11institutetext: Academy of Finland, Hakaniemenranta 6, P.O. Box 131, FI-00531
Helsinki, Finland
11email<EMAIL_ADDRESS>22institutetext: Laboratoire d’Astrophysique
(AIM), CEA, CNRS, Université Paris-Saclay, Université Paris Diderot, Sorbonne
Paris Cité, 91191 Gif-sur-Yvette, France

# ArTéMiS imaging of the filamentary infrared dark clouds G1.75-0.08 and
G11.36+0.80: Dust-based physical properties of the clouds and their
clumps††thanks: This publication is based on data acquired with the Atacama
Pathfinder EXperiment (APEX) under programme ID 108.21W4. APEX is a
collaboration between the Max-Planck-Institut für Radioastronomie, the
European Southern Observatory, and the Onsala Space Observatory. This work is
also based on observations carried out under project number 015-21 with the
IRAM 30-metre telescope. IRAM is supported by INSU/CNRS (France), MPG
(Germany) and IGN (Spain).

O. Miettinen 11 M. Mattern 22 and Ph. André 22

(Received ; accepted)

###### Abstract

Context. Filamentary infrared dark clouds (IRDCs) are a useful class of
interstellar clouds for studying the cloud fragmentation mechanisms on
different spatial scales. Determination of the physical properties of the
substructures in IRDCs can also provide useful constraints on the initial
conditions and early stages of star formation, including those of high-mass
stars.

Aims. We aim to determine the physical characteristics of two filamentary
IRDCs, G1.75-0.08 and G11.36+0.80, and their clumps. We also attempt to
understand how the IRDCs are fragmented into clumps.

Methods. We imaged the target IRDCs at 350 $\mu$m and 450 $\mu$m using the
bolometer called Architectures de bolomètres pour des Télescopes à grand champ
de vue dans le domaine sub-Millimétrique au Sol (ArTéMiS). These data were
used in conjunction with our previous 870 $\mu$m observations with the Large
APEX BOlometer CAmera (LABOCA) and archival Spitzer and Herschel data. The
LABOCA clump positions in G11.36+0.80 were also observed in the
N2H${}^{+}(1-0)$ transition with the Institut de Radioastronomie Millimétrique
(IRAM) 30-metre telescope.

Results. On the basis of their far-IR to submillimetre spectral energy
distributions (SEDs), G1.75-0.08 was found to be composed of two cold ($\sim
14.5$ K), massive (several $\sim 10^{3}$ M☉) clumps that are projectively
separated by $\sim 3.7$ pc. Both clumps are 70 $\mu$m dark, but they do not
appear to be bounded by self-gravity. The G1.75-0.08 filament was found to be
subcritical by a factor of $\sim 14$ with respect to its critical line mass,
but the result is subject to uncertain gas velocity dispersion. The IRDC
G11.36+0.80 was found to be moderately (by a factor of $\sim 2$) supercritical
and composed of four clumps that are detected at all wavelengths observed with
the ground-based bolometers. The SED-based dust temperatures of the clumps are
$\sim 13-15$ K, and their masses are in the range $\sim 232-633$ M☉. All the
clumps are gravitationally bound and they appear to be in somewhat different
stages of evolution on the basis of their luminosity-to-mass ratio. The
projected, average separation of the clumps is $\sim 1$ pc. At least three
clumps in our sample show hints of fragmentation into smaller objects in the
ArTéMiS images.

Conclusions. A configuration that is observed in G1.75-0.08, namely two clumps
at the ends of the filament, could be the result of gravitational focussing
acting along the cloud. The two clumps fulfil the mass-radius threshold for
high-mass star formation, but if their single-dish-based high velocity
dispersion is confirmed, their gravitational potential energy would be
strongly overcome by the internal kinetic energy, and the clumps would have to
be confined by external pressure to survive. Owing to the location of
G1.75-0.08 near the Galactic centre ($\sim 270$ pc), environmental effects
such as a high level of turbulence, tidal forces, and shearing motions could
affect the cloud dynamics. The observed clump separation in G11.36+0.80 can be
understood in terms of a sausage instability, which conforms to the findings
in some other IRDC filaments. The G11.36+0.80 clumps do not lie above the
mass-radius threshold where high-mass star formation is expected to be
possible, and hence lower-mass star formation seems more likely. The
substructure observed in one of the clumps in G11.36+0.80 suggests that the
IRDC has fragmented in a hierarchical fashion with a scale-dependent physical
mechanism. This conforms to the filamentary paradigm for Galactic star
formation.

###### Key Words.:

ISM: clouds – ISM: individual objects: G1.75-0.08 and G11.36+0.80 – Infrared:
ISM – Submillimetre: ISM

##  1 Introduction

Dust continuum imaging in the (sub-)millimetre regime ($\lambda=350-1\,000$
$\mu$m) is a very useful observational approach for characterising the
physical properties of interstellar molecular clouds and their substructures,
where the formation of new stars can take place. For the studies of the so-
called infrared dark clouds (IRDCs), dust continuum observations are of
particular importance because although the clouds are typically selected as
dark absorption features against the Galactic mid-IR background (Pérault et
al. (1996); Egan et al. (1998); Simon et al. (2006); Peretto & Fuller (2009)),
detection of dust emission from the cloud can distinguish it as a true,
physical IRDC from a local mimimum in the Galactic mid-IR background that
might only resemble an IRDC (Wilcock et al. (2012)).

Infrared dark clouds themselves are very useful target sources for the studies
of molecular cloud fragmentation and star formation. One reason for this is
that IRDCs are at an early stage of evolution (e.g. Rathborne et al. (2006)),
and hence their initial fragmentation processes can be studied before any
strong feedback from the newly formed stars has affected the cloud. Second,
IRDCs often exhibit filamentary morphology with substructures of elevated
density along the long axis of the filament (e.g. Jackson et al. (2010);
Kainulainen et al. (2013); Henshaw et al. 2016b ; Miettinen (2018)). Among the
different fragmentation mechanisms that can operate in interstellar clouds,
the so-called sausage instability with a particular length scale has been able
to successfully explain the observed fragment separations in many IRDC
filaments (e.g. the Nessie nebula (Jackson et al. (2010)), G14.225-0.506
(Busquet et al. (2013)), the Snake IRDC (Kainulainen et al. (2013)),
G34.43+00.24 (Foster et al. (2014)), and the Seahorse IRDC (Miettinen
(2018))). Moreover, some filamentary IRDCs are found to exhibit a hierarchical
fragmentation behaviour with spatial-scale-dependent fragmentation mechanism
(e.g. Kainulainen et al. (2013); Wang et al. (2014); Mattern et al. 2018a ;
Liu et al. (2021)). Similarly, other Galactic filamentary molecular clouds
such as those in Orion (e.g. Takahashi et al. (2013); Teixeira et al. (2016))
and NGC 6334 (e.g. Shimajiri et al. (2019)) are found to be fragmented in a
hierarchical fashion, which suggests that the fragmentation of the filament
proceeds at least in a bimodal way (from filaments to clumps, and from clumps
to cores), and that this is a qualitatively similar or quasi-similar
characteristic over a wide range of filament masses and densities (see André
et al. (2014) for a review; André et al. (2019)).

A third reason for IRDCs being useful target sources is that some of the IRDC
population members show evidence of ongoing high-mass ($M>8$ M☉) star
formation (e.g. Rathborne et al. (2006); Beuther & Steinacker (2007); Chambers
et al. (2009); Battersby et al. (2010)). In this process, many details and
even the exact mechanism(s) are still to be deciphered (e.g. Motte et al.
(2018) for a review).

In this paper, we present the results of our dual-band (350 $\mu$m and 450
$\mu$m) submillimetre dust continuum imaging of two filamentary IRDCs, namely
G1.75-0.08 and G11.36+0.80. Selected positions in G11.36+0.80 were also
observed in the $J=1-0$ transition of N2H+, which provides us with information
about the gas kinematics in the cloud. The target IRDCs were already imaged
with the Large APEX BOlometer CAmera (LABOCA; Siringo et al. (2009)) at 870
$\mu$m by Miettinen (2012), and they were part of the sample in the molecular
line study of IRDCs by Miettinen (2014) that was based on the Millimetre
Astronomy Legacy Team 90 GHz (MALT90) survey (Foster et al. (2011), 2013;
Jackson et al. (2013)). The angular resolution of the new dust continuum data
is about twice better than that of our previous LABOCA data, and the new
wavelength bands allow us to construct the spectral energy distributions
(SEDs) of the detected sources and hence determine their physical properties
more accurately than previously.

The observations, data reduction, and ancillary data are described in Sect. 2.
The analysis and results are described in Sect. 3. We discuss the results in
Sect. 4, while in Sect. 5, we summarise our results and main conclusions.

##  2 Observations, data reduction, and ancillary data

###  2.1 ArTéMiS submillimetre continuum imaging

Target fields with sizes of $8\aas@@fstack{\prime}13\times
5\aas@@fstack{\prime}05$ and $6\aas@@fstack{\prime}37\times
7\aas@@fstack{\prime}08$ and centred on the IRDCs G1.75-0.08 and G11.36+0.80,
respectively, were imaged with the bolometer called Architectures de
bolomètres pour des Télescopes à grand champ de vue dans le domaine sub-
Millimétrique au Sol (ArTéMiS; see Revéret et al. (2014); André et al. (2016);
Talvard et al. (2018)) on the Atacama Pathfinder EXperiment (APEX; Güsten et
al. (2006)) telescope. The ArTéMiS bolometer array consists of 2 300
individual bolometers, and it currently operates simultaneously at two
wavebands, namely 350 $\mu$m and 450 $\mu$m. The angular resolutions (half-
power beam width, or HPBW) at these wavelengths are
$8\aas@@fstack{\prime\prime}5$ and $9\aas@@fstack{\prime\prime}4$,
respectively. The field of view (FoV) of the array is
$4\aas@@fstack{\prime}7\times 2\aas@@fstack{\prime}5$ at both wavelengths.

Our observations were carried out on 21–23 August 2021 under very good weather
conditions with a typical precipitable water vapour (PWV) of 0.4 mm. The
zenith opacity at the ArTéMiS frequencies was measured with skydip
observations and was found to lie in the range $\tau_{\rm z}=0.26-0.78$. The
observations were performed using the total power on-the-fly (OTF) mapping
mode with the dump and grid spacings (i.e. along and perpendicular to the
scanning direction) set to $8\arcsec$ and $4\arcsec$, respectively. The
integration time per dump was one second. We note that ArTéMiS uses filled
bolometer arrays that essentially fully sample the aforementioned FoV at any
given time. Scanning was performed in only one direction, namely along the
right ascension axis. This can lead to the beam smearing effects and artefacts
in the maps (see Sect. 4.5 for further discussion). Owing to the need of
subtracting correlated sky noise, the spatial filtering scale for the ArTéMiS
data is typically $\sim 2\arcmin$, which is comparable to the smaller diameter
of the FoV (André et al. (2016)).

The telescope focus and pointing were optimised and checked at regular
intervals on the moon Ganymede, the planets Uranus and Neptune, and multiple
different massive star-forming objects (G5.89-0.39, G10.47+0.03, G10.62-0.38,
G34.26+0.15, and G45.07+0.13). The absolute flux calibration uncertainty is
estimated to be $\sim 30\%$ at 350 $\mu$m and $\sim 20\%$ at 450 $\mu$m. Five
hours of telescope time (including overheads) were spent on G1.75-0.08, while
5.8 hr was used for observations towards G11.36+0.80. Hence, the total
telescope time used for the project was 10.8 hr.

The data were reduced using a dedicated data reduction package in IDL
(Interactive Data Language). The IDL packages are publicly available on the
APEX website111www.apex-telescope.org/ns/data-reduction-for-bolometer-
arrays/#ArtemisPipeline.. The $1\sigma$ rms noise levels in the final 350
$\mu$m and 450 $\mu$m maps of G1.75-0.08 were determined to be 0.88 Jy beam-1
and 0.53 Jy beam-1, respectively, while those for G11.36+0.80 are 0.43 Jy
beam-1 and 0.29 Jy beam-1. The ArTéMiS contour maps of the target IRDCs are
shown in Figs. 1 and 2 along with multi-wavelength images of the IRDCs.

Figure 1: Multi-wavelength images towards G1.75-0.08. In the top panels, the
colour images from left to right show the Spitzer/IRAC 8 $\mu$m, Spitzer/MIPS
24 $\mu$m, and Herschel/PACS 70 $\mu$m maps. The overlaid contours show the
ArTéMiS 350 $\mu$m, ArTéMiS 450 $\mu$m, and LABOCA 870 $\mu$m emission,
respectively. In each case, the contour levels start from $3\sigma$ and
progress in steps of $1\sigma$. In the bottom panels, the colour images from
left to right show the Herschel/PACS 160 $\mu$m, Herschel/SPIRE 350 $\mu$m,
and Herschel/SPIRE 500 $\mu$m maps. The first two are overlaid with the
ArTéMiS 350 $\mu$m contours, while the rightmost map is overlaid with the
ArTéMiS 450 $\mu$m contours. The beam sizes (HPBW) of the APEX bolometer
observations are shown in the bottom left corner of each panel. A scale bar of
1 pc is shown in the right panels. Figure 2: Same as Fig. 1, but towards
G11.36+0.80. All the contour levels start from $3\sigma$ and progress in steps
of $3\sigma$. A source from the Feng et al. (2020) sample, G11.38+0.81 P3, is
labelled P3 in the Spitzer 24 $\mu$m plus ArTéMiS 450 $\mu$m image. The 24
$\mu$m dark features to the south of clump B are also indicated in this image.
We call them subfilaments.

###  2.2 IRAM 30-metre telescope observations

The LABOCA 870 $\mu$m peak positions along the G11.36+0.80 filament, which we
call clumps A–D (Fig. 2), were observed in the $J=1-0$ transition of N2H+ with
the Institut de Radioastronomie Millimétrique (IRAM) 30-metre telescope. The
observations were carried out on 17 October 2021 during a pool observing week.

As a front end, we used the Eight MIxer Receiver (EMIR; Carter et al. (2012))
band E090. The backend was the Versatile SPectrometer Array (VESPA), where the
bandwidth was $18\times 20$ MHz with a corresponding channel spacing of 20
kHz. The N2H${}^{+}(1-0)$ line was tuned at 93 173.7637 MHz, which corresponds
to the frequency of the strongest hyperfine component
($J_{F_{1},\,F}=1_{2,\,3}-0_{1,\,2}$) of the transition (Pagani et al.
(2009)). The aforementioned channel spacing yielded a velocity resolution of
about 64 m s-1. The telescope beam size (HPBW) at the observed frequency is
$26\aas@@fstack{\prime\prime}4$.

The observations were performed in the frequency-switching mode with a
frequency throw of $\pm 3.9$ MHz. Clumps A and B were observed for 22.5 min
each, while clumps C and D were observed for 28.1 min. The telescope focus and
pointing were optimised and checked on the planet Venus and G5.89-0.39. The
PWV during the observations was measured to be between 6.8 mm and 8.6 mm. The
system temperatures during the observations were within the range of $T_{\rm
sys}=103-108$ K. The observed intensities were converted into the main-beam
brightness temperature scale by using a main-beam efficiency factor of
$\eta_{\rm MB}=0.80$.

The spectra were reduced using the Continuum and Line Analysis Single-dish
Software (CLASS90) of the GILDAS software package (version mar19b)222Grenoble
Image and Line Data Analysis Software (GILDAS) is provided and actively
developed by IRAM, and is available at www.iram.fr/IRAMFR/GILDAS/.. The
individual spectra were averaged with weights $w_{i}\propto t_{\rm int}/T_{\rm
sys}^{2}$, where $t_{\rm int}$ is the on-source integration time. Linear
baselines were determined from the velocity ranges free of spectral line
features, and then subtracted from the spectra. The resulting $1\sigma$ rms
noise levels were about 23–32 mK.

The reduced spectra are shown in Fig. 3. The $J=1-0$ transition of N2H+ is
split into 15 hyperfine components, where the aforementioned strongest
component has a relative intensity of $R_{i}=7/15$. The hyperfine components
are blended into three groups in the observed spectra. We fitted the hyperfine
structure in CLASS90 using the rest frequencies from Pagani et al. (2009,
Table 2 therein) and the relative intensities from Mangum & Shirley (2015,
Table 7 therein). The derived spectral line parameters are given in Table 1.
The listed parameters are the local standard of rest (LSR) radial velocity
($v_{\rm LSR}$), full width at half maximum (FWHM; $\Delta v$), peak intensity
($T_{\rm MB}$), and the total optical thickness of the line ($\tau$).

Figure 3: N2H${}^{+}(1-0)$ spectra towards the LABOCA clumps in G11.36+0.80.
The intensity scale is shown in units of the main-beam brightness temperature.
Table 1: Parameters of the N2H${}^{+}(1-0)$ spectral lines detected towards
the LABOCA clumps in G11.36+0.80.

Source | $v_{\rm LSR}$ | $\Delta v$ | $T_{\rm MB}$ | $\tau$
---|---|---|---|---
| [km s-1] | [km s-1] | [K] |
Clump A | $26.7\pm 0.006$ | $1.46\pm 0.02$ | $1.22\pm 0.09$ | $6.76\pm 0.27$
Clump B | $28.1\pm 0.003$ | $1.28\pm 0.01$ | $1.69\pm 0.07$ | $7.35\pm 0.18$
Clump C | $27.7\pm 0.002$ | $1.71\pm 0.08$ | $2.31\pm 0.08$ | $2.29\pm 0.06$
Clump D | $28.1\pm 0.002$ | $1.12\pm 0.02$ | $2.12\pm 0.04$ | $2.37\pm 0.02$

###  2.3 Ancillary data

We complemented our new ArTéMiS submillimetre continuum data with the earlier
LABOCA 870 $\mu$m imaging data of the target IRDCs from Miettinen (2012). We
note that G1.75-0.08 was part of a larger target field of Miettinen (2012)
called G1.87-0.14. In brief, the angular resolution of the LABOCA data is
$19\aas@@fstack{\prime\prime}86$ (HPBW), and the $1\sigma$ rms noise of the
G1.75-0.08 map is 90 mJy beam-1 and that for G11.36+0.80 is about 31 mJy
beam-1 around the IRDC in the centre of the map.

We also employed data from the Herschel Space Observatory (Pilbratt et al.
(2010))333Herschel is an ESA space observatory with science instruments
provided by European-led Principal Investigator consortia and with important
participation from NASA., which were measured with the Photodetector Array
Camera and Spectrometer (PACS; Poglitsch et al. (2010)) in the far-IR regime
(70–160 $\mu$m) with nominal angular resolutions of $\sim 6\arcsec-12\arcsec$,
and the Spectral and Photometric Imaging Receiver (SPIRE; Griffin et al.
(2010)) in the far-IR to submillimetre regime (250–500 $\mu$m) with nominal
angular resolutions of $18\arcsec-35\arcsec$. In particular, we made use of
the data products from the Herschel infrared Galactic Plane Survey (Hi-GAL;
Molinari et al. 2016a ). The PACS beams of the Hi-GAL observations are
elongated along the scan direction, and have sizes of
$5\aas@@fstack{\prime\prime}8\times 12\aas@@fstack{\prime\prime}1$ at 70
$\mu$m and $11\aas@@fstack{\prime\prime}4\times 13\aas@@fstack{\prime\prime}4$
at 160 $\mu$m.

Our study also makes use of the mid-IR observations obtained with the Spitzer
Space Telescope’s (Werner et al. (2004)) Infrared Array Camera (IRAC; Fazio et
al. (2004)) and the Multiband Imaging Photometer for Spitzer (MIPS; Rieke et
al. (2004)). We used the 8 $\mu$m images taken as part of the Galactic Legacy
Infrared Mid-Plane Survey Extraordinaire (GLIMPSE; Churchwell et al. (2009)),
and 24 $\mu$m images from the MIPS Galactic Plane Survey (MIPSGAL; Carey et
al. (2009)). Moreover, we employed the 24 $\mu$m Point Source Catalogue of the
Galactic Plane from Spitzer/MIPSGAL (Gutermuth & Heyer (2015)).

##  3 Analysis and results

###  3.1 Submillimetre dust emission characteristics

In Tables 2 and 3, we list the properties of the ArTéMiS and LABOCA clumps in
the IRDCs G1.75-0.08 and G11.36+0.80, respectively. The parameters given in
the tables are the coordinates of the dust emission peak position, signal-to-
noise ratio (S/N) of the peak position, peak surface brightness, flux density
of the source integrated within the $3\sigma$ level of the emission, the
effective radius, which is defined as $R_{\rm eff}=\sqrt{A_{\rm
proj,\,3\sigma}/\pi}$, where $A_{\rm proj,\,3\sigma}$ is the projected area
within the $3\sigma$ extension of the source, and the FWHM size of the source,
which was measured over the largest diameter of the half-maximum contour. The
quoted uncertainties of the surface brightness and flux density take the
absolute calibration uncertainties into account. The uncertainties in the
effective radii and FWHM sizes were propagated from the average of the $\pm$
uncertainties in the cloud’s kinematic distance (Table 5, column (2)). As seen
in Table 3, the non-deconvolved FWHM extension of clump D in G11.36+0.80 is
comparable to the beam size (HPBW) at each wavelength, and the corresponding
deconvolved sizes make the source a point source. We note that the
aforementioned method for determining the source flux density, that is, above
the $3\sigma$ level, resembles the bijection paradigm presented by Rosolowsky
et al. (2008) and the way how for example the clumpfind algorithm (Williams et
al. (1994)) works.

We also determined the integrated flux densities of the whole filamentary
structures of G1.75-0.08 and G11.36+0.80. The ArTéMiS and LABOCA flux
densities that were all determined inside a region that corresponds to the
LABOCA $3\sigma$ contour are listed in Table 4.

The Spitzer and Herschel counterparts of the clumps in G1.75-0.08 and
G11.36+0.80 were searched from the NASA/IPAC Infrared Science Archive
(IRSA)444irsa.ipac.caltech.edu.. More details and the corresponding flux
densities are given in Appendix A.

Table 2: Submillimetre dust continuum sources in G1.75-0.08.

Source | $\alpha_{2000.0}$ | $\delta_{2000.0}$ | S/N | $I$ | $S$ | $R_{\rm eff}$ | ${\rm FWHM}_{\rm maj}$
---|---|---|---|---|---|---|---
| [h:m:s] | [$\degr$:$\arcmin$:$\arcsec$] |  | [Jy beam-1] | [Jy] | [″] | [pc] | [″] | [pc]
ArTéMiS 350 $\mu$m
Clump A | 17 50 05.19 | -27 28 38.36 | 5.4 | $4.74\pm 1.67$ | $42.6\pm 13.2$ | 18.3 | $0.73\pm 0.02$ | 45.5 | $1.81\pm 0.04$
Clump B | 17 49 59.63 | -27 29 19.54 | 5.5 | $4.85\pm 1.70$ | $33.3\pm 10.6$ | 16.0 | $0.64\pm 0.01$ | 39.3 | $1.57\pm 0.03$
ArTéMiS 450 $\mu$m
Clump A | 17 50 05.78 | -27 28 30.75 | 6.0 | $3.16\pm 0.63$ | $25.3\pm 5.5$ | 19.3 | $0.77\pm 0.02$ | 40.9 | $1.63\pm 0.03$
Clump B | 17 49 59.50 | -27 29 20.75 | 6.8 | $3.60\pm 0.72$ | $29.7\pm 6.4$ | 20.4 | $0.81\pm 0.02$ | 36.9 | $1.47\pm 0.03$
LABOCA 870 $\mu$m
Clump A | 17 50 05.08 | -27 28 32.12 | 8.4 | $0.76\pm 0.63$ | $4.05\pm 0.60$ | 36.5 | $1.46\pm 0.03$ | 49.4 | $1.97\pm 0.04$
Clump B | 17 49 59.38 | -27 29 24.98 | 8.2 | $0.74\pm 0.63$ | $2.73\pm 0.50$ | 31.0 | $1.24\pm 0.03$ | 42.6 | $1.70\pm 0.04$
555The columns are as follows: (1) source name; (2) and (3) coordinates of the
peak surface brightness position; (4) S/N ratio of the source; (5) peak
surface brightness; (6) integrated flux density within the $3\sigma$ extension
of the source; (7) and (8) effective radius of the source in angular and
spatial units; (9) and (10) FWHM size of the source in angular and spatial
units. The FWHM size refers to the largest diameter of the half-maximum
contour, and the tabulated values are not deconvolved from the beam size.

Table 3: Submillimetre dust continuum sources in G11.36+0.80.

Source | $\alpha_{2000.0}$ | $\delta_{2000.0}$ | S/N | $I$ | $S$ | $R_{\rm eff}$ | ${\rm FWHM}_{\rm maj}$
---|---|---|---|---|---|---|---
| [h:m:s] | [$\degr$:$\arcmin$:$\arcsec$] |  | [Jy beam-1] | [Jy] | [″] | [pc] | [″] | [pc]
ArTéMiS 350 $\mu$m
Clump A | 18 07 37.22 | -18 41 09.92 | 7.4 | $3.20\pm 1.03$ | $13.60\pm 4.53$ | 13.6 | $0.21\pm 0.02$ | 31.0 | $0.49\pm 0.04$
Clump B | 18 07 36.30 | -18 42 29.91 | 11.8 | $5.06\pm 2.00$ | $29.45\pm 9.24$ | 17.4 | $0.27\pm 0.02$ | 33.6 | $0.53\pm 0.04$
Clump C | 18 07 35.39 | -18 43 41.53 | 11.8 | $5.09\pm 2.01$ | $17.46\pm 5.80$ | 14.1 | $0.22\pm 0.02$ | 18.1 | $0.28\pm 0.02$
Clump D | 18 07 36.68 | -18 44 24.96 | 13.2 | $5.67\pm 2.38$ | $8.15\pm 3.73$ | 9.16 | $0.14\pm 0.01$ | 11.9 | $0.19\pm 0.02$
ArTéMiS 450 $\mu$m
Clump A | 18 07 37.18 | -18 41 05.42 | 10.0 | $2.89\pm 0.65$ | $17.11\pm 3.73$ | 20.0 | $0.31\pm 0.03$ | 32.6 | $0.51\pm 0.04$
Clump B | 18 07 36.20 | -18 42 36.39 | 13.6 | $3.94\pm 0.84$ | $20.27\pm 4.52$ | 18.9 | $0.30\pm 0.03$ | 30.3 | $0.47\pm 0.04$
Clump C | 18 07 35.30 | -18 43 42.22 | 15.0 | $4.35\pm 0.92$ | $20.95\pm 4.50$ | 21.6 | $0.34\pm 0.03$ | 17.1 | $0.27\pm 0.02$
Clump D | 18 07 36.53 | -18 44 26.20 | 17.0 | $4.92\pm 1.03$ | $10.77\pm 2.87$ | 14.5 | $0.23\pm 0.02$ | 13.5 | $0.21\pm 0.02$
LABOCA 870 $\mu$m
Clump A | 18 07 36.69 | -18 41 13.51 | 19.9 | $0.62\pm 0.07$ | $2.39\pm 0.41$ | 37.8 | $0.59\pm 0.05$ | 51.8 | $0.81\pm 0.07$
Clump B | 18 07 35.86 | -18 42 41.75 | 24.0 | $0.75\pm 0.08$ | $2.07\pm 0.39$ | 31.9 | $0.50\pm 0.04$ | 41.2 | $0.65\pm 0.05$
Clump C | 18 07 35.04 | -18 43 45.75 | 19.9 | $0.62\pm 0.07$ | $2.58\pm 0.42$ | 35.3 | $0.55\pm 0.05$ | 49.1 | $0.77\pm 0.07$
Clump D | 18 07 36.14 | -18 44 25.74 | 19.2 | $0.60\pm 0.07$ | $1.20\pm 0.30$ | 26.2 | $0.41\pm 0.03$ | 25.8 | $0.40\pm 0.03$
666Columns are as in Table 2.

###  3.2 Kinematic distances of the target IRDCs

To determine the kinematic distances of the target IRDCs, we used the Monte
Carlo technique of Wenger et al. (2018). This technique makes use of the Reid
et al. (2014) rotation curve with updated solar motion parameters (see Table 2
in Wenger et al. (2018))777The online calculation tool is available at
www.treywenger.com/kd/..

For G11.36+0.80, we used the mean value of the LSR velocities of
N2H${}^{+}(1-0)$ listed in column (2) in Table 1, that is $27.65\pm 0.33$ km
s-1, where the uncertainty represents the standard error of the mean. For
G1.75-0.08, we do not have single-pointing spectral line observations
available, and hence we relied on the spectral line data presented in
Miettinen (2014), which were taken as part of the MALT90 survey with the 22 m
Mopra telescope in the OTF mapping mode. The HCN$(1-0)$ and HCO${}^{+}(1-0)$
lines were the strongest detections of the spectra extracted towards clump B
in G1.75-0.08 (source G1.87-0.14 SMM 8 in Miettinen (2014); see Fig. C.2
therein). Towards clump A, HCN$(1-0)$ and HCO${}^{+}(1-0)$ were also clearly
detected, but the latter showed a clear double-peaked profile (source
G1.87-0.14 SMM 15 in Miettinen (2014); see Fig. C.6 therein). The mean value
of the LSR velocities of HCN$(1-0)$ towards clump A and HCN$(1-0)$ and
HCO${}^{+}(1-0)$ towards clump B is $53.53\pm 1.05$ km s-1, and this value was
used in the kinematic distance calculation.

The derived distances are given in Table 5. The parameters listed in the table
are the near and far kinematic distance solutions and the Galactocentric
radius ($d_{\rm near}$, $d_{\rm far}$, and $R_{\rm GC}$, respectively). The
quoted uncertainties represent the 68.3% confidence interval. We assume that
the IRDC lies at the near kinematic distance, where it is more likely to be
seen in absorption against the Galactic mid-IR background radiation field.
However, we note that G1.75-0.08 lies so close to the Galactic Centre ($\sim
270$ pc) that the near and far distance solutions differ by only $\sim 4\%$.

###  3.3 SED analysis

The mid-IR or far-IR to submillimetre SEDs of the clumps in G1.75-0.08 and
G11.36+0.80 were fitted by a modified blackbody (MBB) function using the
method and assumptions of Miettinen (2020; Sect. 3.1 therein). For the 24
$\mu$m dark clumps, we used a single-temperature MBB model, while the SEDs of
24 $\mu$m bright clumps were fitted with a two-temperature MBB model. In
brief, the SED analysis assumed a Ossenkopf & Henning (1994) dust model for
grains with thin ice mantles, where, for reference, the dust opacity at 870
$\mu$m is 1.38 cm2 g-1. The dust-to-total gas mass ratio was assumed to be
1/141 (i.e. 1.41 times lower than a dust-to-hydrogen mass ratio of 1/100,
where the helium and metal mass contributions are taken to be 27% and 2%,
respectively).

The SEDs are shown in Figs. 4 and 5. We note that the flux densities at
$160-870$ $\mu$m were determined within apertures that corresponded to the
LABOCA $3\sigma$ contour of emission so that the flux densities can be better
compared to each other. However, the PACS 70 $\mu$m flux densities used in the
fits were taken from the Hi-GAL catalogue (see Appendix A) because those data
are related to 70 $\mu$m point sources. The SPIRE $250-500$ $\mu$m data were
not used in the fits because the 350 $\mu$m band is covered by our ArTéMiS
data, the 500 $\mu$m band is close to our ArTéMiS 450 $\mu$m data point, and
SPIRE is more sensitive to the extended dust emission than our ground-based
bolometric data, and because the present work focusses on the high density
parts of the filaments and their clumps. Nevertheless, as discussed in Sect.
4.2, the SPIRE flux densities are mostly consistent with the SED fits. In Fig.
6 we show the SEDs of the G1.75-0.08 and G11.36+0.80 filaments that were
constructed using the photometric data from Table 4 in conjunction with the
PACS 160 $\mu$m flux densities integrated within the LABOCA $3\sigma$ contour.
The 160 $\mu$m flux densities were determined as described in Appendix A, that
is, by determining the flux conversion factors on the basis of the Hi-GAL
catalogue values, which also take the background emission contribution into
account. We note that because our target IRDCs do not extend over a wide area
in projection (the projected extensions are $\sim 3\aas@@fstack{\prime}4$ for
G1.75-0.08 and $\sim 4\aas@@fstack{\prime}2$ for G11.36+0.80), the background
(and foreground) emission is expected to be fairly uniform across the clouds.

The derived SED parameters are listed in Table 6. Columns (2)–(7) in the table
list the number of photometric data points used in the SED fit, reduced
$\chi^{2}$ value of the fit, wavelength of the peak position of the fitted
SED, dust temperature, total (gas+dust) mass of the source, and the bolometric
luminosity integrated over the fitted SED curve. The quoted uncertainties were
derived from the flux density uncertainties.

Figure 4: Far-IR to submillimetre SEDs of the clumps in G1.75-0.08. The SPIRE
data points shown by square green symbols are plotted for comparison, and they
were not included in the fit (see Sect. 3.3 for details). The dashed black
lines represent the best MBB fits to the data.

Figure 5: Mid-IR to submillimetre SEDs of the clumps in G11.36+0.80. The
Spitzer/MIPS 24 $\mu$m data points are highlighted in red. For the two-
temperature MBB fits (clumps C and D), the dashed black line shows the sum of
the two components. The dashed blue and red lines show the SED fits to the
cold and warm component, respectively.

Figure 6: Far-IR to submillimetre SEDs of the G1.75-0.08 and G11.36+0.80
filaments. Table 4: ArTéMiS and LABOCA flux densities of the G1.75-0.08 and
G11.36+0.80 filaments.

IRDC | $S_{\rm 350\,\mu m}$ | $S_{\rm 450\,\mu m}$ | $S_{\rm 870\,\mu m}$
---|---|---|---
| [Jy] | [Jy] | [Jy]
G1.75-0.08 | $172.9\pm 51.9$ | $99.7\pm 20.0$ | $6.9\pm 0.8$
G11.36+0.80 | $123.3\pm 24.7$ | $99.5\pm 19.9$ | $8.2\pm 0.9$
888All the quoted flux densities were determined within the corresponding
LABOCA $3\sigma$ contour level.

Table 5: Kinematic distances of the target IRDCs.

IRDC | $d_{\rm near}$ | $d_{\rm far}$ | $R_{\rm GC}$
---|---|---|---
| [kpc] | [kpc] | [kpc]
G1.75-0.08 | $8.22^{+0.26}_{-0.09}$ | $8.55^{+0.12}_{-0.19}$ | $0.27^{+0.02}_{-0.02}$
G11.36+0.80 | $3.23^{+0.34}_{-0.21}$ | $12.88^{+0.50}_{-0.25}$ | $5.12^{+0.28}_{-0.28}$
999The quoted uncertainties represent the 68.3% confidence interval.

Table 6: SED parameters and physical properties of the sources.

Source | $n_{\rm SED}^{\rm fit}$ | $\chi_{\rm red}^{2}$ | $\lambda_{\rm peak}$ | $T_{\rm dust}$ | $M$ | $L_{\rm bol}$ | $N({\rm H_{2}})$ | $n({\rm H_{2}})$ | $l$ | $R_{\rm fil}$
---|---|---|---|---|---|---|---|---|---|---
|  |  | [$\mu$m] | [K] | [M☉] | [L☉] | [$10^{23}$ cm-2] | [$10^{3}$ cm-3] | [pc] | [pc]
G1.75-0.08
Filament | 4 | 4.49 | 200.0 | $15.0\pm 0.4$ | $8\,289\pm 1\,196$ | $6\,358^{+2\,189}_{-1\,749}$ | $0.3\pm 0.05$ | $9.3\pm 1.3$ | $8.2\pm 0.2$ | $0.70\pm 0.01$
Clump A | 4 | 3.55 | 214.3 | $14.3\pm 0.4$ | $5\,501\pm 935$ | $3\,140^{+1\,237}_{-965}$ | $1.7\pm 0.6$ | $6.1\pm 1.0$ | … | …
Clump B | 4 | 3.30 | 200.0 | $14.5\pm 0.5$ | $3\,770\pm 738$ | $2\,402^{+1\,107}_{-833}$ | $1.7\pm 0.6$ | $6.8\pm 1.3$ | … | …
G11.36+0.80
Filament | 4 | 2.31 | 200.0 | $15.2\pm 0.5$ | $1\,468\pm 199$ | $1\,232^{+432}_{-342}$ | $0.2\pm 0.03$ | $11.5\pm 1.6$ | $3.9\pm 0.3$ | $0.39\pm 0.03$
Clump A | 5 | 0.90 | 214.3 | $13.6\pm 0.1$ | $532\pm 56$ | $226^{+39}_{-34}$ | $1.3\pm 0.4$ | $8.9\pm 0.9$ | … | …
Clump B | 5 | 2.60 | 230.8 | $13.2\pm 0.2$ | $437\pm 66$ | $157^{+41}_{-36}$ | $2.3\pm 0.9$ | $12.0\pm 1.8$ | … | …
Clump Caa$a$The SED of the source was fitted with a two-temperature MBB model. The temperature and mass of the warm component are shown below the cold dust component values. | 6 | 1.41 | 230.8 | $12.8\pm 0.3$ | $633\pm 112$ | $344^{+88}_{-182}$ | $2.6\pm 1.0$ | $13.0\pm 2.3$ | … | …
| … | … | … | $46.5\pm 0.1$ | $0.19\pm 0.01$ | … | … | … | … | …
Clump Daa$a$The SED of the source was fitted with a two-temperature MBB model. The temperature and mass of the warm component are shown below the cold dust component values. | 6 | 1.87 | 200.0 | $14.8\pm 0.6$ | $232\pm 56$ | $263^{+117}_{-87}$ | $1.8\pm 0.8$ | $11.5\pm 2.8$ | … | …
| … | … | … | $43.1\pm 0.3$ | $0.17\pm 0.01$ | … | … | … | … | …
101010The parameters given in the table are the number of flux density data
points used in the SED fit, reduced $\chi^{2}$ value, wavelength of the peak
position of the fitted SED, dust temperature, total (gas+dust) mass,
bolometric luminosity, H2 column density (source-averaged value within the
LABOCA $3\sigma$ emission for the filaments, and for the clumps in their
ArTéMiS 350 $\mu$m peak position), volume-averaged H2 number density,
projected length of the filament, and the outer radius of the filament (see
Sect. 3.4).

###  3.4 H2 column and number densities

The H2 column densities of the clumps were calculated from the ArTéMiS 350
$\mu$m peak surfaces brightness and using the SED-based dust temperature of
the clump (see e.g. Eq. (3) in Miettinen et al. (2009)). The dust temperature
of the cold component was adopted for clumps C and D in G11.36+0.80. The mean
molecular weight per H2 molecule needed in the calculation was assumed to be
2.82, which results from the hydrogen mass percentage of 71% assumed in the
calculation of the dust-to-gas mass ratio (Sect. 3.3). The volume-averaged H2
number densities of the clumps were calculated assuming a spherical geometry
with the effective radius listed in column (8) in Tables 2 and 3. We note that
the assumption of a uniformly distributed mass can be too simplistic for a
source with a density distribution, and the central density can significantly
exceed the volume-averaged value (see Sect. 4.3.1).

We also calculated the H2 column and number densities of the target IRDC
filaments. The former was calculated by dividing the mass of the filament by
the area enclosed by the LABOCA $3\sigma$ contour and using the aforementioned
mean molecular weight value. To calculate the number densities of the
filaments, we assumed a cylindrical geometry, where the length of the filament
was taken to be the projected length of the long axis of the LABOCA emission
at $3\sigma$ level. These are about 8.2 pc for G1.75-0.08 and 3.9 pc for
G11.36+0.80. The outer radius of the filament was calculated as the effective
radius by assuming a rectangular geometry where the area is equal to that
enclosed by the LABOCA $3\sigma$ contour and the length the aforementioned
extension of the LABOCA $3\sigma$ emission. This yielded outer radii of about
0.70 pc and 0.39 pc for G1.75-0.08 and G11.36+0.80, respectively. In this way,
the filament size refers to the aperture sizes used to derive the flux
densities and hence the mass of the filament.

The derived H2 column and number densities are listed in columns (8) and (9)
in Table 6. The quoted uncertainties in the $N({\rm H_{2}})$ values of the
clumps were propagated from the 350 $\mu$m intensity and dust temperature
uncertainties, while those of the filaments were propagated from the mass
uncertainties. The uncertainties in $n({\rm H_{2}})$ were propagated from the
mass uncertainties and for the clumps also from the uncertainties of the clump
radius. The sizes (projected length and outer radius) of the filaments are
listed in columns (10) and (11) in Table 6.

###  3.5 Virial parameter analysis

The virial parameter of the clumps, which is defined as the ratio of the
virial mass to the clump mass, was calculated using the same approach as in
Miettinen (2020). The clumps were assumed to have a radial density profile of
the form $n(r)\propto r^{-p}$ with $p=1.6$, which is consistent with those
derived for Galactic high-mass star-forming clumps, although we note that
within the usual range of the density power-law indices, $p=1.5-2$, a given
virial mass would vary by only 34% (see Miettinen (2020) and references
therein). The density profiles of our target filaments and clumps could be
derived through combining the available dust continuum data (including the
combination of the ArTéMiS 350 $\mu$m and SPIRE 350 $\mu$m data to probe the
different spatial scales), but an analysis like this is beyond the scope of
this study.

To calculate the total (thermal plus non-thermal) one-dimensional velocity
dispersion, we assumed that the gas kinetic temperature is equal to the dust
temperature derived from the SED fit. The mean molecular weight per free
particle was assumed to be 2.37 to be consistent with the assumptions in the
SED analysis. For the clumps in G11.36+0.80, the spectral line width was taken
to be the FWHM of N2H${}^{+}(1-0)$ detected in the present study (column (3)
in Table 1). For the G1.75-0.08 clumps, we used the FWHM of HCN$(1-0)$
($23.40\pm 1.24$ km s-1 for clump A and $13.50\pm 0.38$ km s-1 for clump B)
from Miettinen (2014). The HCN$(1-0)$ transition was the strongest line
detected towards clump B, while towards clump A, the strongest line was that
of HNCO, which exhibited a very broad profile with an FWHM of $30.40\pm 1.39$
km s-1, however, possibly as a result of shocks (e.g. Yu et al. (2018)).
Hence, HCN$(1-0)$ was selected as a probe of the gas kinematics for clump A as
well.

The derived virial masses and virial parameters are given in Table 7. The
quoted virial mass uncertainties were propagated from the temperature and line
width (FWHM) uncertainties, and those uncertainties along with the clump mass
uncertainties were propagated to the virial parameter uncertainties.

##  4 Discussion

###  4.1 General features of the dust emission and the clumps detected with
the APEX bolometers

The IRDC G1.75-0.08 is composed of two dust clumps separated in projection by
about 3.7 pc (as measured between the LABOCA 870 $\mu$m maxima). The clumps
appear to be connected by a dust-emitting bridge where the central part is
also detected in our ArTéMiS images. If we count the 450 $\mu$m fragment in
between clumps A and B as a third fragment, the average 450 $\mu$m clump
separation would become $1.9\pm 0.2$ pc. The dust clumps in G1.75-0.08 do not
appear to have Spitzer IR counterparts and the clumps also appear dark in the
PACS 70 $\mu$m image, although some diffuse 70 $\mu$m emission can be seen
towards clump A.

The IRDC G11.36+0.80 is filamentary in projected shape, and it has four
prominent LABOCA 870 $\mu$m clumps along its long axis that are all detected
with ArTéMiS. The average projected separation between the LABOCA clumps is
$1.0\pm 0.2$ pc. The interclump dust emission between clumps A and B is
detected at both ArTéMiS wavelengths. That part of the filament exhibits three
dust emission peaks in the ArTéMiS 450 $\mu$m image, and interestingly, those
positions appear to be associated with Spitzer 24 $\mu$m sources. The
northernmost of these 24 $\mu$m sources could be identified with a young
stellar object (YSO) candidate from the Robitaille et al. (2008) catalogue of
intrinsically red mid-IR sources. The middle one of the 24 $\mu$m sources lies
about $1\aas@@fstack{\prime\prime}5$ away from a target position called
G11.38+0.81 P3 in Feng et al. (2000; their Fig. 1). That source is not
detected as a dust emission peak in our ArTéMiS 350 $\mu$m or LABOCA 870
$\mu$m images. The southernmost of the 24 $\mu$m sources is also a prominent
70 $\mu$m source and hence can be considered a candidate YSO embedded in the
filament. The average separation between the 450 $\mu$m maxima along the
G11.36+0.80 filament, including those seen between clumps A and B, is $0.55\pm
0.14$ pc.

The submillimetre clumps A and B in G11.36+0.80 are associated with weak 70
$\mu$m sources, but no counterparts were found from the Spitzer 24 $\mu$m
catalogue (Gutermuth & Heyer (2015)). However, visual inspection of the
Spitzer 24 $\mu$m image suggests that some 24 $\mu$m emission is present in
these clumps. Clumps C and D are both associated with bright 24 $\mu$m sources
and are hence likely to be in a more advanced stage of evolution than clumps A
and B (see Sect. 4.4. for further discussion). We note that the non-detection
of a clump at 24 $\mu$m could be the result of a high column density of the
source as well. However, the H2 column densities towards clumps A–D are
comparable with each other within the derived uncertainties (see column (8) in
Table 6).

The 24 $\mu$m image towards G11.36+0.80 (top middle panel in Fig. 2) shows
that there are dark absoprtion features to the south and south-east of clump
B. The features are partly dark even at 70 $\mu$m, and this region is
coincident with a weak ($3\sigma$) protrusion-like feature in the LABOCA 870
$\mu$m image. These features could be subfilaments attached to G11.36+0.80.
These types of subfilaments have been detected towards some other IRDCs as
well (e.g. Peretto et al. (2013); Miettinen (2018); Leurini et al. (2019)),
and they are denser structures than the striations seen to be connected to
some lower-mass filaments (e.g. Palmeirim et al. (2013); Bonne et al. (2020);
Hacar et al. (2022) and Pineda et al. (2022) for recent reviews). The
subfilaments seen towards G11.36+0.80 could be signatures of the filament
formation through gas flows or gas accretion from the surrounding medium.
However, high-resolution interferometric molecular line observations would be
needed to test these hypotheses and the physical connection of the features
with the main filament in the first place.

Table 7: Virial masses and virial parameters of the clumps in the target
IRDCs.

Source | $M_{\rm vir}$ | $\alpha_{\rm vir}$
---|---|---
| [M☉] |
G1.75-0.08
Clump A | $129\,337\pm 13\,701$ | $23.51\pm 4.71$
Clump B | $36\,596\pm 2\,057$ | $9.71\pm 1.98$
G11.36+0.80
Clump A | $226\pm 6$ | $0.43\pm 0.05$
Clump B | $151\pm 2$ | $0.35\pm 0.05$
Clump C | $280\pm 24$ | $0.44\pm 0.09$
Clump D | $101\pm 3$ | $0.43\pm 0.11$

###  4.2 SED characteristics of clumps and filaments

The physical properties of the clumps were derived using simple single- or
two-temperature MBB models. The angular resolution of the Herschel/PACS data
at 70 $\mu$m and 160 $\mu$m, $5\aas@@fstack{\prime\prime}8\times
12\aas@@fstack{\prime\prime}1$ and $11\aas@@fstack{\prime\prime}4\times
13\aas@@fstack{\prime\prime}4$ (see Sect. 2.3), are comparable to those of the
APEX bolometer data used in the SED fits, namely about
$9\arcsec-19\aas@@fstack{\prime\prime}9$. Moreover, the PACS 160 $\mu$m and
all the APEX bolometer flux densities were integrated within equal apertures,
that is, within the region in which LABOCA emission was detected at $3\sigma$
significance.

Figures 4 and 5 show that a single-temperature MBB and a two-temperature MBB
for 24 $\mu$m bright clumps could be reasonably well fitted to the observed
SEDs between 24 $\mu$m and 870 $\mu$m (the reduced $\chi^{2}$ values range
from 0.90 to 3.55). Although we did not employ the Herschel/SPIRE data in the
fits, they are consistent with the SED fits for the clumps in G1.75-0.08
(especially the 250 $\mu$m flux density). The ArTéMiS 350 $\mu$m flux
densities of the clumps in G1.75-0.08 (as integrated within the LABOCA
$3\sigma$ contour) are $1.5\pm 0.5$ (clump A) and $1.2\pm 0.4$ (clump B) times
the SPIRE 350 $\mu$m catalogue (Molinari et al. 2016a ) flux densities, and
hence agree fairly well within the uncertainties. For the G11.36+0.80 clumps,
the SPIRE catalogue (Molinari et al. 2016a ) flux densities show more scatter
with respect to our SEDs, although in some cases, the agreement between the
two is excellent (e.g. the 250 $\mu$m data point for clump B (factor of 1.13
difference), and the SPIRE 350 $\mu$m data points for clumps C and D (factors
of 1.09 and 0.97 difference, respectively)). The ArTéMiS 350 $\mu$m flux
densities of the clumps in G11.36+0.80 (within the LABOCA $3\sigma$ contour)
are $0.7\pm 0.2$ to $2.0\pm 0.6$ times those of SPIRE at 350 $\mu$m. The best
agreement between the two values is found for clump C, where the ratio of the
two flux density values is $1.0\pm 0.3$.

Figure 5 shows the SEDs of the G11.36+0.80 clumps C and D would not be
consistent with a single-temperature MBB model when the presence of 24 $\mu$m
emission is taken into account. The 24 $\mu$m emission is arising from a
warmer medium near the embedded YSO, while the longer-wavelength far-IR to
submillimetre emission originates in a colder, dusty medium. Hence, a two-
temperature MBB model was adopted for clumps C and D to include the 24 $\mu$m
flux density in the SED fit. This leads to a higher source luminosity than
what would be derived from a single-$T$ model fit at $\lambda\geq 70$ $\mu$m
as was done for clumps A and B.

The SEDs of the target IRDC filaments in the wavelength range of 160–870
$\mu$m were also fitted with a single-temperature MBB model. In terms of the
reduced $\chi^{2}$, the SED fit for G11.36+0.80 was a factor of 1.9 times
better than for G1.75-0.08. For G1.75-0.08, the ArTéMiS flux densities lie
above the model fit by factors of $1.9\pm 0.6$ (350 $\mu$m) and $1.9\pm 0.4$
(450 $\mu$m), while for G11.36+0.80 only the 450 $\mu$m data point lies above
the MBB curve (by a factor of $1.6\pm 0.3$). Within the flux density
uncertainties, these differences are fairly small.

###  4.3 Stability of the filaments and fragmentation into clumps

####  4.3.1 G1.75-0.08 – a 70 $\mu$m dark binary clump system in the Galactic
centre region

The line mass (mass per unit length) of G1.75-0.08 is $M_{\rm line}=1\,011\pm
146$ M☉ pc-1. The dynamical state of the filament can be addressed by
comparing its line mass with the virial or critical line mass (e.g. Fiege &
Pudritz (2000), Eq. (12) therein). To calculate the latter quantity, we used
the total (thermal+non-thermal) velocity dispersion, where the observed
spectral line width (FWHM) was taken to be the FWHM of the HCN$(1-0)$ line
detected towards clump B in G1.75-0.08 from Miettinen (2014), that is,
$13.50\pm 0.38$ km s-1. We note that this is a very broad line width (and it
is even broader (by a factor of 1.73) in clump A; Sect. 3.5), which could be
attributed to multiple factors. From an observational point of view, the
hyperfine structure of HCN was not resolved in the MALT90 spectra, which can
lead to an overestimation of the line width although it was derived through
fitting the hyperfine structure of the transition (Miettinen (2014)).
Moreover, the angular resolution of the Mopra telescope observations employed
by Miettinen (2014) is $38\arcsec$ (HPBW), which corresponds to 1.5 pc at the
cloud distance, and hence the beam might have captured emission from the
turbulent outer parts of the cloud. For example, if G1.75-0.08 follows the
line width–size relation in the central molecular zone, or CMZ, which for HCN
is found to be $\sigma\propto R^{0.62}$ (Shetty et al. (2012); Table 2
therein), the aforementioned HCN$(1-0)$ line width would be expected to be
only $\sim 2.5$ km s-1 on the $\sim 0.1$ pc scale, which is a typical inner
width (FWHM) of filamentary molecular clouds (e.g. Arzoumanian et al. (2019);
Priestley & Whitworth (2022), and references therein). We note that Henshaw et
al. (2016a) derived a velocity dispersion of 11 km s-1 ($\sim 26$ km s-1 FWHM
for a Gaussian profile) for another Galactic centre region IRDC, namely
G0.253+0.016 or the Brick from spectral line observations with Mopra, which is
comparable to the observational results for G1.75-0.08 with the same telescope
(Miettinen (2014)). However, based on much higher angular resolution
($1\aas@@fstack{\prime\prime}7$) observations with the Atacama Large
Millimetre/submillimetre Array (ALMA), Henshaw et al. (2019) derived an
average velocity dispersion of 4.4 km s-1 in the Brick, which demonstrates
that higher resolution spectral line observations towards G1.75-0.08 are also
required. The effect of the angular resolution of the observations on the
derived spectral line widths was also demonstrated by Hacar et al. (2018) in
the case of the Orion integral filament (see e.g. Fig. 6 therein). From a
physical point of view, there are several effects that can lead to line
broadening. First, the HCN lines detected towards the G1.75-0.08 clumps were
not optically thin, and hence the optical thickness effects might contribute
to the broadening of the lines (e.g. Hacar et al. (2016)). On the other hand,
even the optically thin transitions detected towards the clumps had line
widths that are comparable to those of HCN$(1-0)$ (Miettinen (2014), Table 3
therein). Second, G1.75-0.08 might be associated with high-velocity gradients
along the filament (e.g. Federrath et al. (2016); Gong et al. (2018)) that
would be blended in the MALT90 spectra and hence lead to overestimated line
widths. Third, star formation driven outflows and shocks could lead to broad
line profiles. Fourth, G1.75-0.08 is located close to the Galactic centre
($R_{\rm GC}\simeq 270$ pc), where the interstellar medium is highly turbulent
(e.g. Salas et al. (2021) and references therein). At least part of this
turbulent gas could contribute to our single-dish-observed spectral line
widths. Obviously, further spectral line observations are needed to study the
gas kinematics of G1.75-0.08 in more detail.

The gas kinetic temperature was assumed to be equal to the dust temperature
($15.0\pm 0.4$ K) derived from the SED of the filament. For comparison, the
isothermal sound speed at the cloud temperature is only $c_{\rm s}=0.2$ km
s-1, and hence the filament appears to be dominated by supersonic non-thermal
motions with $\sigma_{\rm NT}/c_{\rm s}\simeq 25$. The resulting virial line
mass is $M_{\rm line}^{\rm vir}=15\,306\pm 860$ M☉ pc-1, and the corresponding
$M_{\rm line}/M_{\rm line}^{\rm vir}$ ratio is about $0.07\pm 0.01$. Hence,
G1.75-0.08 as a whole appears to be gravitationally unbound, that is,
subcritical. However, both $M_{\rm line}$ and $M_{\rm line}^{\rm vir}$ are
particularly uncertain in the case of G1.75-0.08 owing to the uncertain dust
opacity in the CMZ (e.g. owing to the change in metallicity) and the very
large uncertainty in the gas velocity dispersion as discussed above. On the
other hand, because G1.75-0.08 lies near the turbulent Galactic centre, its
strongly subcritical (by a factor of $\sim 14$) line mass could be the result
of its environment.

Thermally supercritical filaments are generally found to be self-gravitating
and in rough virial equilibrium (i.e. $M_{\rm line}\sim M_{\rm line}^{\rm
vir}$ within a factor of $\sim 2$; e.g. Arzoumanian et al. (2013); Mattern et
al. 2018b ). The derived $M_{\rm line}/M_{\rm line}^{\rm vir}$ ratio for
G1.75-0.08 is very low compared to the ratio of the general population of
Galactic filaments, but we emphasise again that our estimate of $M_{\rm
line}^{\rm vir}$ is subject to significant uncertainties. Morever, at the
small galactocentric distance of G1.75-0.08, the gas-to-dust mass ratio could
be a factor of $\sim 5$ lower than adopted in our study (Giannetti et al.
(2017); Eq. (2) therein), which would make the line mass of the filament lower
by a similar factor (for a given dust opacity). On the other hand, a higher
metallicity at small Galactocentric distances would make the mean molecular
weight per free particle higher and hence the total velocity dispersion and
virial line mass lower, which would increase the $M_{\rm line}/M_{\rm
line}^{\rm vir}$ ratio. However, even if the metallicity were three times
higher than assumed in the present work (Giannetti et al. (2017); Eq. (3)
therein), the mean molecular weight would increase by less than $\sim 20\%$
(see Kauffmann et al. (2008), Appendix A.1 therein; for this estimate, the He
mass fraction was assumed to increase with metallicity as $Y=0.24+2.4\times Z$
(e.g. Mowlavi et al. (1998))). Hence, if the line mass of G1.75-0.08 is
overestimated, which seems plausible owing to the Galactic gradient in the
gas-to-dust ratio, the cloud could be even more subcritical and an even more
severe outlier compared to the more general filamentary cloud population.

Another factor is that the line mass of G1.75-0.08 is not uniformly
distributed, but the mass is concentrated at both ends of the filamentary
structure. Hence, the system cannot be well approximated by idealised models
of cylindrical cloud fragmentation. The theoretical clump separations
calculated from multiple different cylindrical fragmentation models (Larson
(1985); Nagasawa (1987); Bastien et al. (1991); Jackson et al. (2010);
Fischera & Martin (2012)) were not found to agree well with the observed
projected clump separation of 3.7 pc ($1.9\pm 0.2$ pc if the 450 $\mu$m
fragment in between clumps A and B were counted as a third fragment). If the
gas-to-dust mass ratio is overestimated, as discussed above, the cloud density
needed in the fragmentation analyses would be different. Other caveats that
complicate the comparison between the observed and theoretical clump
separations is that the width of the G1.75-0.08 filament is fairly poorly
defined as the system is composed of two large clumps that are connected by a
narrower dust-emitting bridge, and the filament could be inclined from the
line of sight.

A filament can also collapse along its long axis (longitudinal collapse) in
such a way that the so-called gravitational focussing leads to the
accumulation of gas at the ends of the filament, where gravitational collapse
can then result in the formation of clumps (the so-called edge effect; e.g.
Burkert & Hartmann (2004); Pon et al. (2012); Clarke & Whitworth (2015); Heigl
et al. (2021); Hoemann et al. (2022)). This configuration is qualitatively
consistent with the configuration observed in G1.75-0.08. If the gravitational
focussing scenario is true, the aforementioned possiblity of high-velocity
gradients could well be present owing to the edge effect, where the filament
self-gravity drives the piling up of gas at both ends of the filament. Owing
to the cloud’s location in the Galactic centre region, it might also be
speculated that it was tidally stripped or disrupted, but the volume-averaged
density of over $9\times 10^{3}$ cm-3 of the cloud makes this possibility
questionable (Kruijssen et al. (2014)). However, if the cloud mass is
overestimated, then its true density would be lower and the cloud could be
subject to the aforementioned disruption effects. The red-skewed profiles of
HCN$(1-0)$ and HCO${}^{+}(1-0)$ detected towards clump A (Miettinen (2014))
are indeed an indication of outward motions. For comparison, orbital dynamics
and shearing motions have been suggested to have influenced the structure and
dynamics of the Brick (Henshaw et al. (2019) and references therein), and
these effects could also play a role in G1.75-0.08.

####  4.3.2 Filamentary IRDC G11.36+0.80

For G11.36+0.80, the line mass is $M_{\rm line}=376\pm 59$ M☉ pc-1. When we
use the average FWHM of the detected N2H${}^{+}(1-0)$ lines (Table 1, column
(3)), which is $1.39\pm 0.13$ km s-1, the corresponding critical line mass
becomes $M_{\rm line}^{\rm vir}=186\pm 30$ M☉ pc-1, which yields an $M_{\rm
line}/M_{\rm line}^{\rm vir}$ ratio of $2.0\pm 0.5$. Hence, G11.36+0.80
appears to be marginally supercritical.

Miettinen (2012) concluded that G11.36+0.80 is fragmented into clumps as a
result of the sausage instability. We can revisit this hypothesis using our
improved estimates of the physical properties of the cloud. Our
N2H${}^{+}(1-0)$ data suggest that the filament is moderately dominated by
supersonic non-thermal motions ($\sigma_{\rm NT}/c_{\rm s}=2.6$ on average).
First, we calculated the radial scale height of G11.36+0.80, which depends on
the density at the centre of the filament (e.g. Nagasawa (1987); Eq. (2.3)
therein). The average H2 column density towards clumps A–D is $\langle N({\rm
H_{2}})\rangle=(2.0\pm 0.3)\times 10^{23}$ cm-2 (Table 6, column (8)). We can
estimate the central H2 number density by dividing the column density by the
line-of-sight diameter of the filament, which we assumed to be equal to the
width of the filament. This yields an average central H2 number density of
$(8.3\pm 1.2)\times 10^{4}$ cm-3. The corresponding effective radial scale
height of G11.36+0.80 is $H_{\rm eff}=0.04\pm 0.01$ pc. This can be compared
with the outer radius of the G11.36+0.80 filament as defined with respect to
the LABOCA $3\sigma$ emission level, that is, 0.39 pc. Because the $R_{\rm
fil}/H_{\rm eff}$ ratio is about $9.8\pm 2.6$, G11.36+0.80 appears to be in a
regime of $R\gg H$, where the fluid is compressional, and fragmentation owing
to the sausage fluid instability is expected to produce fragments with a
separation given by the wavelength of the fastest growing unstable mode, which
is $22\times H_{\rm eff}$ (e.g. Jackson et al. (2010); Anathpindika &
Francesco (2021)). The theoretical clump separation, about $0.9\pm 0.2$ pc,
agrees very well with the observed LABOCA clump separation ($1.0\pm 0.2$ pc on
average). As seen in our 450 $\mu$m image, the filament region between clumps
A and B appears to be fragmented into substructures with an average separation
of $0.19\pm 0.08$ pc. This could be a manifestation of secular cloud
fragmentation after the first density enhancements, the clumps seen in our
LABOCA image, were formed via sausage instability (see Sect. 4.5).

A theoretical prediction for the maximum clump mass in G11.36+0.80 is about
$22\times H_{\rm eff}\times M_{\rm line}$, which for the observed line mass is
$331\pm 98$ M☉. The most massive clump in G11.36+0.80 is derived to have a
mass of $633\pm 112$ M☉ (clump C), which exceeds the aforementioned maximum
mass by a factor of $1.9\pm 0.7$. Hence, the two values are comparable within
the uncertainties. We conclude that the mechanism by which G11.36+0.80 has
fragmented into clumps can potentially be assigned to the sausage instability
of the compressible fluid.

###  4.4 Physical properties of the clumps and their potential for high-mass
star formation

The two clumps in G1.75-0.08 are found to be very massive
($\sim(3.8-5.5)\times 10^{3}$ M☉), but their effective radii are also large
($\sim 1.2-1.5$ pc within the LABOCA $3\sigma$ emission). As discussed in
Sect. 4.1, neither of the clumps appears to have a 70 $\mu$m counterpart, and
hence they are candidate high-mass starless clumps. However, because shock-
tracer species such as SiO and HNCO have been detected in the G1.75-0.08
clumps (Miettinen (2014)), they might be associated with star formation
activity, unless the shocks are caused by other processes as discussed above.
Galactic clumps that appear dark at 70 $\mu$m are sometimes found to be
associated with early stages of star formation (e.g. Urquhart et al. (2022)).
Because the aforementioned clumps do not appear to be gravitationally bound
($\alpha_{\rm vir}\gg 2$; Table 7), they cannot be considered high-mass
prestellar clumps even if they were starless. Moreover, as discussed in Sect.
4.3.1, the clump masses might be overestimated because the gas-to-dust ratio
used in the calculation might have been too high, which would increase the
virial parameter values. However, as also discussed in Sect. 4.3.1, the Mopra
(MALT90) spectral line data used in the virial parameter calculation might
overestimate the gas velocity dispersion compared to that within the clump
interiors. For comparison, the MALT90 N2H${}^{+}(1-0)$ lines towards the
LABOCA clumps in G11.36+0.80 are 1.25–1.43 (1.35 on average) times broader
than the present N2H${}^{+}(1-0)$ lines observed with the IRAM 30-metre
telescope. Nevertheless, the G1.75-0.08 clumps do appear to be dominated by
the kinetic energy. That the clumps appear to be starless might be a result of
turbulence, which is suggested to contribute to the suppression of star
formation in the central part of the Galaxy (e.g. Kruijssen et al. (2014)).

The clumps in G11.36+0.80 have masses in the range of $\sim 230-630$ M☉ and
effective radii of $\sim 0.4-0.6$ pc. All the clumps are found to be
gravitationally bound ($\alpha_{\rm vir}<2$; Table 7), and the virial
parameter is remarkably similar among the clumps ($\alpha_{\rm vir}\simeq
0.4$). We note that Feng et al. (2020) derived SED-based dust temperatures for
clumps A and B (their sources P1 and P2) that agree well with our values (the
ratio of the two is $0.85\pm 0.01$ and $0.95\pm 0.02$). Further literature
comparison of the physical properties of the clumps is provided in Appendix B.

The luminosity-to-mass ratios of clumps A–D were derived to be about $0.4\pm
0.1$, $0.4\pm 0.1$, $0.5\pm 0.1$, and $1.0\pm 0.4$ L☉/M☉, respectively (the
uncertainties were propagated from both quantities, and a mean value of the
$\pm$ uncertainties in $L$ was used in the calculation). Because the $L/M$
ratio is expected to increase as the clump evolves, it can be used as an
indicator of the clump evolutionary stage (e.g. Ma et al. (2013); Molinari et
al. 2016b ; Elia et al. (2017)). We note that a similar evolutionary indicator
(bolometric luminosity-to-envelope mass ratio, $L_{\rm bol}/M_{\rm env}$) is
also used for low-mass – class 0 and class I – embedded protostars (e.g. André
et al. (2000) for a review). On the basis of this, clumps A and B appear to be
the least evolved sources in the filament. Clump B in the centre of the
filament is also associated with only a weak 70 $\mu$m source. Clumps C and D
appear to be somewhat more evolved sources owing to their slightly higher
$L/M$ ratio (the differences in $L/M$ are not significant within the
uncertainties). Clumps C and D are both associated with 24 $\mu$m sources,
which also suggests that the clumps are hosting more evolved embedded YSOs.

In Fig. 7 we plot the masses of the studied clumps against their radius. For
comparison, we also plot the mass-radius thresholds for high-mass star
formation derived by Kauffmann & Pillai (2010) and Baldeschi et al. (2017).
The thresholds were scaled to the assumptions we adopted here (see Eqs. (9)
and (10) in Miettinen (2020)). The G1.75-0.08 clumps lie above the thresholds,
for example by factors of $1.9\pm 0.3$ (clump A) and $1.6\pm 0.3$ (clump B)
above the Baldeschi et al. (2017) limit. However, the clumps appear to be
unbound by their self-gravity, and hence it is unclear whether they could
collapse to form massive stars (see also Sect. 4.5).

The G11.36+0.80 clumps lie very close to the Kauffmann & Pillai (2010)
threshold, that is, within factors of $0.7\pm 0.2$ to $1.2\pm 0.1$ ($1.0\pm
0.1$ on average). Assuming a Kroupa (2001) stellar initial mass function and a
star formation efficiency of 0.3, clumps A–D could form a $14\pm 2$ M☉, $12\pm
1$ M☉, $15\pm 1$ M☉, and a $7\pm 0.4$ M☉ star, respectively (see Svoboda et
al. (2016); Sanhueza et al. (2017)). On the basis of this analysis, we cannot
draw firm conclusions about the potential of the clumps in G11.36+0.80 to form
massive stars, but clumps A, B, and C could do so if the clump collapsed to
form only one (massive) star or if the clump or cores within it accreted more
gas from the parent filament.

Figure 7: Clump mass plotted against the effective clump radius. The dashed
green and black lines indicate the empirical mass-radius thresholds for high-
mass star formation proposed by Kauffmann & Pillai (2010) and Baldeschi et al.
(2017), respectively.

###  4.5 Evidence for hierarchical fragmentation

In Fig. 8 we show the ArTéMiS images towards the clumps in our target IRDCs
that show the clearest evidence of fragmention into smaller objects. In Table
8 we list the projected substructure separations in the clumps (column (2))
and the thermal Jeans length of the clumps (column (3); e.g. McKee & Ostriker
(2007) for a review, Eq. (20) therein). The latter parameters agree within
factors of 1.3–1.5 (column (4) in Table 8). The thermal Jeans masses of the
clumps are listed in column (5) in Table 8. The masses of the substructures in
the G1.75-0.08 clumps calculated by assuming that their flux density is equal
to the ArTéMiS 450 $\mu$m peak intensity and that the dust temperature is
equal to that of the parent clump (see e.g. Eq. (2) in Miettinen et al.
(2009)) are $\sim 10^{2}$ times higher than the thermal Jeans mass. For clump
A in G11.36+0.80, the substructure masses are estimated to be $\sim 10$ times
higher than the thermal Jeans mass (column (6) in Table 8). However, the
substructure mass estimates are subject to significant uncertainties owing to
the uncertain dust properties (dust opacity and temperature and the dust-to-
gas mass ratio). We note that the Jeans lengths that take the non-thermal
motions into account would be clearly larger, by factors of $\sim 4-38$
compared to the observed substructure separations. Although further
observations are needed to study the clump fragmentation mechanisms in more
detail, at least clump A in G11.36+0.80 might have fragmented via thermal
Jeans-like instability.

Figure 8 shows that there are stripes in the ArTéMiS images along the right
ascension axis that are artefacts owing to the scanning in only one direction.
In principle, the detected substructure in the clumps could be the result of
such artefacts, or at least the scanning direction could affect the observed
morphology of the substructures (e.g. the southern substructure in clump A of
G1.75-0.08 is elongated in the right ascension direction). Moreover, the
G1.75-0.08 clumps were found to be associated with strongly supersonic non-
thermal motions ($\sigma_{\rm NT}/c_{\rm s}\simeq 26-58$) and be
gravitationally unbound, and hence the presence of substructure and Jeans
fragmentation of the clumps might be questioned. However, Walker et al. (2021)
found evidence that the fragments they detected in the Brick with ALMA at
$0\aas@@fstack{\prime\prime}13$ resolution are the result of thermal Jeans
fragmentation, although the parent cloud is highly turbulent. A similar
conclusion was reached by Lu et al. (2020) on the basis of $\sim
0\aas@@fstack{\prime\prime}2$ resolution ALMA observations of a sample of four
molecular clouds in the Galactic centre region.

Our findings for G11.36+0.80 are consistent with the paradigm where Galactic
filaments fragment in a hierarchical, or at least in a two-level way, that is,
the parent filament is fragmented into clump-size structures via cylindrical
instability, and Jeans fragmentation takes place within the clumps at a higher
density. As discussed in Sect. 1, this behaviour has been found in a wide
variety of Galactic filaments, and the case of G11.36+0.80 with a filament
mass of $\sim 1.5\times 10^{3}$ M☉ and high density of $\sim 10^{4}$ cm-3 adds
another supportive data point to the growing body of evidence for a
filamentary paradigm for star formation.

As discussed in Sect. 4.4, clump A in G11.36+0.80 could in principle be able
to form a high-mass star. However, as the clump appears to be fragmented into
smaller substructures, it could represent the birth site of a stellar cluster
of multiple low- to intermediate-mass stars. Massive star formation is still
possible in this case via competitive accretion, where the gravitational
potential well at the centre of the system allows for an enhanced gas
accretion onto YSOs, and hence to reach high stellar masses (e.g. Bonnell &
Bate (2006)).

Figure 8: Zoom-in images towards the clumps in our sample that show hints of
substructures. ArTéMiS 450 $\mu$m images (colour scale and contours) are shown
for the G1.75-0.08 clumps, while the ArTéMiS 350 $\mu$m image (colour scale
and contours) is shown towards clump A in G11.36+0.80. All the contours are as
in Figs. 1 and 2. The beam sizes (HPBW) are shown above the scale bars in each
panel. Table 8: Clump substructure separations and the Jeans analysis
parameters.

Source | $d_{\rm sep}$ | $\lambda_{\rm J}$ | $d_{\rm sep}/\lambda_{\rm J}$ | $M_{\rm J}$ | $M_{\rm frag}/M_{\rm J}$
---|---|---|---|---|---
| [pc] | [pc] |  | [M☉] |
G1.75-0.08 clump A | $0.44\pm 0.01$ | $0.29\pm 0.02$ | $1.5\pm 0.1$ | $5.6\pm 1.7$ |  $102\pm 38$ and $61\pm 28$
G1.75-0.08 clump B | $0.37\pm 0.01$ | $0.28\pm 0.03$ | $1.3\pm 0.1$ | $5.4\pm 1.9$ |  $117\pm 48$ and $115\pm 50$
G11.36+0.80 clump A | $0.17\pm 0.01$ | $0.24\pm 0.01$ | $0.7\pm 0.03$ | $4.3\pm 0.8$ |  $14\pm 5$ and $7\pm 3$
111111The listed parameters are the observed, projected separation between the
substructures ($d_{\rm sep}$), thermal Jeans length ($\lambda_{\rm J}$), ratio
between the latter two parameters ($d_{\rm sep}/\lambda_{\rm J}$), thermal
Jeans mass ($M_{\rm J}$), and the ratio of the substructure mass to Jeans mass
($M_{\rm frag}/M_{\rm J}$; values are given for the two substructures in the
clump).

##  5 Summary and conclusions

We used ArTéMiS to image the IRDCs G1.75-0.08 and G11.36+0.80 at 350 $\mu$m
and 450 $\mu$m. These new data were used in conjunction with our previous
LABOCA 870 $\mu$m data on the IRDCs and archival Herschel far-IR data.
Spectral line observations of N2H${}^{+}(1-0)$ were also carried out towards
the LABOCA 870 $\mu$m clumps in G11.36+0.80 with the IRAM 30 m telescope. Our
main results are summarised as follows:

  1. 1.

The IRDC G1.75-0.08 is composed of two cold ($T_{\rm dust}\sim 14.5$ K),
massive (several $10^{3}$ M☉) clumps that are projectively separated by $\sim
3.7$ pc. Submillimetre dust emission is also detected in between the clumps,
which in the LABOCA image appears to form a bridge-like structure connecting
the clumps. Both the G1.75-0.08 clumps are 70 $\mu$m dark, and hence they
could be massive starless clumps. Both clumps also fulfil the mass-radius
threshold of high-mass star formation. However, they do not appear to be
bounded by self-gravity and hence cannot be considered prestellar. Moreover,
the ArTéMiS data suggest that the clumps have substructure. No firm
conclusions about the cloud fragmentation mechanism can be drawn on the basis
of the currently available data, but gravitational focussing could lead to a
configuration in which a filament has clumps at both its ends. Owing to the
cloud location near the Galactic centre ($\sim 270$ pc), environmental
conditions such as the elevated level of turbulence and tidal forces could
affect the cloud dynamics. This hypothesis conforms to the finding that the
cloud as whole appears strongly (by an order of magnitude) subcritical, which
would make it an outlier compared to the general population of Galactic
filaments. However, follow-up high-resolution spectral line observations are
needed to confirm or disprove this.

  2. 2.

The filamentary IRDC G11.36+0.80 contains four clear LABOCA 870 $\mu$m clumps
that were all detected in the ArTéMiS images. The clumps have dust
temperatures in the range $\sim 13-15$ K and masses in the range $\sim
230-630$ M☉. The clumps are gravitationally bound, and they appear to be in
somewhat different stages of temporal evolution on the basis of their
luminosity-to-mass ratio. On the basis of the mass-radius comparison, the
clumps do not appear to be candidates for forming high-mass stars. Moreover,
one of the G11.36+0.80 clumps (clump A) shows signatures of further
fragmentation into substructures in our 350 $\mu$m image. The G11.36+0.80
filament was found to be marginally supercritical (by a factor of two), and
the projected separation of the LABOCA clumps is $\sim 1$ pc on average. This
periodic clump separation is consistent with a theoretical prediction of
sausage instability in a compressible fluid. The observed substructure
separation in clump A is roughly consistent (within a factor of $\sim 1.4$)
with the thermal Jeans length of the clump, while the substructure masses
exceed the corresponding thermal Jeans mass by about one order of magnitude.
These findings suggest that the filament is fragmenting in a hierarchical way,
but with different mechanisms on different spatial scales.

Our study is an example of a comparison of filamentary molecular clouds in two
different Galactic environments. Our results for G11.36+0.80 support the
scenario in which filamentary IRDCs are fragmented via sausage instability, as
has been proposed for some other IRDC filaments. A hierarchical or two-level
fragmentation paradigm for Galactic filaments is also consistent with the
detection of substructure in one of the G11.36+0.80 clumps. On the other hand,
the physical characteristics and apparent lack of star formation activity in
G1.75-0.08 might be the result of its location in the Galactic centre region,
where dynamical effects such as tidal forces and shearing motions can affect
the cloud evolution. Studies of larger samples of IRDCs at different
galactocentric distances would be useful to quantify the role played by the
cloud environment in its dynamics, fragmentation mechanism(s), and level of
star formation, and to test the validity of the filament paradigm in the
Galactic centre region.

###### Acknowledgements.

We thank the anonymous referee for providing extensive and insightful comments
and suggestions that helped to improve the quality of this paper. We are
grateful to the staff at the APEX and IRAM 30 m telescopes for performing the
service mode observations presented in this paper. M. M. and Ph. A.
acknowledge support from “Ile de France” regional funding (DIM-ACAV+
Programme) and from the French national programmes of CNRS/INSU on stellar and
ISM physics (PNPS and PCMI). This research has made use of NASA’s Astrophysics
Data System Bibliographic Services, the NASA/IPAC Infrared Science Archive,
which is operated by the Jet Propulsion Laboratory, California Institute of
Technology, under contract with the National Aeronautics and Space
Administration, and Astropy121212www.astropy.org, a community-developed core
Python package for Astronomy (Astropy Collaboration et al. (2013), astropy2018
(2018)). PACS has been developed by a consortium of institutes led by MPE
(Germany) and including UVIE (Austria); KU Leuven, CSL, IMEC (Belgium); CEA,
LAM (France); MPIA (Germany); INAF-IFSI/OAA/OAP/OAT, LENS, SISSA (Italy); IAC
(Spain). This development has been supported by the funding agencies BMVIT
(Austria), ESA-PRODEX (Belgium), CEA/CNES (France), DLR (Germany), ASI/INAF
(Italy), and CICYT/MCYT (Spain). SPIRE has been developed by a consortium of
institutes led by Cardiff University (UK) and including Univ. Lethbridge
(Canada); NAOC (China); CEA, LAM (France); IFSI, Univ. Padua (Italy); IAC
(Spain); Stockholm Observatory (Sweden); Imperial College London, RAL, UCL-
MSSL, UKATC, Univ. Sussex (UK); and Caltech, JPL, NHSC, Univ. Colorado (USA).
This development has been supported by national funding agencies: CSA
(Canada); NAOC (China); CEA, CNES, CNRS (France); ASI (Italy); MCINN (Spain);
SNSB (Sweden); STFC, UKSA (UK); and NASA (USA).

## References

  * Anathpindika & Francesco (2021) Anathpindika, S. V. & Francesco, J. D. 2021, MNRAS, 502, 564
  * André et al. (2000) André, P., Ward-Thompson, D., & Barsony, M. 2000, in Protostars and Planets IV, V. Mannings, A. P. Boss, and S. S. Russell (eds.), University of Arizona Press, Tucson, p. 59
  * André et al. (2014) André, P., Di Francesco, J., Ward-Thompson, D., et al. 2014, in Protostars and Planets VI, H. Beuther, R. S. Klessen, C. P. Dullemond, and Th. Henning (eds.), University of Arizona Press, Tucson, 914 pp., p. 27–51
  * André et al. (2016) André, P., Revéret, V., Könyves, V., et al. 2016, A&A, 592, A54
  * André et al. (2019) André, P., Arzoumanian, D., Könyves, V., et al. 2019, A&A, 629, L4
  * Arzoumanian et al. (2013) Arzoumanian, D., André, P., Peretto, N., et al. 2013, A&A, 553, A119
  * Arzoumanian et al. (2019) Arzoumanian, D., André, P., Könyves, V., et al. 2019, A&A, 621, A42
  * Astropy Collaboration et al. (2013) Astropy Collaboration, Robitaille, T. P., Tollerud, E. J., et al. 2013, A&A, 558, A33
  * (9) Astropy Collaboration, Price-Whelan, A. M., Sipőcz, B. M., et al. 2018, AJ, 156, 123
  * Baldeschi et al. (2017) Baldeschi, A., Elia, D., Molinari, S., et al. 2017, MNRAS, 466, 3682
  * Bastien et al. (1991) Bastien, P., Arcoragi, J.-P., Benz, W., et al. 1991, ApJ, 378, 255
  * Battersby et al. (2010) Battersby, C., Bally, J., Jackson, J. M., et al. 2010, ApJ, 721, 222
  * Beuther & Steinacker (2007) Beuther, H., & Steinacker, J. 2007, ApJ, 656, L85
  * Bonne et al. (2020) Bonne, L., Bontemps, S., Schneider, N., et al. 2020, A&A, 644, A27
  * Bonnell & Bate (2006) Bonnell, I. A. & Bate, M. R. 2006, MNRAS, 370, 488
  * Burkert & Hartmann (2004) Burkert, A. & Hartmann, L. 2004, ApJ, 616, 288
  * Busquet et al. (2013) Busquet, G., Zhang, Q., Palau, A., et al. 2013, ApJ, 764, L26
  * Carey et al. (2009) Carey, S. J., Noriega-Crespo, A., Mizuno, D. R., et al. 2009, PASP, 121, 76
  * Carter et al. (2012) Carter, M., Lazareff, B., Maier, D., et al. 2012, A&A, 538, A89
  * Chambers et al. (2009) Chambers, E. T., Jackson, J. M., Rathborne, J. M., et al. 2009, ApJS, 181, 360
  * Churchwell et al. (2009) Churchwell, E., Babler, B. L., Meade, M. R., et al. 2009, PASP, 121, 213
  * Clarke & Whitworth (2015) Clarke, S. D. & Whitworth, A. P. 2015, MNRAS, 449, 1819
  * Egan et al. (1998) Egan, M. P., Shipman, R. F., Price, S. D., et al. 1998, ApJ, 494, L199
  * Elia et al. (2017) Elia, D., Molinari, S., Schisano, E., et al. 2017, MNRAS, 471, 100
  * Elia et al. (2021) Elia, D., Merello, M., Molinari, S., et al. 2021, MNRAS, 504, 2742
  * Fazio et al. (2004) Fazio, G. G., Hora, J. L., Allen, L. E., et al. 2004, ApJS, 154, 10.
  * Federrath et al. (2016) Federrath, C., Rathborne, J. M., Longmore, S. N., et al. 2016, ApJ, 832, 143
  * Feng et al. (2020) Feng, S., Li, D., Caselli, P., et al. 2020, ApJ, 901, 145
  * Fiege & Pudritz (2000) Fiege, J. D. & Pudritz, R. E. 2000, MNRAS, 311, 85
  * Fischera & Martin (2012) Fischera, J. & Martin, P. G. 2012, A&A, 542, A77
  * Foster et al. (2011) Foster, J. B., Jackson, J. M., Barnes, P. J., et al. 2011, ApJS, 197, 25
  * Foster et al. (2013) Foster, J. B., Rathborne, J. M., Sanhueza, P., et al. 2013, PASA, 30, e038
  * Foster et al. (2014) Foster, J. B., Arce, H. G., Kassis, M., et al. 2014, ApJ, 791, 108
  * Giannetti et al. (2017) Giannetti, A., Leurini, S., König, C., et al. 2017, A&A, 606, L12
  * Gong et al. (2018) Gong, Y., Li, G. X., Mao, R. Q., et al. 2018, A&A, 620, A62
  * Griffin et al. (2010) Griffin, M. J., Abergel, A., Abreu, A., et al. 2010, A&A, 518, L3
  * Gutermuth & Heyer (2015) Gutermuth, R. A. & Heyer, M. 2015, AJ, 149, 64
  * Güsten et al. (2006) Güsten, R., Nyman, L. Å., Schilke, P., et al. 2006, A&A, 454, L13
  * Hacar et al. (2016) Hacar, A., Alves, J., Burkert, A., et al. 2016, A&A, 591, A104
  * Hacar et al. (2018) Hacar, A., Tafalla, M., Forbrich, J., et al. 2018, A&A, 610, A77
  * Hacar et al. (2022) Hacar, A., Clark, S., Heitsch, F., et al. 2022, in Protostars and Planets VII, S. Inutsuka, Y. Aikawa, T. Muto, K. Tomida, and M. Tamura (eds.), University of Arizona Press, Tucson, in press, arXiv:2203.09562
  * Heigl et al. (2021) Heigl, S., Hoemann, E., & Burkert, A. 2021, arXiv:2112.12640
  * (43) Henshaw, J. D., Longmore, S. N., Kruijssen, J. M. D., et al. 2016a, MNRAS, 457, 2675
  * (44) Henshaw, J. D., Caselli, P., Fontani, F., et al. 2016b, MNRAS, 463, 146
  * Henshaw et al. (2019) Henshaw, J. D., Ginsburg, A., Haworth, T. J., et al. 2019, MNRAS, 485, 2457
  * Hoemann et al. (2022) Hoemann, E., Heigl, S., & Burkert, A. 2022, MNRAS, submitted, arXiv:2203.07002
  * Jackson et al. (2010) Jackson, J. M., Finn, S. C., Chambers, E. T., et al. 2010, ApJ, 719, L185
  * Jackson et al. (2013) Jackson, J. M., Rathborne, J. M., Foster, J. B., et al. 2013, PASA, 30, e057
  * Kainulainen et al. (2013) Kainulainen, J., Ragan, S. E., Henning, T., et al. 2013, A&A, 557, A120
  * Kauffmann & Pillai (2010) Kauffmann, J. & Pillai, T. 2010, ApJ, 723, L7
  * Kauffmann et al. (2008) Kauffmann, J., Bertoldi, F., Bourke, T. L., et al. 2008, A&A, 487, 993
  * Kroupa (2001) Kroupa, P. 2001, MNRAS, 322, 231
  * Kruijssen et al. (2014) Kruijssen, J. M. D., Longmore, S. N., Elmegreen, B. G., et al. 2014, MNRAS, 440, 3370
  * Larson (1985) Larson, R. B. 1985, MNRAS, 214, 379
  * Leurini et al. (2019) Leurini, S., Schisano, E., Pillai, T., et al. 2019, A&A, 621, A130
  * Liu et al. (2021) Liu, H.-L., Tej, A., Liu, T., et al. 2021, MNRAS, in press
  * Lu et al. (2020) Lu, X., Cheng, Y., Ginsburg, A., et al. 2020, ApJ, 894, L14
  * Ma et al. (2013) Ma, B., Tan, J. C., & Barnes, P. J. 2013, ApJ, 779, 79
  * Mangum & Shirley (2015) Mangum, J. G. & Shirley, Y. L. 2015, PASP, 127, 266
  * (60) Mattern, M., Kainulainen, J., Zhang, M., et al. 2018a, A&A, 616, A78
  * (61) Mattern, M., Kauffmann, J., Csengeri, T., et al. 2018b, A&A, 619, A166
  * McKee & Ostriker (2007) McKee, C. F. & Ostriker, E. C. 2007, ARA&A, 45, 565
  * Mège et al. (2021) Mège, P., Russeil, D., Zavagno, A., et al. 2021, A&A, 646, A74
  * Miettinen (2012) Miettinen, O. 2012, A&A, 540, A104
  * Miettinen (2014) Miettinen, O. 2014, A&A, 562, A3
  * Miettinen (2018) Miettinen, O. 2018, A&A, 609, A123
  * Miettinen (2020) Miettinen, O. 2020, A&A, 644, A82
  * Miettinen et al. (2009) Miettinen, O., Harju, J., Haikala, L. K., et al. 2009, A&A, 500, 845
  * Molinari et al. (2010) Molinari, S., Swinyard, B., Bally, J., et al. 2010, PASP, 122, 314
  * (70) Molinari, S., Schisano, E., Elia, D., et al. 2016a, A&A, 591, A149
  * (71) Molinari, S., Merello, M., Elia, D., et al. 2016b, ApJ, 826, L8
  * Motte et al. (2018) Motte, F., Bontemps, S., & Louvet, F. 2018, ARA&A, 56, 41
  * Mowlavi et al. (1998) Mowlavi, N., Meynet, G., Maeder, A., et al. 1998, A&A, 335, 573
  * Nagasawa (1987) Nagasawa, M. 1987, Progress of Theoretical Physics, 77, 635
  * Ossenkopf & Henning (1994) Ossenkopf, V. & Henning, T. 1994, A&A, 291, 943
  * Pagani et al. (2009) Pagani, L., Daniel, F., & Dubernet, M.-L. 2009, A&A, 494, 719
  * Palmeirim et al. (2013) Palmeirim, P., André, P., Kirk, J., et al. 2013, A&A, 550, A38
  * Pérault et al. (1996) Pérault, M., Omont, A., Simon, G., et al. 1996, A&A, 315, L165
  * Peretto & Fuller (2009) Peretto, N., & Fuller, G. A. 2009, A&A, 505, 405
  * Peretto et al. (2013) Peretto, N., Fuller, G. A., Duarte-Cabral, A., et al. 2013, A&A, 555, A112
  * Pilbratt et al. (2010) Pilbratt, G. L., Riedinger, J. R., Passvogel, T., et al. 2010, A&A, 518, L1
  * Pineda et al. (2022) Pineda, J. E., Arzoumanian, D., André, P., et al. 2022, in Protostars and Planets VII, S. Inutsuka, Y. Aikawa, T. Muto, K. Tomida, and M. Tamura (eds.), University of Arizona Press, Tucson, in press, arXiv:2205.03935
  * Poglitsch et al. (2010) Poglitsch, A., Waelkens, C., Geis, N., et al. 2010, A&A, 518, L2
  * Pon et al. (2012) Pon, A., Toalá, J. A., Johnstone, D., et al. 2012, ApJ, 756, 145
  * Priestley & Whitworth (2022) Priestley, F. D. & Whitworth, A. P. 2022, MNRAS, 509, 1494
  * Rathborne et al. (2006) Rathborne, J. M., Jackson, J. M., & Simon, R. 2006, ApJ, 641, 389
  * Reid et al. (2014) Reid, M. J., Menten, K. M., Brunthaler, A., et al. 2014, ApJ, 783, 130
  * Revéret et al. (2014) Revéret, V., André, P., Le Pennec, J., et al. 2014, Proc. SPIE, 9153, 915305
  * Rieke et al. (2004) Rieke, G. H., Young, E. T., Engelbracht, C. W., et al. 2004, ApJS, 154, 25
  * Robitaille et al. (2008) Robitaille, T. P., Meade, M. R., Babler, B. L., et al. 2008, AJ, 136, 2413
  * Rosolowsky et al. (2008) Rosolowsky, E. W., Pineda, J. E., Kauffmann, J., et al. 2008, ApJ, 679, 1338
  * Salas et al. (2021) Salas, J. M., Morris, M. R., & Naoz, S. 2021, AJ, 161, 243.
  * Sanhueza et al. (2017) Sanhueza, P., Jackson, J. M., Zhang, Q., et al. 2017, ApJ, 841, 97
  * Shetty et al. (2012) Shetty, R., Beaumont, C. N., Burton, M. G., et al. 2012, MNRAS, 425, 720
  * Shimajiri et al. (2019) Shimajiri, Y., André, P., Ntormousi, E., et al. 2019, A&A, 632, A83
  * Simon et al. (2006) Simon, R., Jackson, J. M., Rathborne, J. M., et al. 2006, ApJ, 639, 227
  * Siringo et al. (2009) Siringo, G., Kreysa, E., Kovács, A., et al. 2009, A&A, 497, 945
  * Svoboda et al. (2016) Svoboda, B. E., Shirley, Y. L., Battersby, C., et al. 2016, ApJ, 822, 59
  * Takahashi et al. (2013) Takahashi, S., Ho, P. T. P., Teixeira, P. S., et al. 2013, ApJ, 763, 57
  * Talvard et al. (2018) Talvard, M., Revéret, V., Le-Pennec, Y., et al. 2018, Proc. SPIE, 10708, 1070838
  * Teixeira et al. (2016) Teixeira, P. S., Takahashi, S., Zapata, L. A., et al. 2016, A&A, 587, A47
  * Urquhart et al. (2022) Urquhart, J. S., Wells, M. R. A., Pillai, T., et al. 2022, MNRAS, 510, 3389
  * Walker et al. (2021) Walker, D. L., Longmore, S. N., Bally, J., et al. 2021, MNRAS, 503, 77
  * Wang et al. (2014) Wang, K., Zhang, Q., Testi, L., et al. 2014, MNRAS, 439, 3275
  * Wenger et al. (2018) Wenger, T. V., Balser, D. S., Anderson, L. D., et al. 2018, ApJ, 856, 52
  * Werner et al. (2004) Werner, M. W., Roellig, T. L., Low, F. J., et al. 2004, ApJS, 154, 1.
  * Wilcock et al. (2012) Wilcock, L. A., Ward-Thompson, D., Kirk, J. M., et al. 2012, MNRAS, 422, 1071
  * Williams et al. (1994) Williams, J. P., de Geus, E. J., & Blitz, L. 1994, ApJ, 428, 693
  * Yu et al. (2018) Yu, N.-P., Xu, J.-L., & Wang, J.-J. 2018, Research in Astronomy and Astrophysics, 18, 015
  * Yuan et al. (2020) Yuan, L., Li, G.-X., Zhu, M., et al. 2020, A&A, 637, A67

##  Appendix A Spitzer and Herschel counterparts of the detected sources

To determine whether the submillimetre dust clumps in the target IRDCs are
associated with YSOs, we cross-matched the ArTéMiS 350 $\mu$m peak positions
of the clumps with the Spitzer/MIPSGAL (Carey et al. (2009)) catalogue
(Gutermuth & Heyer (2015)), which was accessed via the IRSA server. As a
matching radius, we used a value of $10\arcsec$, which corresponds to about
half the LABOCA beam size. Only two matches were found, namely for clumps C
and D in G11.36+0.80 with angular separations between the 350 $\mu$m position
and the MIPS 24 $\mu$m source of $8\aas@@fstack{\prime\prime}3$ and
$5\aas@@fstack{\prime\prime}7$, respectively.

The Herschel source catalogues were also accessed via IRSA, and source
counterparts could be found from the Herschel/Hi-GAL (Molinari et al. (2010))
catalogues (Molinari et al. 2016a ). Again, we used the ArTéMiS 350 $\mu$m
peak positions as the reference positions and used a search radius of
$20\arcsec$, which is comparable to the LABOCA beam size. The 70 $\mu$m, 160
$\mu$m, 250 $\mu$m, 350 $\mu$m, and 500 $\mu$m counterparts were found within
the projected distances of
$4\aas@@fstack{\prime\prime}4-8\aas@@fstack{\prime\prime}7$,
$0\aas@@fstack{\prime\prime}9-8\aas@@fstack{\prime\prime}2$,
$1\aas@@fstack{\prime\prime}9-12\aas@@fstack{\prime\prime}7$,
$1\aas@@fstack{\prime\prime}1-13\aas@@fstack{\prime\prime}2$, and
$5\aas@@fstack{\prime\prime}0-16\aas@@fstack{\prime\prime}7$, respectively.
Clump D in G11.36+0.80 was an exception, where the SPIRE 500 $\mu$m
counterpart position in the Hi-GAL catalogue is offset from the ArTéMiS 350
$\mu$m positions by $27\aas@@fstack{\prime\prime}1$, although the other
Herschel PACS and SPIRE positions lie only
$4\aas@@fstack{\prime\prime}4-7\aas@@fstack{\prime\prime}0$ away from the
ArTéMiS 350 $\mu$m emission maximum. Another special case is clump B in
G1.75-0.08, for which no PACS 160 $\mu$m counterpart could be found from the
Hi-GAL catalogue, but visual inspection of the Herschel 160 $\mu$m image (see
Fig. 1, bottom left panel) clearly shows that the clump is seen in emission at
this wavelength. Hence, we determined the corresponding flux density of the
source from the 160 $\mu$m image using the nearby 160 $\mu$m counterpart of
clump A as a benchmark to determine the same flux conversion factor (including
the zero-level offset and background contribution) for clump B. We note that
because clumps A and B are projectively separated by only about 3.7 pc, the
background (and foreground) contribution to their flux densities can be
expected to be similar.

The photometric data of the Spitzer/MIPS and Herschel/PACS and SPIRE
counterparts of the clumps in G1.75-0.08 and G11.36+0.80 are given in Table 9.

Table 9: Photometric data of the Spitzer/MIPS 24 $\mu$m and Herschel/PACS (70
$\mu$m and 160 $\mu$m) and SPIRE (250 $\mu$m, 350 $\mu$m, and 500 $\mu$m)
counterparts of the clumps in G1.75-0.08 and G11.36+0.80.

Source | $S_{24}$ | $S_{70}$ | $S_{160}$ | $S_{250}$ | $S_{350}$ | $S_{500}$
---|---|---|---|---|---|---
| [mJy] | [Jy] | [Jy] | [Jy] | [Jy] | [Jy]
G1.75-0.08
Clump A | … | … | $17.77\pm 1.09$ | $80.71\pm 2.16$ | $64.36\pm 2.67$ | $38.96\pm 4.40$
Clump B | … | … | …aa$a$No counterpart was found from the Hi-GAL catalogue. If the 160 $\mu$m flux density is determined from the PACS image within an aperture that corresponds to the $3\sigma$ level of the LABOCA emission, the value is $S_{160}=56.1\pm 3.5$ Jy. | $60.41\pm 1.40$ | $51.52\pm 2.15$ | $29.15\pm 2.11$
G11.36+0.80
Clump A | … | $0.38\pm 0.02$ | $2.15\pm 0.18$ | $14.04\pm 0.43$ | $16.69\pm 0.79$ | $25.80\pm 1.38$
Clump B | … | $0.22\pm 0.03$ | $1.26\pm 0.19$ | $26.39\pm 2.15$ | $16.01\pm 0.57$ | $23.04\pm 0.79$
Clump C | $246.61\pm 4.64$ | $6.44\pm 0.09$ | $17.51\pm 0.31$ | $28.29\pm 0.64$ | $26.95\pm 0.71$ | $28.81\pm 1.56$
Clump D | $78.63\pm 1.79$ | $4.51\pm 0.08$ | $16.43\pm 0.26$ | $13.77\pm 0.39$ | $17.85\pm 0.74$ | $10.13\pm 1.74$
131313The flux densities were taken from the MIPSGAL and Hi-GAL catalogues
unless otherwise stated.

##  Appendix B Comparison with the $360\degr$ Hi-GAL catalogue of the physical
properties of the clumps

We cross-matched our clumps with the high-reliability version of the
$360\degr$ Hi-GAL catalogue of the physical properties of the clumps (Elia et
al. (2021)). The cross-matching was performed with respect to the LABOCA peak
positions of our clumps, and the search radius was set to $10\arcsec$, that
is, about half the LABOCA beam size. All of our clumps were found to have a
Hi-GAL counterpart, and they are listed in Table 10 together with the physical
properties as given in the Elia et al. (2021) catalogue.

Elia et al. (2021) derived the dust temperatures and source masses through
fitting the source SEDs with a single-temperature MBB model at wavelengths
$\lambda\geq 160$ $\mu$m. They assumed that the dust opacity is 10 cm2 g-1 at
300 $\mu$m, that the dust emissivity index is $\beta=2$, and that the gas-to-
dust mass ratio is 100 (see Elia et al. (2017)). Hence, the masses from Elia
et al. (2021) should be multiplied by a factor of 1.216 to take the different
assumptions into account when comparing with our results (and additionally to
take the different source distances into account). The bolometric luminosities
in the Elia et al. (2021) catalogue were calculated by integrating the SEDs
from $\lambda=21$ $\mu$m to $\lambda=1.1$ mm, and the luminosities are
proportional to the square of the distance, which also needs to be taken into
account when their values are compared with our results. We provide below a
source-by-source comparison of the physical properties we derived with those
presented in the Hi-GAL catalogue of Elia et al. (2021).

The FWHM size we derived for clump A in G1.75-0.08 from our ArTéMiS image at
350 $\mu$m ($44\aas@@fstack{\prime\prime}7$ when deconvolved from the beam
size) is very similar (only 4% difference) to that in the Hi-GAL catalogue (we
note that 350 $\mu$m is the wavelength band closest to the 250 $\mu$m band
adopted for the catalogue source sizes). No distance value for the clump was
available in the catalogue, but if the 1 kpc based mass is scaled to 8.22 kpc
and multiplied by the aforementioned correction factor, the value ($6\,006\pm
1\,840$ M☉) becomes only $1.1\pm 0.4$ times higher than our value. The dust
temperature we derived for clump A is about $1.1\pm 0.1$ times the catalogue
value and therefore agrees excellently. The bolometric luminosity we derived
is 1.78 times higher than the scaled value from the catalogue, but we note
that our luminosity value is associated with significant uncertainties (see
Table 6, Col. (7)).

For clump B in G1.75-0.08, our deconvolved FWHM size agrees excellently with
the catalogue value (the difference is only 1.7%). A distance value was
available in the Hi-GAL catalogue for this clump, and it is only 1.025 times
lower than ours. The scaled value of the catalogue mass is $1.1\pm 0.2$ times
higher than our estimate, and the two values agree very well within the
uncertainties. The dust temperatures also agree well within the uncertainties
(our value is $1.2\pm 0.1$ times higher). Our luminosity is 2.2 times higher,
but again, we note that our value is associated with large uncertainties. We
note that the source distances in the Elia et al. (2021) catalogue were taken
from Mège et al. (2021), and although the reported distance is comparable to
our value, it was derived from an LSR velocity of 117.2 km s-1, which is
completely different from the LSR velocities of the MALT90 spectral line
detections (around 50 km s-1; Miettinen (2014)). We also note that the Mège et
al. (2021) Hi-GAL catalogue does not contain data on the velocity dispersion
that we could compare with our line width values for the G1.75-0.08 filament.

Only clumps B and D in G11.36+0.80 had distance values in the Hi-GAL
catalogue, and only for clump D this value is similar to ours (the difference
is only a factor of 1.04). The catalogue distance of clump B is 5.5 kpc, but
our N2H${}^{+}(1-0)$ observations show that the clump is physically associated
with the filament at 3.23 kpc. Our deconvolved FWHM sizes at 350 $\mu$m for
the clumps in G11.36+0.80 are 1.17 to 3.51 times larger (2.17 on average) than
the catalogue values. Especially the reported size of clump D in the catalogue
is very compact, only $3\aas@@fstack{\prime\prime}42$, which is about five
times smaller than the SPIRE beam at 250 $\mu$m. Our nominal clump masses are
about 0.4 to 5.1 times the catalogue values (2.4 times higher on average). The
best agreement is found for clump B, where our value is $1.8\pm 0.3$ times the
catalogue value. The poorest agreement is with clump D, for which our mass
estimate is $5.1\pm 1.2$ times higher. However, for clump D, we used a two-
temperature MBB SED fitting approach, and hence a direct comparison with the
Hi-GAL catalogue value is not straightforward and our higher mass estimate is
not surprising. The dust temperatures we derived are $0.7-1.5$ (1.04 on
average) times the catalogue values. The agreement is excellent especially for
clumps B and C, although for clump C, we used a two-temperature MBB fit. The
luminosities we derived are $1.1-3.5$ times higher (1.9 on average) than those
in the catalogue, but owing to the large uncertainties in our derived
luminosities, these differences are not significant. Overall, the physical
properties we derived for our clumps agree well with those reported in the
$360\degr$ Hi-GAL catalogue of Elia et al. (2021).

Table 10: Source counterparts in the $360\degr$ Hi-GAL catalogue of clump
physical properties (Elia et al. (2021)).

Source | Hi-GAL designation | Sep. | $d$ | FWHM | $M$ | $T_{\rm dust}$ | $L_{\rm bol}$
---|---|---|---|---|---|---|---
|  | [$\arcsec$] | [kpc] | [$\arcsec$] | [M☉] | [K] | [L☉]
G1.75-0.08
Clump A | HIGALBM1.7613-0.0857 | 9.94 | … | 42.81 | $73.1\pm 22.4$ | $12.6\pm 1.3$ | 26.1
Clump B | HIGALBM1.7361-0.0762 | 2.87 | 8.02 | 37.67 | $3\,202.2\pm 9.2$ | $12.3\pm 0.5$ | 1 028.5
G11.36+0.80
Clump A | HIGALBM11.3848+0.8110 | 5.27 | … | 18.79 | $96.5\pm 1.1$ | $9.0\pm 0.0$ | 6.1
Clump B | HIGALBM11.3618+0.8016 | 4.93 | 5.5 | 9.26 | $572.4\pm 0.5$ | $13.1\pm 0.2$ | 295.8
Clump C | HIGALBM11.3449+0.7973 | 8.47 | … | 13.62 | $22.0\pm 0.5$ | $13.4\pm 0.1$ | 25.6
Clump D | HIGALBM11.3360+0.7853 | 3.66 | 3.37 | 3.42 | $40.7\pm 0.3$ | $21.4\pm 0.4$ | 257.2
141414The columns are as follows: (1) source name; (2) source designation in
the Elia et al. (2021) Hi-GAL catalogue; (3) projected angular separation
between the LABOCA 870 $\mu$m peak position of our source and the Hi-GAL
catalogue position; (4) source distance in the Hi-GAL catalogue (not available
for clump A in G1.75-0.08 and clumps A and C in G11.36+0.80); (5) circularised
and beam-deconvolved FWHM size measured from the Herschel/SPIRE 250 $\mu$m
image; (6) total clump mass derived from fitting the source SED, and where a
hypothetical distance of 1 kpc was used for sources with unavailable distance;
(7) dust temperature derived from the SED fit; and (8) bolometric luminosity
derived by integrating the SED from $\lambda=21$ $\mu$m to $\lambda=1.1$ mm,
where again a distance of 1 kpc was assumed if it was not available.