baseline covering the OH 84.6 $\mu$m line from the doublet, its emission
cannot be quantified. The OH doublet at $\sim$79.2 $\mu$m seems to be
tentatively detected toward both cores, but is severely affected by the rise
of the baseline on its left side and its flux cannot be properly measured
(Appendix B). The OH doublet at 163.12 and 163.18 $\mu$m, located next to the
CO 16-15 line, is not detected in neither of the two cores.

The [O i] line at 63.18 $\mu$m sits on the edge of a telluric water feature
and the low transmission at $\lambda\ga 63.24$ $\mu$m increases the noise on
the continuum. Since the transmission at the spectral line location is well
known and the S/N is very high, there is no relevant effect on the uncertainty
of the line flux. Possible self-absorption of this [O i] line cannot be
identified at the spectral resolution of FIFI-LS. The higher-resolution
spectra collected toward other YSOs using the German REceiver for Astronomy at
Terahertz Frequencies (GREAT; Risacher et al., 2018) show that self-absorption
could decrease the integrated emission of the line by a factor of 2–3 (Leurini
et al., 2015; Mookerjea et al., 2021).

Similarly, self-absorption in the [C ii] line at 157.7 $\mu$m might affect the
FIFI-LS spectra. Strong self-absorption unresolved by FIFI-LS may result in a
non-detection of this line toward core A. In contrast, the line is clearly
detected toward core-B. Decreases of the flux of [C ii] line due to self-
absorption as high as a factor of 20 have been estimated toward
photodissociation regions (Guevara et al., 2020).

In summary, FIFI-LS spectra provide detections of key FIR cooling lines toward
two dense cores in Gy 3–7: the high$-J$ CO, [O i], and [C ii]. Due to the
atmospheric absorption, the H2O emission could not be targeted; the OH
emission suffers from line blending and poor baselines. In Section 3.2, we
show the distribution of line emission in various species toward the entire
cluster.

###  3.2 Spatial extent of FIR line emission

The FIR range contains several important diagnostic lines, which provide
information about the physical conditions and processes that strongly
contribute to the gas cooling (Goldsmith & Langer, 1978; Kaufman & Neufeld,
1996b). The analysis of spatial extent of various FIR species pinpoints the
presence of shocks and/or UV radiation associated with star formation.

Figures 4 and 5 show the spatial extent of the FIR lines detected toward Gy
3–7 (Section 3.1). The line emission is compared to the FIR dust continuum
emission at 70 or 160 $\mu$m from Herschel/PACS (Figure 4) and the 4.5 $\mu$m
continuum tracing warmer dust from Spitzer/IRAC (Figure 5). Bright
rotationally excited H2 emission associated with shocks might also contribute
to the flux in the 4.5 $\mu$m IRAC band (Cyganowski et al., 2008, 2011).
Appendix B shows additional maps of Gy 3–7 for lines from various species.

The high$-J$ CO emission distribution is elongated in the same direction as
the IR continuum, but its peaks are offset from the continuum peaks at
wavelengths similar to those of the respective CO lines. A similar pattern of
emission is also seen for the [O i] 63 $\mu$m and 145 $\mu$m lines; yet, the
peak of the [O i] lines are almost co-spatial with the core positions, which
is not the case for the CO lines. Nevertheless, both the CO and the [O i]
extend beyond the core positions along the E–W direction. Such elongated
high$-J$ CO morphologies have been commonly interpreted as arising in shocked
outflows from LM and IM protostars (Goicoechea et al., 2012; He et al., 2012;
Kristensen et al., 2012; Karska et al., 2013; Matuszak et al., 2015; Green et
al., 2016; Tobin et al., 2016; Kristensen et al., 2017b). Similarly, extended
[O i] emission has been associated with embedded, atomic outflows (Karska et
al., 2013; Nisini et al., 2015), as recently confirmed by detections of broad
line wings in the 63 $\mu$m line with SOFIA/GREAT (Leurini et al., 2015;
Kristensen et al., 2017a; Yang et al., 2022a). Thus, the high$-J$ CO and [O i]
emission toward Gy 3–7 might also arise from outflows, where shocks and UV
radiation both contribute to gas cooling (see Section 5).

The [C ii] emission also follows the pattern of the continuum emission, but
its two emission peaks are offset by $\sim$11$\arcsec$ in different directions
from the corresponding [O i] 63 $\mu$m peaks. These differences are
illustrated further in Figure 6, in which emission in various species is
directly compared. Clearly, the [C ii] emission traces different regions of Gy
3–7 than the high$-J$ CO and [O i] lines, which is at least partly due to its
lower critical density, $n_{\mathrm{crit}}\sim$(3.7–4.5)$\times 10^{3}$ cm-3
for $T_{\mathrm{kin}}$ of 300-100 K (assuming collisions with H2, Wiesenfeld &
Goldsmith, 2014). As a result, [C ii] is excited in lower density regions and
likely exposed to external UV radiation creating a photodissociation region
(Hollenbach & Tielens, 1997). To some extent, the pattern of [C ii] emission
might be also affected by self-absorption, which is spectrally-unresolved in
the FIFI-LS data (Section 3.1).

##  4 Analysis

Figure 7:  CO rotational diagrams toward core A (left, in blue) and B (right,
in orange). Circles refer to values based on line detections and the triangle
shows the measurement using the upper limit of the CO 31-30 line, which is
blended with OH. Solid lines show fits using transitions belonging to the
”warm” gas component; the CO $30-29$ line at 87.19 $\mu$m is therefore not
included. The CO rotational temperature $T_{\mathrm{rot}}$ derived from the
rotational diagram is labeled in each panel and the value in parenthesis
indicates its uncertainty.

The high$-J$ CO emission allows us to measure the CO rotational temperature of
the warm molecular gas and to estimate the total line cooling by the FIR CO
lines. The mapping capabilities of FIFI-LS provide information about the
spatial distribution of the temperature and gas cooling across the entire
clump.

###  4.1 CO rotational temperatures

We used Boltzmann (or rotational) diagrams to calculate the CO rotational
temperature toward the two dense cores in Gy 3–7 as a proxy of gas kinetic
temperature. Assuming that all these lines are optically thin and thermalized,
their upper level column densities, $N_{\mathrm{u}}$, are estimated using Eq.
1 following Goldsmith & Langer (1999):

$\ln\frac{N_{\mathrm{u}}}{g_{\mathrm{u}}}=\ln\frac{N_{\mathrm{tot}}}{\mathcal{Q}(T_{\mathrm{rot}})}-\frac{E_{\mathrm{u}}}{k_{\mathrm{B}}T_{\mathrm{rot}}},$
(1)

where $g_{\mathrm{u}}$ is the degeneracy of the upper level,
$\mathcal{Q}(T_{\mathrm{rot}})$ is the rotational partition function at a
temperature, $T\rm{{}_{rot}}$, $N_{\mathrm{tot}}$ is the total column density,
and $k_{\mathrm{B}}$ is the Boltzmann constant (see Table 1).

Due to the low spatial resolution of FIFI-LS, the emitting region of the
highly-excited gas is unresolved and thus we calculate instead the number of
emitting molecules, $\mathcal{N}_{\mathrm{u}}$, for each transition (see e.g.,
Herczeg et al., 2012; Karska et al., 2013):

$\mathcal{N}_{\mathrm{u}}=\frac{4\pi d^{2}\lambda F_{\lambda}}{Ahc},$ (2)

where $F_{\lambda}$ is the flux of the line at wavelength $\lambda$, $d$ is
the distance to Gy 3–7, $A$ is the Einstein coefficient, $c$ is the speed of
light, and $h$ is the Planck’s constant. Consequently, the total number of
emitting molecules $\mathcal{N}_{\mathrm{tot}}$ is derived instead of the
$N_{\mathrm{tot}}$ from Eq. 1.

Table 2: CO rotational temperature, the number of emitting molecules, and
total line luminosities of CO and [O i] lines toward dense cores in Gy 3–7 and
IM YSOs from Matuszak et al. (2015)

Source |  $T_{\mathrm{rot}}$aa$a$$T_{\mathrm{rot}}$ and $\mathcal{N}_{\mathrm{tot}}$ calculated using only CO $14-13$, CO $16-15$, CO $18-17$, and CO $22-21$ transitions for more direct comparisons with the cores in Gy 3–7. | log$\mathcal{N}_{\mathrm{tot}}$aa$a$$T_{\mathrm{rot}}$ and $\mathcal{N}_{\mathrm{tot}}$ calculated using only CO $14-13$, CO $16-15$, CO $18-17$, and CO $22-21$ transitions for more direct comparisons with the cores in Gy 3–7. |  $L_{\mathrm{CO}}$(warm)bb$b$The CO line luminosity in the ” warm” component is obtained using all transitions reported in Matuszak et al. (2015) with $J_{\mathrm{up}}$ from 14 to 24 (see Section 4.2). | $L_{\mathrm{[{O\textsc{i}}]63\mu m}}$ | $L_{\mathrm{[{O\textsc{i}}]145\mu m}}$
---|---|---|---|---|---
| (K) |  | ($10^{-3}L_{\odot}$)
Dense cores in Gy 3–7
Core A | $305\pm 85$ | $50.41\pm 0.34$ | $73.31\pm 19.69$ | $92.69\pm 5.01$ | $8.38\pm 0.29$
Core B | $155\pm 20$ | $50.96\pm 0.24$ | $30.47\pm 6.10$ | $76.17\pm 5.37$ | $6.15\pm 0.37$
IM YSOs in the inner Milky Way
AFGL 490cc$c$We adopted the new distance of 0.97 kpc toward this source when calculating $T_{\mathrm{rot}}$, $\mathcal{N}_{\mathrm{tot}}$, and luminositiy of CO and [O i] lines (see Table 9 and references therein).  | $255\pm 25$ | $51.11\pm 0.13$ | $225.17\pm 19.85$ | $336.65\pm 1.47$ | $29.56\pm 0.24$
L1641 | $270\pm 35$ | $50.06\pm 0.17$ | $24.31\pm 2.74$ | $14.69\pm 0.14$ | $1.20\pm 0.05$
NGC 2071 | $295\pm 20$ | $51.23\pm 0.10$ | $425.38\pm 26.11$ | $312.72\pm 0.91$ | $21.59\pm 0.13$
Vela 17 | $215\pm 25$ | $51.08\pm 0.17$ | $125.56\pm 15.44$ | $369.84\pm 1.33$ | $77.85\pm 0.21$
Vela 19 | $255\pm 20$ | $50.66\pm 0.10$ | $86.00\pm 9.27$ | $130.69\pm 0.70$ | $13.60\pm 0.09$
NGC 7129 | $295\pm 35$ | $50.84\pm 0.14$ | $187.26\pm 19.16$ | $66.03\pm 0.78$ | $3.13\pm 0.88$
666

Figure 8: Gas and dust temperatures toward Gy 3–7. Left: CO rotational
temperature map obtained using Boltzmann diagrams. White contour map shows the
CO $17-16$ emission at 25%, 50%, 75%, and 95% of the emission peak. Right: H2
column density contour map overlaid on the dust temperature map derived from
the $\tt ppmap$ tool with the Herschel/Hi-GAL survey (Marsh et al., 2017).
Contour levels of the $N$(H2) are at (5, 10, 30, 50, 55)$\times 10^{21}$ cm-2.

Figure 9: Spatial extent of FIR line luminosities of CO (left) and [O i]
(right) toward Gy 3–7. The calculation procedure is described in Section 4.2.

To measure $T_{\mathrm{rot}}$ over the same physical scales, we convolved the
CO emission maps down to the lowest spatial resolution corresponding to CO
$14-13$ observation (18.3$\arcsec$, see Table 1) and resampled the maps to the
same pixel size. The flux of CO lines toward cores A and B was calculated
within a beam of 20$\arcsec$ (see Table 6). We performed a linear fit on the
rotational diagram using the curve-fit function in Python ($y=ax+b$). The
$T\rm{{}_{rot}}$ and $\mathcal{N}_{\mathrm{tot}}$ values are derived from the
slope $a$ and y-intercept $b$ of the fit (Eq. 1).

Figure 7 shows the CO Boltzmann diagrams toward the two cores in Gy 3–7. We
obtain rotational temperatures of $305\pm 85$ K and $155\pm 20$ K toward dense
cores A and B, respectively, using CO lines with $J_{\mathrm{up}}$ of $14-22$
(see Table 2). The same transitions have been associated with the “warm”
component detected on CO diagrams toward protostars in the inner Galaxy and
corresponding to the widely found $T_{\mathrm{rot}}$ of 300 K (Karska et al.,
2013; Manoj et al., 2013; Green et al., 2013). The CO $30-29$ and CO $31-30$
data at 87.2 and 84.4 $\mu$m, respectively, indicate the presence of the “hot”
component toward core A (Karska et al., 2013). However, the 84.4 $\mu$m line
is blended with OH, which clearly affects the CO line flux (see Figure 3).
Consequently, we are not able to constrain the “hot” component using only the
CO 30-29 line.

The spatial distribution of the “warm” component’s CO rotational temperature
toward the entire Gy 3–7 clump is shown in the left panel of Figure 8. Here,
we calculate $T_{\mathrm{rot}}$ using the three lowest transitions
($J_{\mathrm{up}}$ of $14-17$), which are detected in a large part of the map,
except for the map edges where even CO $17-16$ is not detected (Figure 4). The
resulting $T_{\mathrm{rot}}$ values range from 105 to 230 K across the map,
with a median value of 170(15) $\pm$ 30 K777The temperature in the bracket
shows the mean error of the $T\rm{{}_{rot}}$ distribution and $\pm$ 30 refers
to the standard deviation of the $T\rm{{}_{rot}}$ values.. The CO rotational
temperatures are the highest in the vicinity of the CO $17-16$ emission peaks,
which are offset by $\sim$17$\arcsec$ from core A to the west (Figure 4). The
morphology of the $T_{\mathrm{rot}}$ distribution suggests the origin of
high$-J$ CO in a bipolar outflow driven by core A. Significantly lower CO
temperatures, $\lesssim$150 K, are measured in the surroundings of core B,
without a clear outflow signature from this object. The temperatures around
200 K at the eastern edge of the map are likely caused by higher gas densities
in this region, rather than higher gas kinetic temperatures.

The spatial distribution of dust temperatures $T_{\mathrm{dust}}$ toward Gy
3–7, adopted from the Herschel/Hi-GAL survey, is also shown in the right panel
of Figure 8. The morphology of regions with elevated temperatures is similar
to the extent of $N$(H2), and shows two peaks at 27 and 21 K toward core A and
B, respectively. The pattern differs significantly from the distribution of
warm, $\sim$300 K gas, traced by CO lines, favoring the origin of the CO
emission in a bipolar outflow driven by core A. Nevertheless, the lowest$-J$
CO transitions observed with FIFI-LS might partly trace the extended continuum
emission, for instance, on the eastern part of Gy 3–7. We discuss this issue
further in Section 5.1.

###  4.2 FIR line cooling

The FIR line cooling budget in LM protostars is sensitive to a source’s
evolutionary stage and $L_{\mathrm{bol}}$, and its contributions provide
important information on the shock origin of the FIR emission (Karska et al.,
2013, 2018). Here, we quantify the luminosity of the FIR CO and [O i] lines
toward the two cores in Gy 3–7 and a sample of IM protostars observed with
Herschel/PACS (Matuszak et al., 2015).

####  4.2.1 Calculation procedure

We calculated the line cooling following procedures developed for
Herschel/PACS observations in the same wavelength range (Karska et al., 2018).
Briefly, we determined the line luminosity of CO lines in the ” warm”
component, $L_{\mathrm{CO}}$(warm), from the sum of the individual line fluxes
with $J_{\mathrm{up}}$ from 14 to 24, corresponding to
$E\mathrm{{}_{u}}/k\mathrm{{}_{B}}$=580–1800 K. Since not all CO transitions
are observed, we use linear fits from the Boltzmann diagram to the ”warm”
component to recover the fluxes of the transitions not covered by FIFI-LS or
PACS observations. The flux uncertainties of those transitions were propagated
from the parameters of the linear fit and their uncertainty. The same
procedure could be applied to the ”hot” component toward Gy 3–7 cores due to
the lack of data on a sufficient number of high$-J$ CO transitions. As a
result, we did not calculate the total FIR CO cooling, that is, we did not
account for transitions with $E_{\mathrm{u}}$¿1800 K. The advantage of this
approach is that we avoid significant source-to-source variations in
$L_{\mathrm{CO}}$(hot), which is reflected by the broad range of rotational
temperatures measured using the highest$-J$ CO lines (see Figure 6 in Karska
et al., 2018). The total line luminosity of [O i],
$L_{\mathrm{[{O\textsc{i}}]}}$, is calculated by the addition of the fluxes of
the [O i] 63 and 145 $\mu$m lines. Table 2 shows the FIR line luminosities
obtained for Gy 3–7, as well as those for six IM YSOs from Matuszak et al.
(2015).

Figure 10: Correlations between luminosities of FIR CO and [O i] lines. Top:
Correlation between the luminosities of the 63 and 145 $\mu$m [O i] lines at
from LM to HM YSOs. Cores A and B in Gy 3–7 are marked as grey diamonds, Class
0 and Class I YSOs as orange ”$\times$” signs and red circles (Karska et al.,
2018), respectively, IM YSOs as blue circles (Matuszak et al., 2015), and HM
YSOs as green squares, respectively (Karska et al., 2014a). Black solid line
is the linear fit to all sources except for the HM YSOs, showing a strong
correlation between the two [O i] line luminosities. Bottom: Correlation
between luminosities of the CO lines and the [O i] line at 145 $\mu$m. Black
solid line shows the linear fit to all sources, including the HM YSOs.

Figure 9 shows the spatial distribution of the FIR line luminosity of CO and
[O i] toward Gy 3–7. Most of the CO luminosity orginates in a region west of
core A, suggesting an outflow origin (see also Section 3.2). A similar region
is characterized also by a high [O i] luminosity, which typically follows the
pattern of high$-J$ CO emission around LM protostars (Karska et al., 2013;
Nisini et al., 2015; van Dishoeck et al., 2021). Additionally, a high [O i]
luminosity is measured to the east from core B, toward the direction of the [C
ii] peak (see Figures. 5 and 6). The lack of enhancement of CO line luminosity
in this region could suggest an origin of the [O i] emission in the PDR. We
explore this scenario further in Section 4.3.

Figure 11: Ratio of the CO and [O i] luminosities as a function of pre-shock
density for UV irradiated $C-$shock models and observations of Gy 3–7 cores
and IM YSOs from Matuszak et al. (2015) (light blue box). All models
correspond to UV fields parameterized by $G_{\mathrm{0}}$ of 10 (in blue) and
0.1 (in orange). Solid lines connect models with shock velocities
$\varv_{\mathrm{s}}$ of 20 km s-1, and dashed lines – the models with
$\varv_{\mathrm{s}}$ of 10 km s-1.

####  4.2.2 Flux correlations

We compare the FIR line emission in Gy 3–7 with the data based on
Herschel/PACS measurements toward LM Class 0 and Class I (Karska et al.,
2018), IM (Matuszak et al., 2015), and HM YSOs (Karska et al., 2014a). CO
luminosities are consistently calculated for the ”warm” component on Boltzmann
diagrams. The correlations between luminosities of the two [O i] lines, as
well as the [O i] and CO lines are shown in Figure 10.

A strong power-law correlation is found between the [O i] line luminosities
for all low- and intermediate-mass sources including the two cores in Gy 3–7
(the Pearson coefficient of the correlation is $r=0.996$, which corresponds to
a significance of $8.2\sigma$). The observed ratios of the [O i] 63 $\mu$m/145
$\mu$m lines span a range from 4 to 38 with a median value of $\sim$12\. These
results are quantitatively similar to those shown in Figure 11 of Karska et
al. (2013) for a sub-sample of 18 LM YSOs.

The HM YSOs follow a similar trend, however, several sources show a flux
deficit in the [O i] 63 $\mu$m line, which is likely caused by line-of-sight
contamination and optical depth effects (see, e.g., Liseau et al., 1992;
Leurini et al., 2015). A power-law fit to the sample including HM YSOs shows a
shallower slope ($b=0.91$ versus $b=0.99$, when only LM and IM YSOs are
considered). The correlation strength is comparable to the one for LM and IM
sources alone, with a Pearson coefficient of 0.991 corresponding to
8.5$\sigma$.

The CO line luminosity in the ”warm” component shows a strong correlation with
the 145 $\mu$m [O i] line luminosity. A power–law fit to the entire sample
returns a slope of $b\sim$0.93 and the Pearson coefficient of $0.971$,
corresponding to $\sim$8.2$\sigma$ (Figure 10). Clearly, the line luminosity
of the [O i] 145 $\mu$m line shows a smaller scatter for the HM YSOs with
respect to the 63 $\mu$m line. In case of LM YSOs, a significant scatter in
the ratio of CO and [O i] line luminosities is likely linked to their
different evolutionary stages. The ratio of CO line luminosity over
$L_{\mathrm{bol}}$ is $\sim$2.3 larger for Class 0 than Class I sources,
whereas the [O i] luminosities are similar for both groups (Karska et al.,
2018). Thus, the molecular-to-atomic line cooling is expected to be higher in
Class 0 objects. Indeed, the linear fit using only Class 0 sources results in
a shallower slope, but do not affect the general conclusions.

In summary, we find a strong correlation between the line luminosities of the
[O i] and CO lines for YSOs in a broad mass range, consistent with previous
results for LM YSOs (Karska et al., 2013). Combined with the similar spatial
extent of the [O i] and CO lines (Section 3.2), the correlations suggest a
similar physical origin of the two species.

###  4.3 Properties of a possible photodissociation region

Assuming that the [O i] and [C ii] lines predominately arise from a
photodissociation region (PDR), the UV field strengths, $G_{\mathrm{0}}$, and
hydrogen nucleus number densities, $n_{\mathrm{H}}$, across Gy 3–7 can be
obtained from their ratios. These assumptions might be justified in case of
the eastern part of Gy 3–7 with relatively weak CO line luminosities (see
Section 4.2). On the contrary, the [O i] line luminosity in the surrounding of
core A closely follows the high$-J$ CO emission associated with outflow
shocks.

We determined the physical properties of the PDR using the PDR Toolbox
2.1.1888https://dustem.astro.umd.edu/ (Pound & Wolfire, 2011) based on the PDR
models provided by Kaufman et al. (2006). We used three line ratios involving
the [O i] lines at 63.2 $\mu$m and 145.5 $\mu$m, and [C ii] line at 157.7
$\mu$m, and ran the code at each spaxel of the FIFI-LS maps. We obtained gas
densities of 104–105 cm-3 and UV field strengths of the order of 103-106 times
the average interstellar UV radiation field (Habing, 1968). These physical
conditions are typical for dense, star-forming clumps associated with HM YSOs
(Ossenkopf et al., 2010; Benz et al., 2016; Mirocha et al., 2021). However,
given that Gy 3–7 is associated with IM YSOs (see Section 4.5), UV radiation
fields of 103 or higher are unlikely (e.g., Karska et al., 2018).

The similar spatial extent of the [O i] and high$-J$ CO emission (Section 3.2)
and the strong line luminosity correlation between the two species (Section
4.2) favor the origin of the bulk of [O i] in the outflow shocks rather than
in the photodissociation region. We consider this scenario in Section 4.4.

Table 3: Best-fit models of SED using the Robitaille (2017) classification

Model Set | # of sources | Components | Group
---|---|---|---
s$---$s$-$i | 4 | star | –
sp$--$s$-$i | 2 | star $+$ passive disk; $R_{\rm inner}$ = $R_{\rm sub}$ | $d$
sp$--$h$-$i | 4 | star $+$ passive disk; variable $R_{\rm inner}$ | $d$
s$-$pbhmi | 1 | star $+$ power-law envelope $+$ cavity $+$ medium; variable $R_{\rm inner}$ | $e$
spubsmi | 1 | star $+$ passive disk $+$ Ulrich envelope $+$ cavity $+$ medium; $R_{\rm inner}$ = $R_{\rm sub}$ | $d+e$

Notes: Seven characters in the model set names indicate which component is
present; they are (in order): s: star; p: passive disk, p or u: power-law or
Ulrich envelope; b: bipolar cavities; h: inner hole; m: ambient medium; and i:
interstellar dust. A dash ($-$) is used when a component is absent. $R_{\rm
inner}$ is the inner radius for the disk, envelope, and the ambient medium -
when one or more of these components are present. $R_{\rm sub}$ is the dust
sublimation radius.

Figure 12: SEDs of YSOs in Gy 3–7 well-fitted with Robitaille (2017) models
with envelopes. The best fit model is indicated with the black solid line and
gray lines show the YSO models with $\chi^{2}$ between
$\chi^{2}_{\mathrm{best}}$ and $\chi^{2}_{\mathrm{best}}$\+ $F\times n$, where
$n$ is the number of data points and $F$ is a threshold parameter which we set
to 3 (Sewiło et al., 2019). Filled black circles are valid flux values with
uncertainties. The values of a reduced $\chi^{2}$ and interstellar visual
extinction for the best-fit model are indicated in the plots. Appendix C shows
the SEDs for the remaining YSOs in Gy 3–7.

###  4.4 Comparisons with UV-irradiated shocks

Bright FIR emission detected toward LM YSOs has been interpreted in the
context of continuous ($C-$type) shocks irradiated by UV photons (Karska et
al., 2014b, 2018; Kristensen et al., 2017b). Figure 11 shows a comparison of
the FIR observations toward IM YSOs, including two cores in Gy 3–7, and the
UV-irradiated shock models from Melnick & Kaufman (2015) and Karska et al.
(2018).

Predictions of shock models were previously calculated for high$-J$ CO lines
covered by Herschel/PACS. Here, we show the predictions for UV field
strengths, $G_{\mathrm{0}}$, of 0.1 and 10, and shock velocities of 10 and 20
km s-1 (see also Fig. 14 in Karska et al., 2018). For the sake of comparison,
we calculated the total FIR line luminosity of CO using the transitions in the
”warm” component (Section 4.2.1), as well as high$-J$ CO lines, both for Gy
3–7 and for IM YSOs (Matuszak et al., 2015). The [O i] line luminosity is
calculated from the sum of the two fine-structure [O i] lines.

The models show a good match with observations for pre-shock H2 number
densities of $10^{5}$ cm-3 and the entire range of the considered UV radiation
field strengths. For shock velocities, $\varv_{\mathrm{s}}$, of 20 km s-1, a
possible match is also found for pre-shock densities of 104 cm-3 and G0 of
0.1. The compression factor $>$10 is expected in $C-$type shocks (Karska et
al., 2013), so the gas densities are of the order of 105-106 cm-3, in
agreement with those of LM YSOs (e.g., Kristensen et al., 2012; Mottram et
al., 2017).

Table 4: Physical parameters for a subset of YSO candidates with at least five
photometric data points

ID | Model | Class | $R_{\ast}$ | $T_{\ast}$ | $L_{\ast}$ | $M_{\ast}^{b}$ | Ageb | Notec
---|---|---|---|---|---|---|---|---
| seta |  | ($R_{\odot}$) | (K) | ($L_{\odot}$) | ($M_{\odot}$) | (Myr) |
1 | sp–s-i | II/III | 5.2 | 3592 | 4.0 | … | … |
2 | sp–s-i | II/III | 2.6 | 6791 | 13.2 | … | … | TP01
4 | s—s-i | star | 1.2 | 2955 | 0.1 | … | … |
5 | s—s-i | star | 1.0 | 3388 | 0.1 | … | … |
6 | s—s-i | star | 1.0 | 5120 | 0.6 | … | … | TP03
9 | sp–h-i | II/III | 3.3 | 3359 | 1.2 | … | … |
10 | s-pbhmi | 0 | 25.8 | 6214 | 892.1 | 8.0 | 0.06 | TP06
11 | s—s-i | star | 1.2 | 2955 | 0.1 | … | … |
12 | spubsmi | I | 23.0 | 6194 | 702.1 | … | … | TP07
13 | sp–h-i | II/III | 2.6 | 6112 | 8.4 | … | … | TP09
14 | sp–h-i | II/III | 3.6 | 3736 | 2.2 | … | … |
15 | sp–h-i | II/III | 10.3 | 7101 | 244.9 | … | … | TP08

a See Table 3 footnotes for the description of the model set names.

b We provide stellar masses ($M_{\rm\ast}$) and ages only for sources with
reliable estimation of these parameters (see text for details).

c IDs of the YSO candidates in Gy 3–7 identified by Tapia et al. (1997), see
Figure 1.

Due to the lack of H2O observations from FIFI-LS and non-detections of OH, we
are limited to comparisons between CO and [O i] lines. Some of the [O i]
emission might arise in the PDR (Section 4.3), which would increase the
observed line luminosity ratio; however, the strong correlation of [O i] and
high$-J$ CO tracing outflow shocks does not support this scenario.

###  4.5 Spectral energy distribution analysis

We investigate the physical properties of 15 individual YSO candidates from
Tapia et al. (1997) and Sewiło et al. (2019) in Gy 3–7 to understand their
possible impact on the FIR and submillimeter line emission. Spectral energy
distribution (SED) models of YSOs from Robitaille (2017) are fitted to the
multi-wavelength photometry of YSOs using a dedicated fitting tool (Robitaille
et al., 2007). Multi-wavelength photometry spanning from near- to mid-IR range
of YSO candidates in the IRAS field is presented in Appendix C.

We followed the procedures described in detail in Karska et al. (2022). We
used 18 sets of model SEDs including various physical components of a YSO:
star, disc, in-falling envelope, bipolar cavities, and an ambient medium
(Robitaille, 2017). We used the PARSEC evolutionary tracks produced by the
revised Padova code (Bressan et al., 2012; Chen et al., 2014, 2015; Tang et
al., 2014) to quantify the results of SED model fitting. Models producing YSO
parameters outside of the PARSEC pre-main sequence (PMS) tracks were excluded.
YSOs with models in line with the PARSEC tracks are illustrated on the
Hertzsprung-Russell diagram in Appendix C. For those YSOs, we calculate the
stellar luminosity from the Stefan-Boltzmann law, using the stellar radius and
effective temperature from the SED fitting. The masses and ages are determined
from the closest PMS track; however, we only provide the stellar masses and
ages that are consistent with the SED fitting results (i.e., the evolutionary
stage) and YSO lifetimes from Dunham et al. (2015), respectively (see Section
3.7 in Karska et al., 2022).

Table 3 shows the best-fit SED models for 12 YSOs in Gy 3–7. The SEDs of two
sources require an envelope contribution which is typical for deeply-embedded
Class 0/I YSOs (see Figure 12), while six sources are successfully modeled
with a passive disk and four are normal stars (see Appendix C). The resulting
physical parameters determined from SED models are shown in Table 4.

The two YSOs with envelopes, No. 10 and 12 in Table 4, are located in the
center of Gy 3–7 (see Figures 1 and 5), and their strong IR excess has already
been noted by Tapia et al. (1997). The Class 0 YSO is co-spatial with the IRAS
source and the dense core A (Elia et al., 2021), and might be the source of
the outflow responsible for the FIR emission. Both objects are in the IM
regime based on their luminosities obtained from SEDs. The four objects that
are modeled as stars with foreground extinction have photometry only from 1 to
5 $\mu$m. In this range, it is difficult to distinguish between stars with
foreground extinction and stars with disks, so we do not rule out the latter
explanation.

##  5 Discussion

Far-IR observations from FIFI-LS confirm the status of Gy 3–7 as a deeply-
embedded cluster, as originally proposed by Tapia et al. (1997). Here, we
discuss the likely origin of FIR emission in Gy 3–7 and we search for any
effects coming from metallicity by comparison with YSOs from the inner Galaxy.

###  5.1 Origin of FIR emission in Gy 3–7

A strong correlation of high$-J$ CO and [O i] luminosities and their spatial
extent in Gy 3–7 provides a strong support toward the common origin of the two
species. The modeling of envelopes of HM YSOs showed that emission from shocks
is necessary to reproduce line fluxes of high$-J$ CO lines (Karska et al.,
2014a). The same is certainly the case for LM and IM YSOs as well, which are
characterized by lower envelope densities and temperatures. On the other hand,
the [O i] emission could be associated with a photodissociation region
(Kaufman & Neufeld, 1996a; Hollenbach & Tielens, 1997) or outflow shocks
(Karska et al., 2013; Nisini et al., 2015). Recent velocity-resolved profiles
of the [O i] line at 63 $\mu$m with SOFIA/GREAT support the later scenario for
LM YSOs, which does not suffer from strong self-absorption (Kristensen et al.,
2017a; Yang et al., 2022a). Therefore, we assume that the entire high$-J$ CO
and [O i] emission originates from outflow shocks. As shown in Section 4.4,
the line luminosity ratio of CO and [O i] is consistent with $C-$type shocks
irradiated by UV fields of 0.1–10 times the average interstellar radiation
field and pre-shock densities of $10^{4}$–$10^{5}$ cm-3 (see also, Melnick &
Kaufman, 2015; Karska et al., 2018). The detection of variable H2O maser
further confirms the outflow activity in the region (Furuya et al., 2003a). We
provide details in Appendix A.

The physical conditions in Gy 3–7 are also constrained by CO rotational
temperatures, which provide a good proxy of gas kinetic temperatures (Karska
et al., 2013; van Dishoeck et al., 2021). The high$-J$ transitions are likely
optically thin and thermalized at gas densities ¿$10^{5}$ cm-3 that are
routinely measured for LM and IM YSOs (e.g., Kristensen et al., 2012; Mottram
et al., 2017). The top panel in Figure 13 shows that $T_{\mathrm{rot,CO}}$
associated with the two dense cores in Gy 3–7 is either fully consistent (core
A) or at the low-end of other IM YSOs (core B), for which the median value
from the literature is $320(33)\pm 35$ K (Matuszak et al., 2015). Similar
rotational temperatures have been measured for LM and HM YSOs (Karska et al.,
2013, 2014a, 2018; Green et al., 2013, 2016; Manoj et al., 2013; Yang et al.,
2018), with average values of $328(33)\pm 63$ K (Karska et al., 2018) and
$300(23)\pm 60$ K (Karska et al., 2014a), respectively.

Some differences between the dense core B in Gy 3–7 and the other YSOs studied
in the literature might result from a smaller number of transitions probed by
FIFI-LS. The curvature seen in the rotational diagrams causes the rotational
temperature to be underestimated if high$-J$ CO transitions are not observed.
Therefore, we re-calculated all literature measurements of $T_{\mathrm{rot}}$
for IM and HM YSOs using the same or similar transitions as obtained for Gy
3–7 (Table 2 and Appendix D). As expected, CO rotational temperatures for IM
and HM YSOs are now lower than reported in the literature and consistent even
with core B (see bottom panel of Figure 13).

Figure 13: CO rotational temperatures for dense cores in Gy 3–7 and
intermediate- and high-mass YSOs from the literature (Karska et al., 2014a;
Matuszak et al., 2015). Top: Straightforward comparison with the literature
values. Bottom: Comparison accounting for the number of observed CO lines
considered in the rotational diagrams.

###  5.2 Possible impact of metallicity on FIR line emission in the outer
Galaxy

Far-IR observations of HM YSOs in the low-metallicity environments of the SMC
and LMC show a lower fraction of molecular-to-atomic emission with respect to
Galactic YSOs (Oliveira et al., 2019). Here, we investigate the impact of
bolometric luminosity, $L_{\mathrm{bol}}$, and source Galactocentric radius,
$R_{\mathrm{GC}}$, on the ratio of CO and [O i] line luminosity.

Figure 14 shows the molecular-to-atomic ratio toward Gy 3–7 cores and other
Galactic IM and HM YSOs as a function of $L_{\mathrm{bol}}$. The values of
$L_{\mathrm{bol}}$ for the two cores in Gy 3–7 are adopted from Elia et al.
(2021), and are equal to 75.9 and 324.2 $L_{\odot}$, respectively. The
correlation is characterized by a Pearson coefficient of 0.19, corresponding
to 1.6$\sigma$. Since the number of LM YSOs exceeds by far the number of IM
and HM sources, we also search for trends in the binned datasets. We bin
sources in equal intervals of log$L_{\mathrm{bol}}$=1, and adopt a 1$\sigma$
variance of the distribution as the uncertainty inside the bin. As a result,
we confirm a weak correlation between the ratio of line luminosity and
$L_{\mathrm{bol}}$ ($r\sim$0.59, corresponding to 1.5$\sigma$). Thus, our
large sample of YSOs starts to reveal a relationship between the mass of YSOs
and the fraction of molecular-to-atomic cooling. This has not been possible
with a sample of 18 LM YSOs and 10 HM YSOs analyzed in Karska et al. (2014a),
where a ratio of $\sim$4 was reported for YSOs in the entire mass regime. We
note, however, that the Herschel/PACS measurements from the literature
benefited from additional detections of H2O and OH, which were included in the
molecular FIR cooling.

Figure 14: Correlations between the ratio of CO and [O i] line luminosities
and bolometric luminosity of YSOs in the two cores in Gy 3–7 (gray diamonds),
Class 0 and Class I YSOs (orange ‘$\times$’ signs and red circles,
respectively; Karska et al., 2018), IM YSOs (blue circles; Matuszak et al.,
2015), and HM YSOs (green squares; Karska et al., 2014a). The bolometric
luminosities for Gy 3–7 cores are adopted from Elia et al. (2021). The top
panel shows a power-law fit to all individual data points, and the bottom
panel shows the fit to the data bins (both shown as black dashed line).
Figure 15: Ratio of molecular to atomic line luminosities as a function of the
Galactocentric radius. The light coral box indicates the range of the line
luminosity ratio and the Galactocentric radius for LM YSOs in the nearby
clouds, which are excluded from the power-law fit to the remaining YSOs
(dashed black line). Gray dotted horizontal line indicates where the molecular
luminosity is equal to the atomic luminosity. Figure 16: Ratio of molecular to
atomic line luminosities as a function of metallicity, $Z$. HM YSOs in the SMC
and LMC are shown with cross and ” $\times$” symbols in orange and purple
colors, respectively. The light coral box indicates the range of the line
luminosity ratio and the metallicity for LM YSOs in the nearby clouds, which
are excluded from the power-law fit to the remaining YSOs (dashed black line).
The gray dotted horizontal line indicates where the molecular luminosity is
equal to the atomic luminosity.

Figure 15 shows the molecular-to-atomic ratio as a function of Galactocentric
radius. In addition, Figure 16 shows the same ratio as a function of
metallicity, $Z$, estimated assuming elemental abundance gradients derived
from H ii regions (Balser et al., 2011). Appendix D provides heliocentric
distances, Galactocentric radii, and metallicities of Galactic sources, as
well the measurements of CO and [O i] for YSOs in both the Milky Way and the
Magellanic Clouds.

The molecular-to-atomic ratio decreases very weakly with the Galactocentric
radius (Fig. 15) but the Pearson coefficient is too small to confirm the
correlation ($r=0.08$, corresponding to 0.7$\sigma$). Additional observations
of YSOs in the outer Galaxy would be necessary to constrain the molecular-to-
atomic ratio, in particular, in the gap between 10 and 12 kpc. Due to the lack
of such observations, we searched for trends as a function of metallicity by
including YSOs in the Magellanic Clouds (Fig. 16). We note, however, that the
[O i] line at 145 $\mu$m was not observed toward sources in the Magellanic
Clouds, and the ratio of the two [O i] lines calculated for LM YSOs from
Karska et al. (2018) was adopted to estimate its fluxes (Oliveira et al.,
2019). As seen in Figure 10, this ratio may differ for HM YSOs due to self-
absorption or optical depth effects in the [O i] 63 $\mu$m line. Additionally,
the total CO luminosity of YSOs in the SMC and LMC was estimated using the
detections in SPIRE and the conversion factor from Yang et al. (2018), also
based on LM YSOs.

In conclusion, Gy 3–7 follows closely the correlations set by YSOs observed in
the Milky Way and the Magellanic Clouds. The ratio of molecular-to-atomic line
emission is dominated by source bolometric luminosities, and only a very weak
decreasing trend with the Galactocentric radius was detected.

##  6 Conclusions

We investigated the SOFIA/FIFI-LS maps of the CO transitions from $J=14-13$ to
$J=31-30$, the [O i] lines at 63.2 $\mu$m and 145.5 $\mu$m, and the [C ii] 158
$\mu$m line toward the outer Galaxy cluster Gy 3–7. Spatial information on the
FIR emission enabled us to quantify such physical parameters as temperatures,
densities, and UV radiation fields, as well as to associate them with
identified YSOs. Our conclusions are as follows:

  * •

The CO $J=14-13$ to $J=16-15$ emission lines are detected in a significant
part of Gy 3–7, where Herschel/PACS 160 $\mu$m continuum emission is also
strong. Higher$-J$ CO lines up to $31-30$ are clearly detected only toward the
dense core A.

  * •

The spatial extent of the [O i] emission at 63 and 145 $\mu$m is similar to
that of CO $14-13$ and the 70 and 160 $\mu$m continuum emission. The [C ii]
emission is also extended, but shows systematic shifts in the emission peaks
away from the FIR continuum, tracing lower density gas.

  * •

The CO rotational diagrams show the warm components toward two cores with
$T\rm{{}_{rot}}$ of 305 and 155 K, and a range of $T\rm{{}_{rot}}$ from
$\sim$105 to $230$ K throughout the cluster, where only three high$-J$ CO
lines are unambiguously detected. Similar rotational temperatures have been
detected toward IM and HM YSOs in the inner Milky Way calculated using the
same or similar CO transitions.

  * •

A strong correlation of the CO and [O i] line luminosities and their similar
spatial extent point at the common origin in the outflow shocks. The CO / [O
i] line luminosity ratio of Gy 3–7 cores and other intermediate-mass YSOs is
consistent with C-type shocks propagating at pre-shock densities of
$10^{4}-10^{5}$ cm-3 and UV fields of 0.1-10 times the average interstellar
radiation field.

  * •

Physical parameters for 15 YSO candidates in the Gy 3–7 cluster are obtained
from a YSO SED model fitting (Robitaille, 2017). Two sources, corresponding to
Hi-GAL dense cores from Elia et al. (2021), are well-fitted with YSO models
including the envelope, confirming their early evolutionary stage (Class 0/I).
The location of the Class 0 source at the center of Gy 3–7 cluster suggests
that it might be the driving source of the outflow revealed by FIR emission.

  * •

The ratio of warm CO and [O i] at 145 $\mu$m line luminosities from
protostellar envelopes shows a weak decreasing trend with the bolometric
luminosity and Galactocentric radius. We do not identify any significant
dependence of the line cooling in Gy 3–7 on metallicity.

High-resolution submillimeter observations would be necessary to unambiguously
associate FIR emission from FIFI-LS with candidate YSOs and their outflows.
Alternatively, the efficient imaging of Gy 3-7 with the Mid-Infrared
Instrument on board James Webb Space Telescope in F560W and/or F770 W filters
tracing H2 emission would unveil the details in the outflows (Yang et al.,
2022b).

###### Acknowledgements.

NL, AK, MF, MG, MK, and KK acknowledge support from the First TEAM grant of
the Foundation for Polish Science No. POIR.04.04.00-00-5D21/18-00 (PI: A.
Karska). AK also acknowledges support from the Polish National Agency for
Academic Exchange grant No. BPN/BEK/2021/1/00319/DEC/1. This article has been
supported by the Polish National Science Center grant 2014/15/B/ST9/02111 and
2016/21/D/ST9/01098. The material is based upon work supported by NASA under
award number 80GSFC21M0002 (MS). The research is supported by a research grant
(19127) from VILLUM FONDEN (LEK). This work is based (in part) on observations
made with the NASA/DLR Stratospheric Observatory for Infrared Astronomy
(SOFIA). SOFIA is jointly operated by the Universities Space Research
Association, Inc. (USRA), under NASA contract NNA17BF53C, and the Deutsches
SOFIA Institut (DSI) under DLR contract 50 OK 2002 to the University of
Stuttgart. The 32 m radio telescope is operated by the Institute of Astronomy,
Nicolaus Copernicus University and supported by the Polish Ministry of Science
and Higher Education SpUB grant.

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##  Appendix A Water masers in CMa$-l$224

###  A.1 Survey results

Water masers are important signposts of both low- and high-mass star
formation; they are collisionally excited, tracing warm molecular gas behind
shock waves in the environment of YSOs and H ii regions (e.g., Litvak 1969;
Elitzur et al. 1989; Furuya et al. 2003b; Ladeyschikov et al. 2022 and
references therein). Figure 17 shows the pointings of our 22 GHz water maser
survey of the CMa$-l$224 star-forming region with the 32-m radio telescope in
Toruń (RT4; a half-power beam width, HPBW$\sim$106′′, corresponding to
$\sim$0.5 pc at 1 kpc; see Section 2.3), overlaid on the CO 1–0 integrated
intensity image from the Forgotten Quadrant Survey (FQS, Benedettini et al.
2020; HPBW$\sim$55′′ at 115 GHz). The RT4 water maser survey targeted the
fields harboring YSO candidates identified by Sewiło et al. (2019). Out of 205
observed fields, (185, 18, 2) were observed in (2, 3, 6) epochs. Water masers
were detected toward two fields, one centered on Gy 3–7 and the other on
source IRAS 07069$-$1026 in the main filament in CMa$-l$224\. The very low
detection rate is likely the result of the low sensitivity of our observations
and the variability of the maser emission (see Section 2.3).

The detection of the water maser toward IRAS 07069$-$1026 constitutes the
first maser detection toward this source, while the detection toward Gy 3–7
has previously been reported in the literature. Urquhart et al. (2011)
detected the 22 GHz water maser emission toward Gy 3–7 (their source
G224.6075$-$01.0063) with the 100-m Green Bank Telescope (GBT;
HPBW$\sim$30′′). Water masers were detected toward two other locations in
CMa$-l$224 by Valdettaro et al. (2001) with the Medicina 32-m radio telescope
(a similar angular resolution as for RT4), toward IRAS 07077$-$1026 and IRAS
07054$-$1039 (see Figure 17).

Figure 17: CO $1-0$ integrated intensity map of CMa$-l$224\. Yellow circles show all RT4 pointings, while white circles indicate those with the 22 GHz water maser detections: toward Gy 3–7 and source IRAS 07069$-$1026 in the main filament. Blue and red circles indicate the positions of water masers detected toward CMa$-l$224 with the GBT (Urquhart et al. 2011) and the Medicina telescope (Valdettaro et al. 2001), respectively (see text for details). Figure 18: 22 GHz H2O maser spectra obtained toward Gy 3–7 with the 32-m radio telescope in Toruń from December 2019 to May 2020. The date of observations is indicated above each spectrum. Except for the spectrum corresponding to the observation on May 15, 2020, spectra are shifted vertically by 30 Jy (for the observation on March 21, 2020) and 15 Jy (for the remaining observations) to improve the clarity of the figure. Dashed vertical lines show central velocities of the lines detected in the latest (May 15, 2020) and earliest (December 11, 2019) of the six epochs of observations (see also Table 5). Figure 19: 22 GHz H2O maser emission detected toward IRAS 07069$-$1026 (see also Table 5). The H2O maser was detected only in one of the two observing epochs. The spectrum observed on May 21, 2020 is shifted vertically by 10 Jy to improve the clarity of the figure. Table 5: Parameters of the 22 GHz H2O maser lines detected toward CMa$-l224$ with RT4 Obs. date | $V_{\mathrm{p}}$ | $S_{\mathrm{p}}$ | FWHM | $\int SdV$
---|---|---|---|---
| (km s-1) | (Jy) | (km s-1) | (Jy km s-1)
Gy 3–7
2019 Dec 11 | $1.76\pm 0.01$ | $117.33\pm 0.43$ | 0.79$\pm$0.003 | $41.65\pm 0.32$
2020 Feb 21 | $2.28\pm 0.03$ | $7.27\pm 0.25$ | $1.42\pm 0.06$ | $4.67\pm 0.33$
2020 Feb 25 | $2.30\pm 0.04$ | $5.87\pm 0.35$ | $1.36\pm 0.09$ | $3.61\pm 0.45$
2020 Feb 27 | $2.23\pm 0.08$ | $4.41\pm 0.43$ | $1.59\pm 0.18$ | $3.16\pm 0.66$
2020 Mar 21 | $2.53\pm 0.06$ | $5.87\pm 0.33$ | $1.90\pm 0.12$ | $5.04\pm 0.60$
2020 May 15 | $3.94\pm 0.01$ | $23.80\pm 0.33$ | $1.11\pm 0.02$ | $11.93\pm 0.35$
IRAS 07069$-$1026
2020 May 21 | $15.73\pm 0.02$ | $12.15\pm 0.40$ | $0.80\pm 0.03$ | $4.38\pm 0.31$

The water maser spectra for Gy 3–7 and IRAS 07069-1026 for all epochs of the
RT4 observations are shown in Figure 18 and Figure 19, respectively. Gy 3–7
was observed with RT4 six times over five months with the time interval
between subsequent epochs ranging from 2 days to over two months (see Table 5
for observing dates). The water maser emission was detected in all epochs.
IRAS 07069$-$1026 was observed in two epochs separated by about six months
(see Table 5). The water maser emission toward IRAS 07069$-$1026 was only
detected in the second epoch, illustrating a high variability of the water
maser phenomenon. The RT4 spectra of Gy 3–7 also show clear evidence for a
variable nature of the maser emission in this region. Both the peak flux
density and central velocity of the maser spots varied significantly between
April 2019 and May 2020, covering the flux density and velocity ranges of
4.4–117.2 Jy and 1.8–3.9 km s-1, respectively.

The Gauss functions were fitted to the 22 GHz water maser lines observed
towards Gy 3–7 and IRAS 07069$-$1026\. The Gauss fitting results are shown
Table 5: the central velocities ($V_{\mathrm{p}}$) of the water maser lines,
their peak flux densities ($S_{\mathrm{p}}$), full-widths at half maximum
(FWHM), and the integrated flux densities ($\int SdV$). The measured flux
densities are likely underestimated by 10–20% due to the atmospheric
conditions at the observatory site; no correction for atmospheric attenuation
was applied.

###  A.2 Water maser emission in Gy 3–7

Urquhart et al. (2011) report the detection of a single water maser spot
toward Gy 3–7 (G224.6075$-$01.0063) with the peak flux density of 2.57 Jy and
velocity of 28.7 km s-1 ($\sim$25 km s-1 higher than the emission detected
with RT4). However, the inspection of the spectrum provided by Urquhart et al.
(2011) reveals a second, much fainter and broader water maser line centered at
the velocity of $\sim$0 km s-1, only $\sim$2–4 km s-1 lower than the peak
velocities of the lines detected in our observations. No water maser emission
at higher velocities was detected in any of the six RT4 epochs.

The systemic velocity of Gy 3–7 is 16.7 km s-1 based on the NH3 observations
of Urquhart et al. (2011), in a very good agreement with the CO 1–0 velocity
of $\sim$16 km s-1 (Benedettini et al. 2020), indicating that NH3 and CO trace
the same molecular gas. The water maser emission toward Gy 3–7 is blueshifted
(this study) and redshifted (Urquhart et al. 2011) with respect to the
systemic velocity. The blueshifted emission has the similar velocity offset
($\sim$13–16 km s-1) from the systemic velocity as the redshifted emission (12
km s-1), extending the velocity range over which the maser spots are found
toward Gy 3–7 from $\sim$2 km s-1 reported in Urquhart et al. (2011) to
$\sim$30 km s-1.

About 60% of the sources from the Urquhart et al. (2011) sample have total
velocity ranges of the water maser emission of $\lesssim$20 km s-1. The
mean/median velocity range of the entire sample is 24.7/15.2 km s-1. Six
sources have velocity ranges of $>$100 km s-1. The Gy 3–7’s water maser
velocity range of $\sim$30 km s-1 is larger than the velocity ranges of the
majority of the sources from the Urquhart et al. (2011) sample, but it is well
within the observed values.

The water maser data currently available for Gy 3–7 show that the blueshifted
water maser spots are brighter than the redshitfed ones (and thus detected
more often), in agreement with the results obtained by Urquhart et al. (2011)
for a large sample of $\sim$300 YSOs and H ii regions. The higher relative
velocities of blueshifted masers also agree with the Urquhart et al. (2011)’
results.

The difference of $\sim$12–16 km s-1 between the systemic velocity of Gy 3–7
and the velocities of maser spots is small enough to assume that there is the
physical association between the molecular gas traced by NH3 and CO and the
maser emission (possibly originating in outflow shocks); it is less likely
that the maser emission is arising from a different region located along the
same line of sight. The mean difference between the maser and systemic
velocities for the Urquhart et al. (2011) sample of YSOs and H ii regions is
$-$3.8 km s-1, but with the large standard deviation ($\sim$20 km s-1). It is
possible that the water maser emission with smaller offsets from the systemic
velocity of Gy 3–7 exists, but remained undetected due to the maser
variability and/or limited sensitivity of the observations.

The spatial resolution of the existing water maser observations of Gy 3–7 is
too low to accurately pinpoint the location of the maser spots in the region.
Higher resolution observations are needed to investigate the distribution of
the maser spots toward Gy 3–7 and the origin of the maser emission. Systematic
monitoring of the water maser activity at high angular resolution would be
necessary to constrain any periodicity in the maser line intensity in one or
more velocity components (Szymczak et al. 2016).

##  Appendix B Spatial extent of FIR line emission

Figures 20 and 21 show the spatial extent of the [O i] lines at 63 and 145
$\mu$m, the [C ii] line at 157 $\mu$m, the CO lines with
$J_{\mathrm{up}}=14-31$, and the OH line at 79.2 $\mu$m. We calculated the
flux of these emission line toward the two dense cores A and B, within a beam
size of 20$\arcsec$ (Table 6). Since all lines are spectrally-unresolved, the
line fluxes are obtained using Gaussian fits and their uncertainties are
estimated as 1$\sigma$ of the distribution of 10 000 Gaussian profiles
generated based on the parameters of the Gaussian fit and their uncertainty.

Figure 20: FIFI-LS integrated intensity maps of the [O i] lines at 63.2 and
145.5 $\mu$m, the [C ii] line at 157.7 $\mu$m, the CO lines $J=14-13$,
$16-15$, $17-16$ at 186, 162.8, 153.3 $\mu$m, respectively. The white contours
show line emission at 25$\%$, 50$\%$, 75$\%$, and 95% of the corresponding
line emission peak. The ”+” and ”x” signs show the positions of the dense
cores A and B, respectively.

Figure 21: FIFI-LS integrated intensity maps of the CO $22-21$, CO $30-29$ , and CO $31-30$ transitions at 118, 84, and 87 $\mu$m, respectively, and the OH line at 79.2 $\mu$m. The white contours show line emission at 25$\%$, 50$\%$, 75$\%$, and 95% of the corresponding line emission peak, see also Figure 20. Table 6: Flux SOFIA FIFI-LS toward the two dense cores within a beam size of 20$\arcsec$.  $\lambda$ ($\mu$m) | Flux (10-16 W m-2)
---|---
| core A | core B
63.18 | $29.65\pm 1.60$ | $24.37\pm 1.72$
84.41 | $4.93\pm 1.80$ | –
87.19 | $1.81\pm 0.53$ | –
118.58 | $2.68\pm 0.21$ | –
157.74 | $1.15\pm 0.07$ | $1.90\pm 0.09$
145.53 | $2.68\pm 0.09$ | $1.97\pm 0.12$
153.27 | $2.09\pm 0.05$ | $1.09\pm 0.07$
162.81 | $2.84\pm 0.05$ | $1.80\pm 0.07$
186.00 | $2.73\pm 0.08$ | $2.56\pm 0.13$

##  Appendix C Multi-wavelength photometry and SED fitting results

Figure 22: SEDs of YSO candidates with well-fitted Robitaille (2017) YSO
models. The best-fit model is indicated with the black solid line. Gray lines
show the YSO models with $\chi^{2}$ between $\chi^{2}_{\mathrm{best}}$ and
$\chi^{2}_{\mathrm{best}}$\+ $F\times n$, where $n$ is the number of data
points and $F$ is a threshold parameter which we set to 3 (Sewiło et al.
2019). Filled circles and triangles are valid flux values and flux upper
limits, respectively. The values of a reduced $\chi^{2}$ and interstellar
visual extinction for the best-fit model are indicated in the plots. Figure
23: HR diagram with YSOs in Gy 3–7 (blue ‘$\times$’ symbols) and the PARSEC
evolutionary tracks (Bressan et al. 2012; Chen et al. 2014, 2015; Tang et al.
2014).

Table 7 shows the multi-wavelength photometry for 15 YSO candidates in Gy 3–7.
Figure 22 shows the SEDs of YSO candidates in Gy 3–7 with the best-fit
Robitaille (2017) models. Figure 23 shows the Hertzsprung-Russell diagram with
the positions of the YSOs obtained from the SED modeling in line with the PMS
tracks.

We note that sources numbered 3 – 8 and source 11 lack continuum measurements
at $>4.5$ $\mu$m, which limits the SED modeling. The best-fit models for
sources No. 4 – 6 and 11 are stars (Table 4), since the envelope emission
could not be traced. Near-IR observations show some IR excess toward source
No. 6 (Tapia et al. 1997), but the confirmation of its YSO status would
require additional observations, which are outside of the scope of this paper.

Table 7: Multi-wavelength photometry of YSO candidates in the IRAS field. The columns represent the 2MASS $JHK_{s}$, Spitzer IRAC 3.6 and 4.5 $\mu$m, AllWISE, Herschel PACS, and SPIRE ID | RA | Dec | $S_{J}$ | $S_{H}$ | $S_{K}$ | $S_{3.6}$ | $S_{4.5}$ | $S_{\rm W1}$ | $S_{\rm W2}$
---|---|---|---|---|---|---|---|---|---
| (deg) | (deg) | (mJy) | (mJy) | (mJy) | (mJy) | (mJy) | (mJy) | (mJy)
1 | 107.32592 | -10.82772 | … | $\phantom{1}1.33\pm 0.06$ | $\phantom{1}4.56\pm 0.13$ | $11.11\pm 0.43$ | $13.57\pm 0.30$ | $\phantom{1}8.59\pm 0.18$ | $\phantom{1}14.28\pm 0.27$
2 | 107.32975 | -10.82944 | $86.55\pm 2.07$ | $86.28\pm 2.07$ | $60.13\pm 1.27$ | $27.44\pm 1.05$ | $16.67\pm 0.57$ | $33.60\pm 0.72$ | $\phantom{1}20.02\pm 0.43$
3 | 107.33013 | -10.83094 | … | $\phantom{1}2.33\pm 0.33$ | $\phantom{1}3.15\pm 0.26$ | $\phantom{1}2.61\pm 0.10$ | $\phantom{1}2.63\pm 0.07$ | … | …
4 | 107.33121 | -10.82658 | $\phantom{1}0.79\pm 0.07$ | $\phantom{1}1.33\pm 0.07$ | $\phantom{1}1.85\pm 0.12$ | $\phantom{1}1.89\pm 0.08$ | $\phantom{1}1.82\pm 0.05$ | … | …
5 | 107.33121 | -10.86256 | $\phantom{1}1.72\pm 0.06$ | $\phantom{1}2.86\pm 0.12$ | $\phantom{1}2.45\pm 0.13$ | $\phantom{1}1.67\pm 0.06$ | $\phantom{1}1.38\pm 0.04$ | … | …
6 | 107.33275 | -10.83064 | $\phantom{1}1.02\pm 0.06$ | $\phantom{1}1.92\pm 0.09$ | $\phantom{1}1.98\pm 0.12$ | $\phantom{1}1.31\pm 0.06$ | $\phantom{1}0.82\pm 0.04$ | … | …
7 | 107.33379 | -10.84222 | $\phantom{1}2.33\pm 0.24$ | $14.68\pm 1.47$ | $23.61\pm 2.36$ | … | … | … | …
8 | 107.33467 | -10.84192 | $\phantom{1}2.59\pm 0.35$ | $16.75\pm 1.67$ | $32.03\pm 3.20$ | … | … | … | …
9 | 107.33529 | -10.85725 | $\phantom{1}1.55\pm 0.08$ | $\phantom{1}4.01\pm 0.11$ | $\phantom{1}5.40\pm 0.19$ | $\phantom{1}5.79\pm 0.21$ | $\phantom{1}5.58\pm 0.16$ | $\phantom{1}5.16\pm 0.12$ | $\phantom{11}5.31\pm 0.11$
10 | 107.33563 | -10.84117 | $\phantom{1}1.17\pm 0.16$ | $\phantom{1}4.57\pm 0.65$ | $21.06\pm 1.46$ | … | … | $87.24\pm 1.79$ | $403.08\pm 7.87$
11 | 107.33771 | -10.83483 | $\phantom{1}1.41\pm 0.07$ | $\phantom{1}2.02\pm 0.07$ | $\phantom{1}2.03\pm 0.10$ | $\phantom{1}1.98\pm 0.07$ | $\phantom{1}2.02\pm 0.07$ | … | …
12 | 107.33879 | -10.84178 | $\phantom{1}1.47\pm 0.15$ | $\phantom{1}1.16\pm 0.20$ | $\phantom{1}3.22\pm 0.30$ | … | … | $24.38\pm 2.28$ | $\phantom{1}52.26\pm 1.41$
13 | 107.34054 | -10.85011 | … | $\phantom{1}1.08\pm 0.07$ | $\phantom{1}3.15\pm 0.10$ | $\phantom{1}4.89\pm 0.21$ | $\phantom{1}5.30\pm 0.11$ | $\phantom{1}4.91\pm 0.10$ | $\phantom{11}7.53\pm 0.10$
14 | 107.34071 | -10.85858 | $\phantom{1}0.35\pm 0.05$ | $\phantom{1}1.48\pm 0.09$ | $\phantom{1}2.82\pm 0.13$ | $\phantom{1}3.25\pm 0.12$ | $\phantom{1}3.62\pm 0.14$ | $\phantom{1}2.93\pm 0.06$ | $\phantom{11}3.38\pm 0.07$
15 | 107.34100 | -10.83386 | $\phantom{1}8.58\pm 0.29$ | $21.79\pm 0.60$ | $23.65\pm 0.50$ | $12.97\pm 0.61$ | $\phantom{1}9.56\pm 0.36$ | $14.79\pm 0.27$ | $\phantom{1}11.00\pm 0.18$
ID | RA | Dec | $S_{\rm W3}$ | $S_{\rm W4}$ | $S_{70}$ | $S_{160}$ | $S_{250}$ | $S_{350}$ | $S_{500}$
| (deg) | (deg) | (mJy) | (mJy) | (Jy) | (Jy) | (Jy) | (Jy) | (Jy)
1 | 107.32592 | -10.82772 | $\phantom{11}19.26\pm 0.41\phantom{1}$ | $\phantom{11}79.72\pm 2.92\phantom{1}$ | … | … | … | … | …
2 | 107.32975 | -10.82944 | $\phantom{111}7.16\pm 0.28\phantom{1}$ | $\phantom{11}23.74\pm 1.56\phantom{1}$ | … | … | … | … | …
3 | 107.33013 | -10.83094 | … | … | … | … | … | … | …
4 | 107.33121 | -10.82658 | … | … | … | … | … | … | …
5 | 107.33121 | -10.86256 | … | … | … | … | … | … | …
6 | 107.33275 | -10.83064 | … | … | … | … | … | … | …
7 | 107.33379 | -10.84222 | … | … | … | … | … | … | …
8 | 107.33467 | -10.84192 | … | … | … | … | … | … | …
9 | 107.33529 | -10.85725 | $\phantom{111}4.18\pm 0.19\phantom{1}$ | $\phantom{11}66.19\pm 2.99\phantom{1}$ | … | … | … | … | …
10 | 107.33563 | -10.84117 | $4017.01\pm 33.44$ | $28890.97\pm 160.10$ | $165.8\pm 0.5$ | $1030.6\pm 0.4$ | $57.0\pm 0.5$ | $23.5\pm 0.1$ | $10.6\pm 0.1$
11 | 107.33771 | -10.83483 | … | … | … | … | … | … | …
12 | 107.33879 | -10.84178 | $117.25\pm 2.62$ | $4541.11\pm 46.24$ | $\phantom{1}42.5\pm 0.2$ | $\phantom{11}60.8\pm 0.2$ | $62.0\pm 0.6$ | $20.8\pm 0.1$ | …
13 | 107.34054 | -10.85011 | $\phantom{111}8.38\pm 0.17\phantom{1}$ | $144.66\pm 1.34$ | … | … | … | … | …
14 | 107.34071 | -10.85858 | $\phantom{111}3.38\pm 0.20\phantom{1}$ | $\phantom{11}66.43\pm 2.62\phantom{1}$ | … | … | … | … | …
15 | 107.34100 | -10.83386 | $\phantom{11}11.04\pm 0.14\phantom{1}$ | $743.98\pm 1.37$ | … | … | … | … | …

##  Appendix D CO rotational temperature of the intermediate-to high-mass YSOs
in the Milky Way

We compare the rotational temperatures towards the two dense cores to the
results found in the samples of the IM and HM YSOs in the Milky Way. There are
six IM and ten HM YSOs presented in Karska et al. (2014a) and Matuszak et al.
(2015), respectively using the data from Herschel/PACS. To be consistent with
our SOFIA/FIFI-LS data, we recalculated the rotational temperatures derived by
fitting the rotational diagram using 4 CO transitions with $J_{\mathrm{up}}$ =
(14, 16, 18, 22) and (14, 16, 17, 22) for the IM and HM YSOs, respectively.

We checked the new results with the ones presented in the aforementioned
studies and find that using only four CO transitions returns a lower
$T\rm{{}_{rot}}$ than the case of using all the observable CO transitions. The
relative difference is in between 4-27$\%$ with a median of 5$\%$ for the HM
YSOs. The comparison of these differences is shown in Table 8.

Table 9 shows heliocentric distances, Galactocentric radii, metallicities, and
luminosities of the sources in the Milky Way and Magellanic Clouds. We used
heliocentric distances together with source coordinates to calculate
Galactocentric radii, assuming a distance from the Sun to the Galactic center
of 8.34 kpc. We estimated the metallicity toward the Milky Way sources using
the O/H galactocentric radial gradient based on H ii regions in the Galactic
disk (Balser et al. 2011).

Table 8: CO rotational temperature and number of emitting molecules for HM
YSOs, derived by fitting the rotational diagram of CO using all and only four
available CO transitions.

Source | $T\mathrm{{}_{rot,4}}^{a}$ | $T\mathrm{{}_{rot,all}}^{b}$ |  $T\mathrm{{}_{rot}}$ Diff.c
---|---|---|---
| (K) | (K) | ($\%$)
G327-0.6 | $275\pm 60$ | $265\pm 30$ | 4
W51N-e1 | $270\pm 15$ | $260\pm 15$ | 4
DR21OH | $225\pm 25$ | $235\pm 15$ | 4
W33A | $135\pm 5$ | $185\pm 25$ | 27
G34.26+0.15 | $330\pm 60$ | $320\pm 25$ | 3
NGC6334-I | $320\pm 55$ | $300\pm 25$ | 7
NGC7538-I1 | $175\pm 25$ | $185\pm 15$ | 5
AFGL2591 | $175\pm 10$ | $195\pm 15$ | 10
W3IRS5 | $335\pm 40$ | $320\pm 15$ | 5
G5.89-0.39 | $345\pm 40$ | $285\pm 15$ | 21
999 (a) Rotational temperatures calculated using 4 CO transitions with
$J_{\rm{up}}=(14,16,17,22)$. (b) Rotational temperatures calculated using all
detected CO transitions, data is adopted from Karska et al. (2014a). (c) The
relative difference of $T\rm{{}_{rot}}$ is taken as
$|1-T_{\rm{rot,4}}/T_{\rm{rot,all}}|\times 100$.

Table 9: FIR line cooling luminosity of the sample in the Milky Way (MW), LMC,
and SMC used to compare with results of this study

Source | $D$ |  $D$ Ref. | $R{\rm{}_{GC}}^{a}$ | $Z^{b}$ |  log($L\mathrm{{}_{bol}}$/$L\mathrm{{}_{\odot}}$) |  $L\mathrm{{}_{bol}}$ Ref. | $L\mathrm{{}_{CO}}^{c}$ |  $L\mathrm{{}_{[{O\textsc{i}}]}}^{d}$ |  $L\mathrm{{}_{CO}}$/$L\mathrm{{}_{[{O\textsc{i}}]}}$ |  $L\mathrm{{}_{lines}}$ Ref.
---|---|---|---|---|---|---|---|---|---|---
| (kpc) |  | (kpc) |  ($Z\mathrm{{}_{\odot}}$) |  |  |  (10${}^{-2}L\mathrm{{}_{\odot}}$) |  (10${}^{-2}L\mathrm{{}_{\odot}}$) |  |
LMC YSOs |
IRAS04514$-$6931 |  50.0 $\pm$ 1.1 | 1 | — | 0.4 | 4.8 | 16 | 5.3 | 8.8 | 0.605 | 16
N113 YSO3 |  50.0 $\pm$ 1.1 | 1 | — | 0.4 | 5.4 | 16 | 29.5 | 57.8 | 0.511 | 16
SAGE045400.9$-$691151.6 |  50.0 $\pm$ 1.1 | 1 | — | 0.4 | 5.1 | 16 | 14.6 | 29.2 | 0.5 | 16
SAGE051351.5$-$672721.9 |  50.0 $\pm$ 1.1 | 1 | — | 0.4 | 5.1 | 16 | 5.5 | 15.4 | 0.357 | 16
SAGE052202.7$-$674702.1 |  50.0 $\pm$ 1.1 | 1 | — | 0.4 | 4.5 | 16 | 4.3 | 9.4 | 0.459 | 16
SAGE052212.6$-$675832.4 |  50.0 $\pm$ 1.1 | 1 | — | 0.4 | 5.5 | 16 | 14.9 | 25.4 | 0.585 | 16
SAGE053054.2$-$683428.3 |  50.0 $\pm$ 1.1 | 1 | — | 0.4 | 4.9 | 16 | 6.8 | 5.4 | 1.24 | 16
ST01 |  50.0 $\pm$ 1.1 | 1 | — | 0.4 | 4.6 | 16 | 7 .1 | 8.5 | 0.833 | 16
N113 YSO1 |  50.0 $\pm$ 1.1 | 1 | — | 0.4 | 5.4 | 16 | 47.5 | 47.5 | 1 | 16
SMC YSOs |
IRAS00464$-$7322 |  62.1 $\pm$ 2.0 | 2 | — | 0.2 | 4.1 | 16 | 3.0 | 2.8 | 1.077 | 16
IRAS00430$-$7326 |  62.1 $\pm$ 2.0 | 2 | — | 0.2 | 4.9 | 16 | 1.9 | 10.9 | 0.176 | 16
N81 |  62.1 $\pm$ 2.0 | 2 | — | 0.2 | 4.7 | 16 | 2\. 4 | 36.4 | 0.067 | 16
SMC012407-730904 |  62.1 $\pm$ 2.0 | 2 | — | 0.2 | 5.3 | 16 | 6.8 | 71.2 | 0.096 | 16
MW IM YSOs |
AFGL 490 |  0.97 $\pm$ 0.40 | 3 | 9.12 | 0.7 | 3.7 | 17 |  24.7 $\pm$ 2.5 |  36.6 $\pm$ 0.2 |  0.7 $\pm$ 0.1 | 32
L 1641 |  0.428 $\pm$ 0.010 | 4 | 8.69 | 0.8 | 1.8 | 18 |  2.7 $\pm$ 0.4 |  1.6 $\pm$ 0.1 |  1.7 $\pm$ 0.3 | 32
NGC 2071 |  0.422 $\pm$ 0.050 | 4 | 8.71 | 0.8 | 2.7 | 19 |  44.8 $\pm$ 4.3 |  33.4 $\pm$ 0.1 |  1.3 $\pm$ 0.1 | 32
Vela 17 | 0.7 | 5 | 8.44 | 0.8 | 2.9 | 20 |  13.2 $\ pm$ 2.0 |  44.8 $\pm$ 0.1 |  0.3 $\pm$ 0.1 | 32
Vela 19 | 0.7 | 5 | 8.45 | 0.8 | 2.9 | 5 |  9.7 $\pm$ 2.0 |  14.4 $\pm$ 0.1 |  0.7 $\pm$ 0.1 | 32
NGC 7129 | 1.25 | 6 | 8.75 | 0.8 | 2.6 | 21 |  20.9 $\pm$ 3.0 |  6.9 $\pm$ 0.1 |  3.0 $\pm$ 0.4 | 32
MW HM YSOs |
G327-0.6 | 3.3 | 7 | 5.84 | 1.0 | 4.7 | 22 |  180 $\pm$ 60 |  30 $\pm$ 10 |  6.0 $\pm$ 2.8 | 33
W51N-e1 |  5.1 $\pm$ 2.2 | 8 | 12.42 | 0.5 | 5.0 | 23 |  2530 $\pm$ 680 |  760 $\pm$ 180 |  3.3 $\pm$ 1.2 | 33
DR21(OH) |  1.50 $\pm$ 0.08 | 9 | 8.26 | 0.8 | 4.1 | 24 |  120 $\pm$ 30 |  9 $\pm$ 3 |  13.3 $\pm$ 5.6 | 33
W33A |  2.40 $\pm$ 0.16 | 10 | 6.02 | 1.0 | 4.6 | 25 |  60 $\pm$ 20 |  5 $\pm$ 2 |  12.0 $\pm$ 6.2 | 33
G34.26+0.15 | 3.30 | 11 | 5.91 | 1.0 | 5.5 | 26 |  770 $\pm$ 240 |  190 $\pm$ 50 |  4.1 $\pm$ 1.7 | 33
NGC6334I |  1.7 $\pm$ 0.3 | 12 | 6.66 | 0.9 | 5.4 | 27 |  340 $\pm$ 100 |  4 $\pm$ 3 |  85.0 $\pm$ 68.5 | 33
NGC7538-IRS1 |  2.7 $\pm$ 0.1 | 13 | 9.66 | 0.7 | 5.1 | 28 |  280 $\pm$ 50 |  1100 $\pm$ 250 |  0.2 $\pm$ 0.1 | 33
AFGL2591 |  3.3 $\pm$ 0.08 | 9 | 8.36 | 0.9 | 5.3 | 29 |  160 $\pm$ 50 |  340 $\pm$ 90 |  0.5 $\pm$ 0.2 | 33
W3-IRS5 |  2.00 $\pm$ 0.07 | 14 | 9.83 | 0.7 | 5.2 | 30 |  1050 $\pm$ 250 |  420 $\pm$ 10 |  2.5 $\pm$ 0.8 | 33
G5.89-0.39 |  1.30 $\pm$ 0.09 | 15 | 7.05 | 0.9 | 4.7 | 31 |  390 $\pm$ 90 |  370 $\pm$ 90 |  1.1 $\pm$ 0.4 | 33

References: (1) Pietrzyński et al. (2013), (2) Graczyk et al. (2014)), (3)
Purser et al. (2021), (4) Wilson et al. (2005), (5) Liseau et al. (1992),(6)
Shevchenko & Yakubov (1989), (7) Minier et al. (2009), (8) Xu et al. (2009),
(9) Rygl et al. (2012) , (10) Immer et al. (2013), (11) Kuchar & Bania (1994),
(12) Neckel (1978), (13) Moscadelli et al. (2009), (14) Hachisuka et al.
(2006), (15) Motogi et al. (2011), (16) Oliveira et al. (2019),(17) Navarete
et al. (2015), (18) Stanke et al. (2000) , (19) Butner et al. (1990), (20)
Giannini et al. (2005), (21)Johnstone et al. (2010), (22) Urquhart et al.
(2012), (23) van Dishoeck et al. (2011),(24) Jakob et al. (2007), (25) Faúndez
et al. (2004), (26) Hatchell & van der Tak (2003), (27) Sandell (2000), (28)
Sandell & Sievers (2004), (29) van der Tak et al. (1999), (30) Ladd et al.
(1993), (31) van der Tak (2012), (32) Matuszak et al. (2015), (33) Karska et
al. (2014a).

a Galactocentric radius toward the sources in the Milky Way is calculated
using the information from their coordinate and distance to the Sun, assuming
the distance from the Sun to the Galactic center is 8.34 kpc.

b Metallicity towards the sources in the Milky Way is calculated using the
relation of O/H gradients inferred from the H ii regions:
$12+\log\rm{(O/H)}=-0.0446R_{\rm GC}+8.962$ (Balser et al. 2011).

c Total FIR luminosity of CO derived from the rotational diagram fitting.

d Total FIR luminosity of [O i] calculated by summing the luminosity of [O i]
lines at 63 and 145 $\mu$m.