# Galactic H ii region IRAS 17149$-$3916 - A multiwavelength study

Ajay Potdar1, Swagat R Das2, Namitha Issac3, Anandmayee Tej1, Sarita Vig1, C.
H. Ishwara Chandra4
1Indian Institute of Space Science and Technology, Thiruvananthapuram 695 547,
Kerala, India
2Indian Institute of Science Education and Research (IISER) Tirupati, Tirupati
517507, India
3Indian Institute of Astrophysics, Koramangala II Block, Bangalore 560 034,
India
4National Centre for Radio Astrophysics, Tata Institute of Fundamental
Research, Pune University Campus, Pune 411007, India
Present affiliation: Space Physics Laboratory, Vikram Sarabhai Space Centre,
Thiruvananthapuram 695 022, Kerala, IndiaCorresponding author:
<EMAIL_ADDRESS>(SRD)Corresponding author<EMAIL_ADDRESS>(AT)

(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

This paper presents a multiwavelength investigation of the Galactic H ii
region IRAS 17149$-$3916\. Using the Giant Meterwave Radio Telescope, India,
first low-frequency radio continuum observations at 610 and 1280 MHz for this
region are presented. The ionized gas emission displays an interesting
cometary morphology which is likely powered by the early type source, E4
(IRS-1). The origin of the cometary morphology is discussed under the
framework of the widely accepted bow shock, champagne flow, and clumpy cloud
mechanisms. The mid- and far-infrared data from Spitzer-GLIMPSE and Herschel-
Hi-GAL reveal a complex network of pillars, clumps, bubble, filaments, and
arcs suggesting the profound influence of massive stars on the surrounding
medium. Triggered star formation at the tip of an observed pillar structure is
reported. High-resolution ALMA continuum data show a string of cores detected
within the identified clumps. The core masses are well explained by thermal
Jeans fragmentation and support the hierarchical fragmentation scenario. Four
‘super-Jeans’ cores are identified which, at the resolution of the present
data set, are suitable candidates to form high-mass stars.

###### keywords:

stars: formation - H II region - ISM: individual objects (IRAS 17149$-$3916)-
radio continuum: ISM

††pubyear: 2020††pagerange: Galactic H ii region IRAS 17149$-$3916 - A
multiwavelength study–Galactic H ii region IRAS 17149$-$3916 - A
multiwavelength study

##  1 Introduction

H ii regions, that are an outcome of the photoionization of newly forming
high-mass stars ($\textup{M}\gtrsim 8~{}M_{\odot}$), not only play a crucial
role in understanding processes involved in high-mass star formation but also
reveal the various feedback mechanisms at play on the surrounding ambient
interstellar medium (ISM) and the natal molecular cloud. Numerous
observational and theoretical studies of H ii regions have been carried out in
the last two decades. However, dedicated multiwavelength studies of star-
forming complexes add to the valuable observational database that provide a
detailed and often crucial insight into the intricacies involved.

In this work, we study the massive star-forming region IRAS 17149$-$3916\.
This region is named RCW 121 in the catalog of $\rm H{\alpha}$ emission in the
Southern Milky Way (Rodgers et al., 1960). The mid-infrared (MIR) dust bubble,
S6, from the catalog of Churchwell et al. (2006) is seen to be associated with
this complex. IRAS 17149$-$3916 has a bolometric luminosity of
$\sim~{}9.4\times 10^{4}~{}L_{\odot}$ (Beltrán et al., 2006). In literature,
several kinematic distance estimates are found for this complex. The near and
far kinematic distance estimates range between 1.6 – 2.2 and 14.5 – 17.7 kpc,
respectively (Walsh et al., 1997; Sewilo et al., 2004; Beltrán et al., 2006;
Watson et al., 2010). In a recent paper, Tapia et al. (2014a) use the spectral
classification of the candidate ionizing star along with near-infrared (NIR)
photometry to place this complex at 2 kpc. This is in agreement with the near
kinematic distance estimates and conforms to the argument of Beltrán et al.
(2006) for assuming the near kinematic distance of 2.1 kpc based on the height
above the Galactic plane. Based on the above discussion, we assume a distance
of 2.0 kpc in this work.

This star-forming region has been observed as part of several radio continuum
surveys at 2.65 GHz (Beard et al., 1969), 4.85 GHz (Wright et al., 1994), and
more recently at 18 and 22.8 GHz by Sánchez-Monge et al. (2013). Using NIR
data, Roman-Lopes & Abraham (2006) detect a cluster of young massive stars
associated with this IRAS source. These authors also suggest IRS-1, the
brightest source in the cluster, to be the likely ionizing star of the H ii
region detected in radio wavelengths. Arnal et al. (2008) probed the
${}^{12}\rm CO$ molecular gas in the region. Based on this observation, they
conclude that RCW 121 and RCW 122 are possibly unrelated star-forming regions
belonging to a giant molecular complex while negating the speculation of these
being triggered by Wolf-Rayet stars located in the HM 1 cluster. In the most
recent work on this source, Tapia et al. (2014a) re-visit the cluster detected
by Roman-Lopes & Abraham (2006). These authors also detect three bright
Herschel clumps, the positions of which coincide with three of the 1.2 mm
clumps of Beltrán et al. (2006).

Introducing the IRAS 17149$-$3916 complex, in Fig. 1, we show the colour
composite image of the associated region. The 5.8 $\rm\mu m$ IRAC band, which
is mostly a dust tracer (Watson et al., 2010), displays an almost closed,
elliptical ring emission morphology. The extent of the bubble, S6, as
estimated by Churchwell et al. (2006) traces this. The cold dust component,
revealed by the Herschel 350 $\rm\mu m$ emission, is distributed along the
bubble periphery with easily discernible cold dust clumps. Ionized gas sampled
in the SuperCosmos $\rm H{\alpha}$ survey (Parker et al., 2005) fills the
south-west part of the bubble and extends beyond towards south. MIR emission
at 21 $\rm\mu m$ is localized towards the south-west rim of the bubble. This
emission is seen to spatially correlate with the central, bright region of
ionized gas (see Fig. 2) and is generally believed to be due to the Ly$\alpha$
heating of dust grains to temperatures of around 100 K (Hoare et al., 1991).

In this paper, we present an in-depth multiwavelength study of this star-
forming region. In discussing the investigation carried out, we present the
radio observations and the related data reduction procedure followed in
Section 2. This section also briefly discusses the salient features of the
archival data used in the study. Section 3 presents the results obtained for
the associated ionized gas and dust environment. Section 4 delves into the
detailed discussion and interpretation related to the observed morphology of
the ionized gas, investigation of the pillar structures, dust clumps in the
realm of triggered star formation, and the nature of the detected dust clumps
and cores. Section 4.5 highlights the main results obtained in this study.

Figure 1: Colour-composite image towards IRAS 17149$-$3916 with the MSX 21
$\rm\mu m$ (red), SuperCosmos $\rm H_{\alpha}$ (blue) and IRAC 5.8 $\rm\mu m$
(green) are shown overlaid with the contours of the Herschel 350 $\rm\mu m$
map. The contour levels are 600, 700, 1000, 1500, 2000, 2500, 4000, and 6000
MJy/sr. The ellipse shows the position and the extent of the bubble, S6, as
identified by Churchwell et al. (2006).

##  2 Observation, data reduction and archival data

###  2.1 Radio Continuum Observation

To probe the ionized component of IRAS 17149$-$3916, we have carried out low-
frequency radio continuum observations of the region at 610 and 1280 MHz with
the Giant Meterwave Radio Telescope (GMRT), Pune, India. GMRT offers a hybrid
configuration of 30 antennas (each of diameter 45 m) arranged in a Y-shaped
layout. The three arms contain 18 evenly placed antennas and provide the
largest baseline of $\sim 25$ km. The central $\rm 1\,km^{2}$ region houses 12
antennas that are randomly oriented with shortest possible baseline of $\sim
100$ m. A comprehensive overview of GMRT systems can be found in Swarup
(1991). The target was observed with the full array for nearly full-synthesis
to maximize the uv coverage which is required to detect the extended, diffuse
emission. Observations were carried out during August 2014 at 610 and 1280 MHz
with a bandwidth of 32 MHz over 256 channels. In the full-array mode, the
resolution is $\sim$5 and 2 arcsec and the largest detectable spatial scale is
$\sim$17 and 7 arcmin at 610 and 1280 MHz, respectively. Radio sources 3C 286
and 3C 48 were selected as the primary flux calibrators. The phase calibrator,
1714-252, was observed after each 40-min scan of the target to calibrate the
phase and amplitude variations over the full observing run. The details of the
GMRT radio observations and constructed radio continuum maps are listed in
Table 1.

Astronomical Image Processing System (AIPS) was used to reduce the radio
continuum data where we follow the procedure detailed in Das et al. (2017) and
Issac et al. (2019). The data sets are carefully examined to identify bad data
(non-working antennas, bad baselines, RFI, etc.), employing the tasks, TVFLG
and UVPLT. Following standard procedure, the gain and bandpass calibration is
carried out after flagging bad data. Subsequent to bandpass calibration,
channel averaging is done while keeping bandwidth smearing negligible.
Continuum maps at both frequencies are generated using the task IMAGR,
adopting wide-field imaging technique to account for w-term effects. Several
iterations of self-calibration (phase-only) are performed to minimize phase
errors and improve the image quality. Subsequently, primary beam correction
was carried out for all the generated maps.

In order to obtain a reliable estimate of the flux density, contribution from
the Galactic diffuse emission needs to be accounted for. This emission follows
a power-law spectrum with a steep negative spectral index of $-2.55$ (Roger et
al., 1999) and hence has a significant contribution at the low GMRT
frequencies (especially at 610 MHz). This results in the increase in system
temperature, which becomes particularly crucial when observing close to the
Galactic plane as is the case with our target, IRAS 17149$-$3916\. The flux
calibrators lie away from the Galactic plane and for such sources at high
Galactic latitudes, the Galactic diffuse emission would be negligible. This
makes it essential to quantify the system temperature correction to be applied
in order to get an accurate estimate of the flux density. Since measurement of
the variation in the system temperature of the antennas at GMRT are not
automatically implemented during observations, we adopt the commonly used
Haslam approximation discussed in Marcote et al. (2015) and implemented in
Issac et al. (2019).

The sky temperature, ${T_{\rm sky}}$, at frequency $\nu$ for the location of
our source is determined using the equation

$T_{\rm sky,\nu}=T_{\rm
sky}^{408}\bigg{(}\frac{\nu}{408~{}\textrm{MHz}}\bigg{)}^{\gamma}$ (1)

where $\gamma=-2.55$ is the spectral index of the Galactic diffuse emission
and $\it{T_{\rm sky}^{\rm 408}}$ is the sky temperature at 408 MHz obtained
from the all-sky 408 MHz survey of Haslam et al. (1982). Using this method, we
estimate the scaling factors of 2.2 and 1.24 at 610 and 1280 MHz,
respectively, which are used to rescale and obtain the final maps.

Table 1: Details of the GMRT radio continuum observations. | 610 MHz | 1280 MHz
---|---|---
Date of Obs. | 8 August 2014 | 14 August 2014
Flux Calibrators | 3C286, 3C48 | 3C286, 3C48
Phase Calibrators | 1714-252 | 1714-252
On-source integration time |  $\sim$4 h |  $\sim$4 h
Synth. beam | 10.69″$\times$6.16″ | 4.41″$\times$2.24″
Position angle. (deg) | 7.04 | 6.37
rms noise (mJy/beam) | 0.41 | 0.07
Int. Flux Density (Jy) | $14.1\pm 1.4$ | $12.6\pm 1.3$
(integrated within $3\sigma$ level) |  |

###  2.2 Other available data

####  2.2.1 Near-infrared data from 2MASS

NIR ($\rm JHK_{s}$) data for point sources around our region of interest has
been obtained from the Two Micron All Sky Survey (2MASS) Point Source Catalog
(PSC). Resolution of 2MASS images is $\sim$5.0 arcsec. We select those sources
for which the “read-flag” values are 1 - 3 to ensure a sample with reliable
photometry. This data is used to study the stellar population associated with
IRAS 17149$-$3916.

####  2.2.2 Mid-infrared data from Spitzer

The MIR images towards IRAS 17149$-$3916 are obtained from the archives of the
Galactic Legacy Infrared Midplane Survey Extraordinaire (GLIMPSE) survey of
the Spitzer Space Telescope. We retrieve images towards the region in the four
Infrared Array Camera (IRAC; Fazio et al. 2004) bands (3.6, 4.5, 5.8, 8.0
$\rm\mu m$). These images have an angular resolution of $\lesssim 2$ arcsec
with a pixel size of $\sim 0.6$ arcsec. We utilize these images in our study
to present the morphology of the MIR emission associated with the region.

####  2.2.3 Far-infrared data from Herschel

The far-infrared (FIR) data used in this paper have been obtained from the
Herschel Space Observatory archives. Level 2.5 processed 70 - 500 $\rm\mu m$
images from Spectral and Photometric Imaging Receiver (SPIRE; Griffin et al.
2010) and JScanam images from the Photodetector Array Camera and Spectrometer
(PACS; Poglitsch et al. 2010), that were observed as part of the Herschel
infrared Galactic plane Survey (Hi-GAL; Molinari et al. 2008), were retrieved.
We use the FIR data to examine cold dust emission and investigate the cold
dust clumps in the regions.

####  2.2.4 Atacama Large Millimeter Array archival data

We make use of the 1.4 mm (Band 6) continuum maps obtained from the archives
of Atacama Large Millimeter Array (ALMA) to identify the compact dust cores
associated with IRAS 17149$-$3916\. These observations were made in 2017 (PI:
A.Sanchez-Monge #2016.1.00191.S) using the extended 12m-Array configuration
towards four pointings, S61, S62, S63, and S64. Each of these pointings sample
different regions of the IRAS 17149$-$3916 complex. The retrieved maps have an
angular resolution of 1.4 arcsec $\times$ 0.9 arcsec and a pixel scale of 0.16
arcsec.

####  2.2.5 Molecular line data from MALT90 survey

The Millimeter Astronomy Legacy Team 90 GHz survey (MALT90; Foster et al.
2011; Jackson et al. 2013) was carried out using the Australia Telescope
National Facility (ATNF) Mopra 22-m telescope with an aim to characterize
molecular clumps associated with H ii regions. The survey dataset contains
molecular line maps of more than 2000 dense cores lying in the plane of the
Galaxy, the corresponding sources of which were taken from the ATLASGAL 870
$\rm\mu m$ continuum survey. The MALT90 survey covers 16 molecular line
transitions lying near 90 GHz with a spectral resolution of $0.11$ km s-1 and
an angular resolution of 38 arcsec. In this study we use the optically thin
$\rm N_{2}H^{+}$ line spectra to carry out the virial analysis of the detected
dust clumps associated with this complex.

##  3 Results

###  3.1 Ionized gas

Figure 2: Top: 610 MHz GMRT map of the region associated with IRAS
17149$-$3916 overlaid with the 3$\sigma$ ($\sigma=0.6\,\rm mJy\,beam^{-1}$)
level contour in light blue. White contours correspond to the 18 GHz emission
mapped with ATCA (Sánchez-Monge et al., 2013), with contour levels starting
from 3$\sigma$ and increasing in steps of 8$\sigma$ ($\sigma=1.2\times
10^{-2}\,\rm Jy\,beam^{-1}$). Bottom: Same as the top panel but for 1280 MHz
GMRT ($\sigma=0.1\,\rm mJy\,beam^{-1}$) and 22.8 GHz ATCA map with the contour
levels starting from 3$\sigma$ and increasing in steps of 5$\sigma$
($\sigma=1.2\times 10^{-2}\,\rm Jy\,beam^{-1}$). The GMRT maps presented here
have been convolved to a resolution of 12 arcsec for 610 MHz and 5 arcsec for
1280 MHz. The beams sizes of the GMRT and ATCA maps are shown in the lower
left and right hand corner, respectively.

We present the first low-frequency radio maps of the region associated with
IRAS 17149$-$3916 obtained using the GMRT. The continuum emission mapped at
610 and 1280 MHz are shown in Fig. 2. The ionized gas reveals an interesting,
large-extent cometary morphology where the head lies in the west direction
with a fan-like diffuse tail opening in the east. The tail has a north-south
extension of $\sim$ 6 arcmin. The bright radio emission near the head displays
a ‘horse shoe’ shaped structure that opens towards the north-east and mostly
traces the south-western portion of the dust ring structure presented in Fig.
1. This is enveloped within the extended and faint, diffuse emission. In
addition, there are several discrete emission peaks seen at both frequencies.
The rms noise and the integrated flux density values estimated are listed in
Table 1. For the latter, the flux density is integrated within the respective
3$\sigma$ contours. The quoted errors are estimated following the method
discussed in Das et al. (2018).

Also included in the figure are contours showing the high-frequency ATCA
observations at 18 and 22.8 GHz, from Sánchez-Monge et al. (2013). These
snapshot ($\sim$ 10 mins) ATCA maps sample only the brightest region towards
the head and the emission at 18 GHz is seen to be more extended. The GMRT and
ATCA maps reveal the presence of several distinct peaks which are likely to be
externally ionized density enhanced regions thus suggesting a clumpy medium.
The fact that some of these peaks could also be internally ionized by newly
formed massive stars cannot be ruled out. Hence, further detailed study is
required to understand the nature of these radio peaks. A careful search of
the SIMBAD/NED database rules out any possible association with
background/foreground radio sources in the line of sight. Comparing with Fig.
1, the ionized emission traced in the $\rm H{\alpha}$ image agrees well with
the GMRT maps. Roman-Lopes & Abraham (2006) and Tapia et al. (2014a) present
the ionized emission mapped in the Br$\gamma$ line which is localized to the
central part and mostly correlates with the bright emission seen in the GMRT
maps.

Assuming the radio emission at 1280 MHz to be optically thin and emanating
from a homogeneous, isothermal medium, we derive several physical parameters
of the detected H ii region using the following expressions from Schmiedeke et
al. (2016),

Lyman continuum photon flux ($N_{\rm Ly}$):

$\begin{split}\left(\frac{N_{\rm Ly}}{\rm s^{-1}}\right)=4.771\times
10^{42}\left(\frac{F_{\nu}}{\rm Jy}\right)\left(\frac{T_{\rm e}}{\rm
K}\right)^{-0.45}\\\ \times\left(\frac{\nu}{\rm
GHz}\right)^{0.1}\left(\frac{D}{\rm pc}\right)^{2}\end{split}$ (2)

Electron number density ($n_{\rm e}$):

$\begin{split}\left(\frac{n_{\rm e}}{\rm cm^{-3}}\right)=2.576\times
10^{6}\left(\frac{F_{\nu}}{\rm Jy}\right)^{0.5}\left(\frac{T_{\rm e}}{\rm
K}\right)^{0.175}\left(\frac{\nu}{\rm GHz}\right)^{0.05}\\\
\times\left(\frac{\theta_{\rm source}}{\rm
arcsec}\right)^{-1.5}\left(\frac{D}{\rm pc}\right)^{-0.5}\end{split}$ (3)

Emission measure (EM):

$\begin{split}\left(\frac{\rm EM}{\rm pc\ cm^{-6}}\right)=3.217\times
10^{7}\left(\frac{F_{\nu}}{\rm Jy}\right)\left(\frac{T_{\rm e}}{\rm
K}\right)^{0.35}\\\ \times\left(\frac{\nu}{\rm
GHz}\right)^{0.1}\left(\frac{\theta_{\rm source}}{\rm
arcsec}\right)^{-2}\end{split}$ (4)

where, $F_{\nu}$ is the integrated flux density of the ionized region, $T_{\rm
e}$ is the electron temperature, $\nu$ is the frequency, $\theta_{\rm source}$
is the angular diameter of the H ii region and D is the distance to this
region. $T_{\rm e}$ is taken to be 5000 K from the radio recombination line
estimates by Caswell & Haynes (1987). Approximating the emission region to an
ellipse, the angular source size ($\theta_{\rm source}$) is taken to be the
geometric mean of the axes of the ellipse and is estimated to be 6.25 arcmin
(3.6 pc). The derived physical parameters are listed in Table 2.

Table 2: Derived physical parameters of the H ii region associated with IRAS 17149$-$3916. Size | log $N_{\rm Ly}$ | EM | $n_{\rm e}$ | Spectral Type
---|---|---|---|---
(pc) |  | (cm-6pc) | (cm-3) |
3.6 | 48.73 | 5.8$\times$105 | 1.3$\times$102 | O6.5V – O7V

If we assume a single star to be ionizing the H ii region and compare the
Lyman-continuum photon flux obtained from the 1280 MHz map with the parameters
of O-type stars presented in Martins et al. (Table 1; 2005), we estimate its
spectral type to be O6.5V – O7V. This can be considered as a lower limit as
the emission at 1280 MHz could be optically thick as well. In addition, one
needs to account for dust absorption of Lyman continuum photons, which can be
significant as shown by many studies (e.g. Paron et al., 2011). The estimated
spectral type suggests a mass range of $\sim 20-40~{}M_{\odot}$ for the
ionizing star (Martins et al., 2005).

To decipher the nature of the ionized emission, we determine the spectral
index, $\alpha$ which is defined as $F_{\nu}\propto\nu^{\alpha}$. The flux
density, $F_{\nu}$, is calculated from the GMRT radio maps. For this, we
generate two new radio maps at 610 and 1280 MHz by setting the $\it uv$ range
to a common value of $0.14-39.7$ k$\lambda$. This ensures similar spatial
scales being probed at both frequencies. Further, the beam size for both the
maps is set to $\rm 12~{}arcsec\times 12~{}arcsec$. $F_{\nu}$ is obtained by
integrating within the area defined by the 3$\sigma$ contour of the new 610
MHz map. The integrated flux density values are estimated to be $\rm 13.7\pm
1.3\,Jy\,$, $\rm 12.1\pm 1.2\,Jy\,$ at 610 and 1280 MHz, respectively. These
yield a spectral index of $-0.17\pm 0.19$. Similar values are obtained for the
central, bright radio emission as well. Within the quoted uncertainties, the
average spectral index is fairly consistent with optically thin, free-free
emission as expected from H ii regions which are usually dominated by thermal
emission. Spectral index estimate of $-0.1$, consistent with optically thin
thermal emission, is also obtained by combining the GMRT flux density values
with the available single dish measurements at 2.65 GHz (Beard et al., 1969)
and 4.85 GHz (Wright et al., 1994).

###  3.2 The dust environment

Figure 3: Dust emission associated with IRAS 17149$-$3916 from the Spitzer-
GLIMPSE and the Herschel-Hi-GAL surveys.

The warm and the cold dust emission associated with IRAS 17149$-$3916 unravels
interesting morphological features like a bubble, pillars, filaments, arcs,
and clumps, that strongly suggest this to be a very active star forming
complex where the profound radiative and mechanical feedback of massive stars
on the surrounding ISM is clearly observed. Fig. 3 compiles the MIR and FIR
emission towards IRAS 17149$-$3916 from the GLIMPSE and Hi-GAL surveys. Apart
from the stellar population probed at the shorter wavelengths, the diffuse
emission seen in the IRAC-GLIMPSE images would be dominated by the emissions
from polycyclic aromatic hydrocarbons (PAHs) excited by the UV photons in the
photodissociation regions (Anderson et al., 2012; Paladini et al., 2012).
Close to the hot stars, there would also be significant contribution from
thermally emitting warm dust that is heated by stellar radiation (Watson et
al., 2008). In H ii regions, emission from dust heated by trapped Ly$\alpha$
photons (Hoare et al., 1991), would also be present in these IRAC bands. In
the wavelength regime of the 21 $\rm\mu m$ MSX, the emission is either
associated with stochastically heated Very Small Grains (VSGs) or thermal
emission from hot big grains (BGs). As we go to the FIR Hi-GAL maps, cold dust
emission dominates and shows up as distinct clumps and filamentary structures
where, emission in the 70 $\rm\mu m$ band is dominated by the VSGs and the
longer wavelength bands like 250 $\rm\mu m$ band trace emissions from the BGs
(Paladini et al., 2012).

#### Dust temperature and column density maps

To understand the nature of cold dust emission, we generate the dust
temperature and the molecular hydrogen column density maps following the
procedure detailed in Das et al. (2018) and Issac et al. (2019) and briefly
stated here. A pixel-by-pixel modified blackbody modelling to the observed
spectral energy distribution is carried out. As discussed in these two papers,
the 70 $\rm\mu m$ data is not used because there would be appreciable
contribution from the warm dust component hence rendering a single modified
blackbody fit inaccurate. Thus, we have the FIR emission at 160, 250, 350, and
500 $\rm\mu m$ mostly on the Rayleigh-Jeans part to constrain the model given
by

$F_{\nu}-I_{\rm bg}=B_{\nu}(\nu,T_{\rm d})~{}\Omega~{}(1-{\rm
e}^{-\tau_{\nu}})$ (5)

where, $F_{\nu}$ is the observed flux density, $I_{\rm bg}$ is the background
flux density, $B_{\nu}(\nu,T_{\rm d})$ is the Planck function at the dust
temperature $\rm T_{\rm d}$, $\Omega$ is the solid angle subtended by a pixel
(all maps are convolved to a common resolution of 35.7 arcsec and regridded to
a common pixel size of $\rm 14~{}arcsec\times 14~{}arcsec$). The background
flux is estimated from a nearby region relatively free of clumpy and bright
emission. The optical depth $\tau_{\nu}$ in Eqn. 5 can be expressed as

$\tau_{\nu}=\mu_{\rm H_{2}}~{}N({\rm H_{2}})~{}m_{\rm H}~{}\kappa_{\nu}$ (6)

where, $\mu_{\rm H_{2}}$ is the mean molecular weight which is taken as 2.8
(Kauffmann et al., 2008), $N({\rm H_{2}})$ is the column density, $m_{\rm H}$
is the mass of hydrogen atom and $\kappa_{\nu}$ ($\rm cm^{2}\,g^{-1}$) is the
dust opacity which is given as (Hildebrand, 1983)

$\kappa_{\nu}=0.1\left(\frac{\nu}{1200~{}{\rm GHz}}\right)^{\beta}$ (7)

Here, $\beta$ denotes the dust emissivity spectral index and a typical value
of 2 estimated in several star regions is assumed. In fitting the modified
blackbody to the observed flux densities, $N({\rm H_{2}})$ and $T_{\rm d}$ are
kept as free parameters.

Figure 4: Column density (a), dust temperature (b), and chi-square
($\chi^{2}$) (c) maps of the region associated with the H ii region. The 610
MHz ionized emission overlaid as magenta and gray contours on column density
and dust temperature maps, respectively. The contour levels are 0.01, 0.03,
0.06, 0.12, and 0.2 Jy/beam. The retrieved clump apertures are shown on the
column density (in blue) and dust temperature (in black) maps following the
nomenclature as discussed in the text. The black contour in (a) shows the area
integrated for estimating $n_{0}$ (refer Section 4.3.)

The column density and dust temperature maps generated are shown in Fig. 4.
The goodness of the fits for each pixel can be seen in the $\chi^{2}$ map
where the maximum $\chi^{2}$ value is seen to be $\sim 8$. The column density
map presents a triangular morphology with three distinct, bright and dense
regions. A network of broad filaments are also seen in the map. The dust
temperature map is relatively patchy with regions of higher temperature within
the radio nebula. A region with warm temperature is seen to be located towards
the south-east of IRAS 17149$-$3916, the signature of which can be seen in the
Herschel maps shown in Fig. 3. The western side of the H ii region shows
comparatively cold temperatures. Furthermore, the filamentary features seen in
the column density map are mostly revealed as distinct low temperature lanes.

#### Dust clumps and cores

The FIR and the column density maps show the presence of dust clumps. These
clumps are identified using the Herschel 350 $\rm\mu m$ map and the
Dendrogram111https://dendrograms.readthedocs.io/en/stable/ algorithm. Using
this algorithm, we identify the smallest structures, called the ‘leaves’, in
the 350 $\rm\mu m$ map, which in this case are the cold dust clumps. The key
input parameters for the identification of the clumps are (1) min_value =
$3\sigma$ and (2) min_delta = $\sigma$, where
$\rm\sigma(=191.2\,MJy\,sr^{-1})$ is the rms level of the 350 $\rm\mu m$ map.
An additional parameter, min_pix = N, is also used, which is the minimum
number of pixels required for a ‘leaf’ to be considered an independent entity.
To ensure that the clumps are resolved, the value of N is chosen to be 7
pixels, the beam area of the 350 $\rm\mu m$ map. Setting these parameters, we
extract three cold dust clumps. The central panel of Fig. 5 shows the 350
$\rm\mu m$ map overlaid with the retrieved apertures of the three detected
clumps labelled 1, 2, and 3. The physical parameters of the detected clumps
are listed in Table 3. These are derived from the 350 $\rm\mu m$, column
density and dust temperature maps. The peak positions are determined from the
350 $\rm\mu m$ map. The clump radii, $r=(A/\pi)^{0.5}$, where $A$ is the
enclosed area within the retrieved clump apertures.

Figure 5:  The central panel is the Herschel 350 $\rm\mu m$ map overlaid with
the Dendrogram retrieved clump apertures in black. White circles represent
four pointings of the ALMA observations, S61, S62, S63, and S64. The 1.4 mm
ALMA map towards each pointing is placed at the corners of the 350 $\rm\mu m$
image. The apertures of the dense cores extracted using the Dendrogram
algorithm are overlaid and labelled on these maps. The beam sizes of the 1.4
mm maps are given towards the lower left hand corner of each image. Table 3:
Physical parameters of the detected dust clumps associated with IRAS
17149$-$3916

Clump | Peak position | Mean $T_{\rm d}$ | $\Sigma N({\rm H_{2}})$ | Radius | Mass | Mean $N({\rm H_{2}})$ | No. density $(n_{\textup{H}_{2}})$ | $M_{\rm vir}$ | $\alpha_{\rm vir}$
---|---|---|---|---|---|---|---|---|---
| RA (J2000) | DEC (J2000) | (K) |  $(\times 10^{23}$ cm${}^{-2})$ | (pc) | ($M_{\odot}$) |  $(\times 10^{22}$ cm-2) |  $(\times 10^{4}$ cm-3) | ($M_{\odot}$) |
1 | 17 18 24.32 | -39 19 28.04 | 27.8 | 6.1 | 0.3 | 250 | 5.5 | 5.0 | 452 | 1.75
2 | 17 18 18.77 | -39 19 04.54 | 25.0 | 7.1 | 0.3 | 292 | 5.4 | 5.1 | 600 | 2.15
3 | 17 18 22.57 | -39 18 36.44 | 27.3 | 2.8 | 0.2 | 117 | 3.6 | 4.6 | 450 | 4.08

To estimate the masses of the detected clumps, we utilize the column density
map and use the following expression

$M_{\rm clump}=\mu_{\rm H_{2}}~{}\Sigma N({\rm H_{2}})~{}A_{\rm
pixel}~{}m_{\rm H}$ (8)

where, $\mu_{\rm H_{2}}$ is the mean molecular weight taken as 2.8, $\Sigma
N({\rm H_{2}})$ is the integrated column density over the clump area, $A_{\rm
pixel}$ is the pixel area in $\rm cm^{2}$ and $m_{\rm H}$ is the mass of
hydrogen atom. The number density is determined using the expression $n_{\rm
H_{2}}=3\,N({\rm H_{2}})/4r$. The peak positions of the Clumps 1, 2, and 3
agree fairly well with Clumps III, I, and II, respectively, detected by Tapia
et al. (2014a) using Herschel maps. Comparing the masses presented in Table 3,
the estimates given in Tapia et al. (2014b) are higher by a factor of 8, 1.5,
and 3 for the Clumps 1, 2, and 3, respectively.

ALMA continuum data enables investigation of the detected clumps at high
resolution. In Fig. 5, we show ALMA dust continuum maps at 1.4 mm towards the
four pointings marked and labelled as S61, S62, S63, and S64 where the first
three pointings lie mostly within three clumps and S64 lies outside towards
the east. Using the same Dendrogram algorithm, several cores are identified.
The key input parameters to the Dendrogram algorithm are, min_value =
$3\sigma$, min_delta = $\sigma$, and min_pix = N, where $\sigma$ is the rms
level and $N(=60)$ is the beam area of the 1.4 mm maps. In order to avoid
detection of spurious cores, we only retain those with peak flux density
greater than $5\sigma$. Applying these constraints, seven cores are identified
towards S61 and S62 each, one towards S63 and two towards S64.

To further study these dense cores, we estimate their physical parameters.
Adopting the formalism described by Cheng et al. (2018) and assuming the
emission at 1.4 mm to be optically thin, the masses are estimated using the
following equation

$\displaystyle M$ $\displaystyle=$ $\displaystyle\displaystyle
0.0417\,M_{\odot}\left({\textrm{e}}^{0.533(\lambda/{1.3\,\textrm{mm}})^{-1}(T/{20\,\textrm{K}})^{-1}}-1\right)\left(\frac{F_{\nu}}{\textrm{mJy}}\right)$
(9)
$\displaystyle\displaystyle\times\left(\frac{\kappa_{\nu}}{0.00638\,\textrm{cm}^{2}\,\textrm{g}^{-1}}\right)^{-1}\left(\frac{d}{\textrm{kpc}}\right)^{2}\left(\frac{\lambda}{1.3\,\textrm{mm}}\right)^{3}$

Here, $F_{\nu}$ is the integrated flux density of each core, $d$ is the
distance to the source and $\lambda$ is the wavelength. Opacity,
$\kappa_{\nu}$ is estimated using equation 7 with the dust emissivity spectral
index $\beta$ fixed at 2.0. For cores detected in S61, S62, and S63, mean dust
temperatures of the respective clumps are taken. For S64, the mean dust
temperature for the region covering the S64 pointing is used. The effective
radius, $r=(A/\pi)^{0.5}$, of each core is also estimated where $A$ is area
enclosed within each core aperture. The identified cores with the retrieved
apertures are shown in Fig. 5 and the estimated physical parameters are list
in Table 4. The uncertainties related to the missing flux effect are not taken
into account in deriving these parameters since it may not be significant
given that the largest recoverable scales is quoted to be $\sim$10 arcsec for
this ALMA dataset which is appreciably larger than the typical sizes of the
detected cores. Barring the largest detected core which has an angular size of
$\sim$7 arcsec, the average size of the cores is 3 arcsec.

Table 4: Parameters of the detected dust cores extracted associated with IRAS 17149$-$3916 | Core | Peak position | Flux density | Radius | Mass
---|---|---|---|---|---
|  | RA (J2000) | DEC (J2000) | (mJy) | (pc) | ($M_{\odot}$)
S61 | 1 | 17 18 23.03 | -39 19 12.95 | 9.7 | 0.01 | 1.7
| 2 | 17 18 23.75 | -39 19 18.47 | 62.2 | 0.02 | 10.7
| 3 | 17 18 23.45 | -39 19 19.53 | 32.8 | 0.02 | 5.7
| 4 | 17 18 24.17 | -39 19 22.35 | 11.2 | 0.01 | 1.9
| 5 | 17 18 24.34 | -39 19 25.09 | 106.7 | 0.02 | 18.4
| 6 | 17 18 24.76 | -39 19 19.50 | 4.5 | 0.01 | 0.8
| 7 | 17 18 24.93 | -39 19 28.12 | 31.8 | 0.02 | 5.5
S62 | 1 | 17 18 20.38 | -39 18 54.32 | 46.0 | 0.02 | 9.4
| 2 | 17 18 20.09 | -39 18 54.43 | 9.8 | 0.01 | 2.0
| 3 | 17 18 19.64 | -39 18 56.07 | 7.9 | 0.02 | 1.6
| 4 | 17 18 19.63 | -39 19 04.70 | 2.1 | 0.01 | 0.4
| 5 | 17 18 18.82 | -39 19 04.94 | 3.1 | 0.01 | 0.6
| 6 | 17 18 18.35 | -39 19 01.65 | 19.2 | 0.02 | 3.9
| 7 | 17 18 18.42 | -39 19 05.37 | 8.2 | 0.01 | 1.7
S63 | 1 | 17 18 23.48 | -39 18 40.70 | 357.7 | 0.03 | 62.3
S64 | 1 | 17 18 29.98 | -39 19 11.22 | 7.0 | 0.02 | 1.1
| 2 | 17 18 29.10 | -39 18 54.11 | 1.9 | 0.01 | 0.3

###  3.3 Molecular line observation of identified clumps

We use the optically thin $\rm N_{2}H^{+}$ line emission to determine the
$V_{\textup{LSR}}$ and the line width, $=\Delta V$ of the clumps. The line
spectra are extracted by integrating over the retrieved apertures of the
clumps as shown in Fig. 6. The $\rm N_{2}H^{+}$ spectra have seven hyperfine
structures and the hfs method of CLASS90 is used to fit the observed spectra.
The line parameters retrieved from spectra are listed in Table 5. The
$V_{\textup{LSR}}$ determined agrees well with the value of $\rm
13.7\,kms^{-1}$ obtained from $\rm CS(2-1)$ observations of the region by
Bronfman et al. (1996)

Figure 6: Spectra of the optically thin N2H+ line emission extracted over the three identified clumps associated with IRAS 17149$-$3916\. The red curves are the hfs fit to the spectra. The estimated $V_{\textup{LSR}}$ is denoted by the dashed blue line and the location of the hyperfine components by magenta lines. Table 5: The retrieved $\rm N_{2}H^{+}$ line parameters, $V_{\textup{LSR}}$, $\Delta V$, $T_{\rm mb}$ and $\int T_{\textup{mb}}{\rm dV}$ for the identified clumps associated with IRAS 17149$-$3916. Clump | $V_{\textup{LSR}}$ | $\Delta V$ | $T_{\rm mb}$ | $\int T_{\textup{mb}}{\rm dV}$
---|---|---|---|---
| (km s-1) | (km s-1) | (K) | (K km s-1)
1 | -13.0 | 3.2 | 1.5 | 5.4
2 | -13.8 | 3.7 | 1.6 | 6.8
3 | -13.8 | 3.9 | 0.8 | 3.1

##  4 Discussion

###  4.1 Understanding the morphology of the ionized gas

As seen, the GMRT maps reveal the large extent and prominent cometary
morphology of the H ii region associated with IRAS 17149$-$3916 which was
earlier discussed as a roundish H ii region by Tapia et al. (2014a). In this
section, we attempt to investigate the likely mechanism for this observed
morphology. The widely accepted models to explain the formation of cometary H
ii regions are (1) the bow-shock model (e.g. van Buren & Mac Low 1992), (2)
the champagne-flow model (e.g. Tenorio-Tagle 1979), and (3) the mass loading
model (e.g. Williams & McKee 1997). However, subsequent studies, like those
conducted by Cyganowski et al. (2003) and Arthur & Hoare (2006), find the
‘hybrid’ models, that are a combination of these, to better represent the
observed morphologies.

The bow-shock model assumes a wind-blowing, massive star moving supersonically
through a dense cloud. Whereas, the champagne-flow model invokes a steep
density gradient encountered by the expanding H ii region around a newly
formed stationary, massive star possibly located at the edge of a clump. Here,
the ionized gas flows out towards regions of minimum density. In comparison,
the model proposed by Williams & McKee (1997) invokes the idea of strong
stellar winds mass loading from the clumpy molecular cloud and the cometary
structure unfolds when a gradient in the geometrical distribution of mass
loading centres are introduced. In this model the massive, young star is
considered to be stationary as in the case of the champagne-flow model.

While observation of ionized gas kinematics is required to understand the
origin of the observed morphology, in the discussion that follows, we discuss
a few aspects based on the radio, column density, and FIR maps of the region
associated with IRAS 17149$-$3916 along with the identification of E4 as the
likely ionizing star (refer Section 4.2). Following the simple analytic
expressions discussed in Das et al. (2018), we derive a few shock parameters
to probe the bow-shock model. Taking the spectral type of E4 to be O6.5V – O7V
as estimated from the radio flux density, and assuming it to move at a typical
speed of $\rm 10~{}km/s$ through the molecular cloud, we calculate the ‘stand-
off’ distance to range between $\rm
2.6~{}arcsec(0.02~{}pc)-3.1~{}arcsec(0.03~{}pc)$. This is defined as the
distance from the star at which the shock occurs and where the momentum flux
of the stellar wind equals the ram pressure of the surrounding cloud. The
theoretically estimated value is significantly less than the observed distance
of $\sim\rm 84~{}arcsec(0.8~{}pc)$ between E4 and the cometary head. Taking
viewing angle into consideration would decrease the theoretical estimate thus
widening the disparity further. Based on the above estimations, it is unlikely
that the bow-shock model would explain the cometary morphology. To confirm
further, we determine the trapping parameter which is the inverse of the
ionization fraction. As the ionizing star moves supersonically through the
cloud, the swept off dense shells trap the H ii region within it and its
expansion is eventually inhibited by the ram pressure. Trapping becomes more
significant when recombinations far exceed the ionizing photons. Studies of a
large number of cometary H ii regions show the trapping parameter to be much
greater than unity (Mac Low et al., 1991). For IRAS 17149$-$3916, we estimate
the value to lie in the range $3.2-3.5$ which indicates either weak or no bow
shock. Similar interpretations are presented in Das et al. (2018) and Veena et
al. (2016). The trapping parameters obtained by these authors lie in the range
1.2 – 4.3.

To investigate the other models, namely the champagne-flow and clumpy/mass
loading wind models, we compare the observed spatial distribution of the dust
component and the ionized gas. The FIR and column density maps presented in
Section 3.2 show a complex morphology of pillars, arcs, filaments in the
region with detected massive clumps towards the cometary head. The steep
density gradient towards the cometary head is evident. Without the ionized and
molecular gas kinematics information, it is difficult to invoke the champagne-
flow model. However, the maps do show the presence of clumps towards the
cometary head which could act as potential mass loading centres and thus
support the clumpy cloud model. Further observations and modelling are
essential before one can completely understand the mechanisms at work.

###  4.2 Ionizing massive star(s)

Roman-Lopes & Abraham (2006) have studied the associated stellar population
towards IRAS 17149$-$3916 in the NIR. Using the colour-magnitude diagram,
these authors show the presence of a cluster of massive stars within the
infrared nebula and suggest IRS-1 to be the likely ionizing source. In a later
study, Tapia et al. (2014a) have supported this view citing the spectroscopic
classification of IRS-1 as O5 – O6 by Bik et al. (2005) and consistency with
the Lyman continuum photon flux estimated from the radio observations by
Sánchez-Monge et al. (2013).

Figure 7: (a) J vs J-H colour magnitude diagram of the sources (cyan dots)
associated with IRAS 17149$-$3916 and located within the 3$\sigma$ radio
contour. The nearly vertical solid lines represent the ZAMS loci with 0, 15,
and 30 magnitudes of visual extinction corrected for the distance. The
slanting lines show the reddening vectors for spectral types B0 and O5. (b)
J-H vs H-K colour-colour diagram for sources (magenta dots) in the same region
as (a). The cyan and black curves show the loci of main sequence and giants,
respectively, and are taken from Koornneef (1983) and Bessell & Brett (1988).
The locus of classical T Tauri adopted from Meyer et al. (1997) is shown as
long dashed line. The locus of Herbig AeBe stars shown as short dashed line is
adopted from Lada & Adams (1992). The parallel lines are the reddening vectors
where cross marks indicate intervals of 5 mag of visual extinction. The
colour-colour plot is divided into three regions, namely, ‘F’, ‘T’, and ‘P’
(see text for more discussion). The interstellar reddening law assumed is from
Rieke & Lebofsky (1985). The magnitudes, colours and various loci plotted in
both the diagrams are in the Bessell & Brett (1988) system. The identified
early type (earlier than B0) and YSOs candidates are highlighted as blue and
red stars, respectively. Of these, the ones located towards the central,
bright radio emission are labelled as ‘E’ and ‘Y’, respectively. (c) An
enlarged view of the bottom left portion of (b) showing the position of
spectral types on the main sequence locus. The location of the source E4 and
the errors on the colour are also shown. Figure 8: Panel a: Spitzer 5.8
$\rm\mu m$ image in grey scale overlaid by 610 MHz radio contours. The contour
levels are 0.002, 0.01, 0.03, 0.06, 0.12, and 0.2 Jy/beam. The YSOs and
massive stars detected towards the central ionized region are shown in red and
blue coloured stars, respectively. Panel b: This shows the zoom-in view of
central part of the H ii region.

While the spectral type estimated from the GMRT radio emission at 1280 MHz is
consistent with that of IRS-1, we investigate the stellar population within
the radio emission for a better understanding. In Figs. 7(a) and (b), we plot
the NIR colour-magnitude and colour-colour diagrams, respectively, of 2MASS
sources located within the 3$\sigma$ radio contour. Fig. 7(c) shows an
enlarged view of the bottom left portion of Fig. 7(b) to highlight the
location of the star E4 with respect to the main sequence locus. Following the
discussion given in Tej et al. (Fig. 7; 2006) the colour-colour plot is
classified into ‘F’, ‘T’, and ‘P’ regions. The ‘F’ region is occupied by
mostly field stars or Class III sources, the ‘T’ region is for T-Tauri stars
(Class II YSOs) and protostars (Class I YSOs) populate the ‘P’ region. As seen
from the figures, there are sixteen sources earlier than spectral type B0 and
eighteen identified YSOs. The sample of identified YSOs fall in the ‘T’
region, the sources of which are believed to be Class II objects with NIR
excess. Sources that lie towards the central bright, radio emitting region are
labelled in the figures with prefixes of ‘E’ for the sources earlier than B0
and ‘Y’ for the YSOs. As indicated in Fig. 7(b), early type sources E1 and E2
are also the identified YSOs, Y1 and Y3, respectively. The coordinates and NIR
magnitudes of these selected sources are listed in Table 6. Fig. 8 shows the
spatial distribution of the above sources with respect to the radio and 5.8
$\rm\mu m$ emission.

As seen from the above analysis, in addition to the presence of possible
discrete radio sources that could be internally ionized, several massive stars
are also identified from the NIR colour-magnitude and colour-colour diagrams.
Hence, it is likely that ionization in this H ii region is the result of this
cluster of massive stars. However, the observed symmetrical, cometary
morphology of the ionized emission strongly suggests that the ionization is
mostly dominated by a single star. As seen in Fig. 7(a), out of the early type
sources that lie towards the central, bright radio emission, the colour and
magnitude of the source E4 is consistent with a spectral type of $\sim$O6. A
careful scrutiny of the Fig. 7(b) show that early type stars E1 and E2 are
embedded Class II sources and hence unlikely to be the main driving source.
Sources E3, E5, E6 are possibly reddened giants or field stars. The location
of early type star, E4 (which is the source IRS-1) in the colour-colour
diagram (see enlarged view shown in Fig. 7c) agrees fairly well with the
spectral type estimate of $\sim$O6 obtained from the colour-magnitude diagram
and strongly advocates it as the dominating exciting source. This is
consistent with the identification of IRS-1 as the ionizing star in previous
studies. Spatially also, the location of E4 clearly suggests its role in the
formation of the network of pillar like structures observed (see Section 4.3).
As mentioned earlier, the spectral type of E4, estimated from NIR
spectroscopy, is in good agreement with the radio flux. Location wise,
however, it is 30 arcsec away from the radio peak. This offset could be
attributed to density inhomogeneity or clumpy structure of the surrounding,
ambient ISM. Supporting this scenario of E4 being the dominant player, are the
interesting pillar like structures revealed in the MIR images discussed in the
next section.

Table 6: Early type and YSOs detected within the central, bright, radio
emission of IRAS 17149$-$3916

Source | Coordinates | J | H | K
---|---|---|---|---
| RA (J2000) | DEC (J2000) |  |  |
|  | Early-type sources |  |  |
E1 (Y1) | 17 18 22.21 | -39 18 42.24 | 11.364 | 10.039 | 8.957
E2 (Y3) | 17 18 22.85 | -39 18 22.45 | 15.356 | 12.086 | 10.086
E3 | 17 18 25.11 | -39 18 46.41 | 15.398 | 12.486 | 11.361
E4 | 17 18 25.45 | -39 19 08.61 | 8.654 | 8.208 | 7.927
E5 | 17 18 25.68 | -39 18 26.65 | 11.790 | 9.617 | 8.606
E6 | 17 18 25.94 | -39 18 00.89 | 7.715 | 7.191 | 7.021
|  | YSOs |  |  |
Y1 (E1) | 17 18 22.21 | -39 18 42.24 | 11.364 | 10.039 | 8.957
Y2 | 17 18 22.28 | -39 18 12.69 | 15.064 | 14.161 | 13.578
Y3 (E2) | 17 18 22.85 | -39 18 22.45 | 15.356 | 12.086 | 10.086
Y4 | 17 18 22.87 | -39 18 58.45 | 14.413 | 13.346 | 12.393
Y5 | 17 18 23.28 | -39 19 07.77 | 12.993 | 11.906 | 11.135
Y6 | 17 18 24.44 | -39 19 11.04 | 14.324 | 12.958 | 12.166
Y7 | 17 18 25.14 | -39 19 25.95 | 14.741 | 13.657 | 12.664
Y8 | 17 18 25.28 | -39 18 24.76 | 12.818 | 11.372 | 10.420
Y9 | 17 18 28.30 | -39 19 49.71 | 14.677 | 13.449 | 12.680

###  4.3 Triggered star formation

#### Pillar Structures

In Fig. 9, we illustrate the identification of pillar structures in the IRAC 8
$\rm\mu m$ image. The MIR emission presents a region witnessing a complex
interplay of the neutral, ambient ISM with the ionizing radiation of newly
formed massive star(s). The boxes labelled ‘A’ and ‘B’ show prominent pillar
structures, the orientation of these are clearly pointed towards E4. This
strongly suggests E4 as the main sculptor of the detected pillars.
Furthermore, it also supports the identification of E4 as the main ionizing
source of the H ii region .

Figure 9: (a) 8.0 $\rm\mu m$ map of IRAS 17149$-$3916 from the Spitzer-GLIMPSE
survey. Blue arrows highlight the pillar structures identified within Boxes
‘A’ and ‘B’. The white ‘+’ mark shows the position of E4 (IRS-1). (b) A zoom-
in on ‘A’ at IRAC 5.8 $\rm\mu m$. (c) 1280 MHz map covering the pillar ‘A’.

One of the mechanisms widely accepted to explain the formation of these
pillars is the radiative driven implosion (RDI) (Lefloch & Lazareff, 1994).
Here, pre-existing clouds exposed to newly forming massive star(s) are
sculpted into pillars by slow photoevaporation caused due to strong
impingement of ionizing radiation. The other being the classical collect and
collapse model of triggered star formation proposed by Elmegreen & Lada
(1977). Under this framework, the expanding H ii region sweeps up the
surrounding material, creating dense structures that could eventually form
pillars in their shadows.

Figs. 9(b) and (c) show the zoomed in IR and radio view of pillar ‘A’. Clearly
seen is a slightly elongated and bright radio source at the pillar head. To
ascertain the nature of the bright radio source, we estimate few physical
parameters using the 1280 MHz GMRT map. Using the 2D fitting tool of CASA
viewer we fit a 2D Gaussian and determine the deconvolved size and flux
density of this source to be $\rm 19.05~{}arcsec\times 8.07~{}arcsec$
($\theta_{\rm source}=12.4\rm arcsec;0.12\rm pc$) and $\rm 445\,mJy\,$,
respectively. Inserting these values in equations 2, 3, and 4, we get
$logN_{\rm Ly}=47.30$ , $n_{\rm e}=3.9\times 10^{3}~{}\rm cm^{-3}$ and EM=
$\rm 1.8\times 10^{6}~{}pc~{}cm^{-6}$. The ionized mass ($\rm
M_{ion}=\frac{4}{3}\pi r^{3}\mathit{n_{\rm e}}m_{p}$, where $\rm r$ is the
radius of the source and $\rm m_{p}$ is the mass of proton) is also calculated
to be $\rm 0.09~{}M_{\odot}$. The estimated values of these physical
parameters lie in between the typical values of compact and UCHII region
(Kurtz, 2002; Martín-Hernández et al., 2005; Yang et al., 2021). Hence this
radio source could well represent an intermediate evolutionary stage between a
compact and an UCH ii region thus indicating a direct signature of triggered
star formation at the tip of the pillar. An alternate picture for the bright
radio emission at the head of the pillar could also be external ionization by
the ionizing front emanating from E4. Such externally ionized, tadpole-shaped
structures have been studied in the Cygnus region, where the ionized front
heads point towards the central, massive Cygnus OB2 cluster (Isequilla et al.,
2019). In support of the former scenario of the UCH ii region, a bright and
compact 5.8 $\rm\mu m$ emission region is seen that is co-spatial with radio
emission. This compact IR emission is seen in all IRAC bands and Herschel
images. While it is not listed as an IRAC point source in the GLIMPSE catalog,
it is included in the PACS 70 $\rm\mu m$ point source catalog (Marton et al.,
2017). There also exists a 2MASS counterpart (within $\sim$3 arcsec) but has
been excluded from the YSO identification procedure owing to poor quality
photometry in one or more 2MASS bands. It is thus likely that the compact IR
emission sampled in the Spitzer-GLIMPSE and Herschel images is the massive YSO
powering the UCH ii region.

Several studies (e.g. Billot et al. 2010; Paron et al. 2017) have shown
evidence of star formation in the pillar tips in the form jets, outflows, YSO
population, etc. The driving mechanism for this triggered star formation, RDI,
is initiated when the propagating ionizing front traverses the pillar head
creating a shell (known as the ionized boundary layer, IBL) of ionized gas. If
the pressure of the IBL exceeds the internal pressure of the neutral gas
within the pillar head then shocks are driven into it. This leads to
compression and subsequent collapse of the clump leading to star formation.
However, to comment further on the detected UCH ii region on the tip of pillar
‘A’ and link its formation to RDI, one needs to conduct pressure balance
analysis using molecular line data as discussed in Paron et al. (2017) and
Ortega et al. (2013). These authors have used $\rm{}^{13}CO$ transitions for
their analysis which is not possible in our case as CO molecular line data
with adequate spatial resolution is not available for this region of interest.
Furthermore, attempting any study with the detected MALT90 transitions is
difficult given the limited spatial resolution. In a recent study, Menon et
al. (2020) have carried out involved hydrodynamical simulations to study
pillar formation in turbulent clouds. As discussed by these authors, star
formation triggered in pillar heads can be explained without invoking the RDI
mechanism. Gravitational collapse of pre-existing clumps can lead to star
formation without the need for ionizing radiation to play any significant
role. From their simulations, they conclude that compressive turbulence driven
in H ii regions, which competes with the reverse process of photoevaporation
of the neutral gas, ultimately dictates the triggering of star formation in
these pillars. Further high-resolution studies are required to understand the
nature of the compact radio and IR emission at the head of the pillar ‘A’.

#### Dust clumps and the collect and collapse mechanism

The detection of clumps and the signature of fragmentation to cores is evident
from the FIR and sub-mm dust continuum maps presented in Fig. 5. Investigating
the collect and collapse hypothesis is necessary to explain whether the dust
clumps are a result of swept-up material that is accumulated or these clumps
are pre-existing entities. Towards this, we carry out a simple analysis and
evaluate few parameters such as the dynamical age ($t_{\rm dyn}$) of the H ii
region and the fragmentation time ($t_{\rm frag}$) of the cloud.

Assuming that the H ii region associated with IRAS 17149$-$3916 expands in a
homogeneous cloud, the dynamical timescale can be estimated using the
classical expressions from Spitzer (1978) and Dyson & Williams (1980),

$t_{\rm dyn}=\frac{4}{7}\frac{R_{\rm St}}{a_{\rm Hii}}\left[\left(\frac{R_{\rm
if}}{R_{\rm St}}\right)^{7/4}-1\right]$ (10)

where, $a_{\rm Hii}$ is the isothermal sound speed and is assumed to be 10 km
s-1, $R_{\rm if}=1.8$ pc is the radius of the H ii region determined from the
geometric mean of an ellipse visually fit to encompass the ionized emission in
the 610 MHz GMRT map. $R_{\rm St}$ is the Strömgren radius, given by the
following equation

$R_{\rm St}=\left(\frac{3N_{\rm Ly}}{4\pi n^{2}_{0}\alpha_{\rm
B}}\right)^{1/3}$ (11)

where, $N_{\rm Ly}$ is the Lyman continuum flux, $n_{0}$ is the initial
particle density of the ambient gas. To derive $n_{0}$, we assume that the
dense, bright region seen in the column density map is swept-up material due
to the expansion of the H ii region and it was initially homogeneously
distributed within the radius of the H ii region. To estimate the mass of this
swept-up material, we use equation 8 and the column density map (refer Section
3.2). Integrating within the black contour shown in Fig. 4(a), the mass is
estimated to be $1645~{}{M_{\odot}}$. Taking the observed estimate of the
radius, we calculate $n_{0}$ to be $9.5\times 10^{2}$ cm-3. $\alpha_{B}$ is
the coefficient of radiative recombination and is determined using the
expression (Dyson & Williams, 1980):

$\alpha_{\rm B}=2.6\times 10^{-13}\left(\frac{10^{4}\textup{K}}{T_{\rm
e}}\right)^{0.7}\textup{cm}^{3}\>\textup{s}^{-1}$ (12)

where, $T_{\rm e}=5000$ K, is the electron temperature. Using the above
parameters and from the spectral type of the ionizing source (O6.5V - O7V), we
estimate the $t_{\rm dyn}$ to be $\sim$ 0.2 Myr.

Using the formalism discussed in Whitworth et al. (1994), we proceed next to
estimate the fragmentation time scale of a cloud and that can be written as

$t_{\rm frag}=1.56\>a^{7/11}_{{\rm s},0.2}\>\mathit{N^{-1/11}_{{\rm
Ly},\textup{49}}}\>n^{-5/11}_{0,3}\>\textup{Myr}$ (13)

where, $a_{\rm s}=a_{\rm s,0.2}\times 0.2$ km s-1 is the speed of sound in the
shocked layer and is taken as 0.3 km s-1 (Das et al., 2017). $N_{\rm
Ly}=N_{\rm Ly,\textup{49}}\times 10^{49}$ s-1 is the ionizing photon flux and
$n_{0}=n_{0,3}\times 10^{3}$ cm-3 is the initial particle density of the
ambient gas. Plugging in the values in the above expression, we estimate
$t_{\rm frag}$ to be $\sim$ 2.2 Myr. Comparing the estimates of the two time
scales involved, it is seen that the fragmentation time scale is more than a
factor of 10 larger than the dynamical time scale of the H ii region. This
essentially indicates that if the clumps detected are the result of swept-up
material due to expansion of the H ii region, then the shell has not got
enough time to fragment thus making the collect and collapse process highly
unlikely here. Such a scenario has been invoked by Ji et al. (2012) for the
dust bubble N22. In contrast, Das et al. (2017), in their investigation of
bubble CS51 found support for the collect and collapse hypothesis. Thus,
further studies, as indicated earlier, are required to probe the RDI process
not only with regards to the pillar structures but also the detected clumps.

###  4.4 Nature of the detected dust clumps and cores

#### Virial analysis of the dust clumps

Here, we investigate the gravitational stability of the identified dust clumps
associated with IRAS 17149$-$3916\. This would enable us to determine whether
these clumps are gravitationally bound or not. The virial mass, $M_{\rm vir}$,
of a dust clump is the amount of mass that can be supported against self-
gravity purely by thermal and non-thermal gas motion. This is given by
Contreras et al. (2016)

$M_{\rm vir}=\frac{5\ r\ \Delta V^{2}}{8\ {\rm ln}(2)\ a_{1}\ a_{2}\ G}\sim
209\ \frac{1}{a_{1}\ a_{2}}\left(\frac{\Delta V}{\rm km\ s^{-1}}\right)^{2}\
\left(\frac{r}{\rm pc}\right)M_{\odot}$ (14)

In the above equation, $\Delta V$ is the line width of the optically thin $\rm
N_{2}H^{+}$ line, $r$ is the radius of clumps taken from Table 3, the constant
$a_{1}$ accounts for the correction for power-law density distribution, and is
given by $a_{1}=(1-p/3)/(1-2p/5)$, for $p<2.5$ (Bertoldi & McKee, 1992) where
we adopt $p=1.8$ (Contreras et al., 2016). The constant $a_{2}$ accounts for
the shape of the clump which we assume to be spherical and take $a_{2}$ as 1.
We also calculate the virial parameter, $\alpha_{\rm vir}=M_{\rm vir}/M_{\rm
clump}$. The estimated values of $M_{\rm vir}$ and $\alpha_{\rm vir}$ are
listed in Table 3. As discussed in Kauffmann et al. (2013) and Tang et al.
(2019), $\alpha_{\rm vir}=2$ sets a lower threshold for gas motions to prevent
collapse in the absence of magnetic field and/or external pressure. The virial
parameter estimate for Clump 1 is $<2$ indicating that it is gravitationally
bound and hence likely to collapse. However, Clump 2 is marginally above this
threshold and Clump 3, which shows signature of star formation in the form of
an UCH ii region has a higher virial parameter value of 4.1. Similar values of
$\alpha>2$ has been observed for protostellar and prestellar dense cores by
Tang et al. (2019). These authors have used the $\rm C^{18}O$ line and discuss
the contribution from turbulence as a primary factor that would significantly
affect the line width and hence overestimate the virial mass. While turbulence
gets dissipated in the densest region of molecular clouds and the $\rm
N_{2}H^{+}$ line used here is a dense gas tracer, it is likely that the
resolution of the MALT90 survey does not probe the inner dense cores and the
observed velocity dispersion is influenced by the outer and more turbulent
region. High-resolution molecular line observations are thus essential to
probe the nature of the clumps.

#### Clump fragmentation and the detected cores

ALMA 1.4 mm continuum map (Fig. 5) is seen to resolve the identified dust
clumps into a string of cores, with masses ranging between
$0.3-62.3~{}M_{\odot}$ and radii $\sim 0.01$ pc, thus indicating a scenario of
hierarchical fragmentation. If we assume that fragmentation of the clumps is
governed by thermal Jeans instability, then the initially homogeneous gas
clump has a Jeans length and mass given by Sanhueza et al. (2019)

$\lambda_{J}=\sigma_{\rm th}\left(\frac{\pi}{G\rho}\right)^{1/2}$ (15)

and

$M_{J}=\frac{4\pi\rho}{3}\left(\frac{\lambda_{J}}{2}\right)^{3}=\frac{\pi^{5/2}}{6}\frac{\sigma_{\rm
th}^{3}}{\sqrt{G^{3}\rho}}$ (16)

where $\rho$ is the mass density, $G$ the gravitational constant and
$\sigma_{\rm th}$ the thermal velocity dispersion (the isothermal sound speed)
and is given by

$\sigma_{\rm th}=\left(\frac{k_{B}T}{\mu m_{\rm H}}\right)^{1/2}$ (17)

where $k_{B}$ is the Boltzmann constant and $\mu$ the mean molecular weight.
As the thermal velocity dispersion will be dominated by $\rm H_{2}$ and He, we
consider $\mu=2.37$ (Kauffmann et al., 2008; Wang et al., 2014; Sanhueza et
al., 2019). Using the clump parameters tabulated in Table 3, we estimate
$\lambda_{J}$ and $M_{J}$ of the clumps which are listed in Table 7. If
turbulence drives the fragmentation instead, then the turbulent Jeans length
and mass for each clump is derived by replacing the thermal velocity
dispersion with the clump velocity dispersion estimated from the observed line
width of the dense gas tracer $\rm N_{2}H^{+}$ which is a good approximation
for the turbulent line width. The calculated $\lambda_{\rm turb}$ and $M_{\rm
turb}$ values are given in Table 7.

Table 7: Hierarchical fragmentation of clumps associated with IRAS$-$3916

Clump | $\sigma_{\rm th}$ | $\sigma^{a}$ | $\lambda_{J}$ | $M_{J}$ | $\lambda_{\rm turb}$ | $M_{\rm turb}$
---|---|---|---|---|---|---
| ($\rm km\,s^{-1}$) | ( $\rm km\,s^{-1}$) | (pc) | ($M_{\odot}$) | (pc) | ($M_{\odot}$)
1 | 0.3 | 1.4 | 0.2 | 6.8 | 0.8 | 543
2 | 0.3 | 1.6 | 0.2 | 5.3 | 0.9 | 807
3 | 0.3 | 1.7 | 0.1 | 5.6 | 0.8 | 819

a $\sigma=\Delta V/\sqrt{8\,{\rm ln}2};\Delta V$ being the $\rm N_{2}H^{+}$
line width.

The turbulent Jeans masses are $\sim 80-150$ times larger than the thermal
Jeans mass. Comparing with the derived core masses, it is seen that 11 out of
the 15 detected cores ($\sim 73\%$) in the three clumps have masses less than
the Jeans mass. This suggests that the observed cores are consistent with the
prediction of Jeans fragmentation without invoking turbulence indicating that
it does not play a significant role in the fragmentation process. Similar
results are obtained by Sanhueza et al. (2019) who studied the 70 $\rm\mu m$
dark massive clumps in early stages using ALMA data. As discussed by these
authors, the majority of detected cores having masses less than the thermal
Jeans mass supports competitive accretion and hierarchical fragmentation
frameworks. The four cores whose mass exceeds the Jeans mass (the ‘super-
Jeans’ cores) are suitable candidates for forming high-mass cores. However,
further high-resolution observations are essential to completely understand
the fragmentation process, if any, at the core level.

Figure 10: Masses of the dense cores identified from the 1.4 mm ALMA maps and
the cold dust clumps identified using the 350 $\rm\mu m$ Herschel map of IRAS
17149$-$3916 are plotted as a function of their effective radii, depicted by
circles and ‘$\star$’s, respectively. The shaded area corresponds to the low-
mass star-forming region that do not satisfy the condition,
$M>870\,M_{\odot}(r/\rm pc)^{1.33}$ (Kauffmann & Pillai, 2010). Black-dashed
lines indicate the surface density thresholds of 0.05 and 1 $\rm g\,cm^{-2}$
defined by Urquhart et al. (2014) and Krumholz & McKee (2008), respectively.
The red lines represent the surface density thresholds of 116 $M_{\odot}\,\rm
pc^{-2}$ ($\sim 0.024\,\rm g\,cm^{-2}$) and 129 $M_{\odot}\,\rm pc^{-2}$
($\sim 0.027\,\rm g\,cm^{-2}$) for active star formation proposed by Lada et
al. (2010) and Heiderman et al. (2010), respectively.

In Fig. 10, we plot the estimated mass and radius of the identified clumps and
cores. The plot also compiles several surface density thresholds proposed by
various studies to identify clumps/cores with efficient and active star
formation (Lada et al., 2010; Heiderman et al., 2010; Urquhart et al., 2014).
In addition, criteria for these to qualify as high-mass star-forming ones are
also included. All the detected clumps and cores associated with IRAS
17149$-$3916 are seen to be active star-forming regions. The three clumps
satisfy the empirical mass-radius criteria, $M>870\,M_{\odot}(r/\rm
pc)^{1.33}$ defined by Kauffmann & Pillai (2010), and hence are likely to
harbour massive star formation. At the core scale ($\rm<0.1~{}pc$), Krumholz &
McKee (2008) have posed a theoretical surface density threshold of $\rm
1\,g\,cm^{-2}$, below which cores would be devoid of high-mass star formation.
From the figure, we see that there are four cores (2 in Clump 1, 1 each in
Clumps 2 and 3) which have masses $\gtrsim 10~{}M_{\odot}$ and above this
surface density limit. These are the ‘super-Jeans’ cores discussed above.
High-resolution molecular line observations are essential to shed better light
on the nature of the cores, the gas kinematics involved and for accurate
determination of physical parameters like temperature, mass, etc.

###  4.5 Conclusion

Using multiwavelength data, we have carried out a detailed analysis of the
region associated with IRAS 17149$-$3916\. The important results of this study
are summarized below.

  1. 1.

Using the GMRT, we present the first low-frequency radio continuum maps of the
region mapped at 610 and 1280 MHz. The H ii region, previously believed to be
nearly spherical, displays a large-extent cometary morphology. The origin of
this morphology is not explained by the bow shock model. The presence of dense
clumps towards the cometary head indicates either the champagne flow or the
clumpy cloud model but further observations of the ionized gas kinematics are
essential to understand the observed morphology.

  2. 2.

The integrated flux densities yields an average spectral index value of
$-0.17\pm 0.19$ consistent with thermal free-free emission. If powered by a
single massive star, the estimated Lyman continuum photon flux suggests an
exciting star of spectral type O6.5V – O7V star.

  3. 3.

NIR colour-magnitude and colour-colour diagrams show the presence of a cluster
of massive stars (earlier than spectral type B0) located within the bright,
central radio emitting region. MIR and FIR images show complex and interesting
features like a bubble, pillars, clumps, filaments,and arcs revealing the
profound radiative and mechanical feedback of massive stars on the surrounding
ISM.

  4. 4.

The spatial location of source, E4 (IRS-1), and the orientation of observed
pillar structures with respect to it, strongly suggest it as the dominant
driving source for the cometary H ii region. This view finds support from the
position of E4 in the colour-magnitude and colour-colour diagrams. Further,
its spectral type estimation from literature agrees well with that estimated
for the exciting source of the H ii region from GMRT data.

  5. 5.

The column density map reveals the presence of dust clumps towards the
cometary head while the dust temperature map appears to be relatively patchy
with regions of higher temperature within the radio nebula. The dust clumps
identified using the Herschel 350 $\rm\mu m$ map have masses ranging between
$\sim$100 - 300 $\rm M_{\odot}$ and radius $\sim$0.2 - 0.3 pc. Virial analysis
using the $\rm N_{2}H^{+}$ shows that the south-east clump (#1) is
gravitationally bound. For the other two clumps (# 2 and 3), the line widths
would possibly have contribution from turbulence thus rendering larger values
of the virial parameter.

  6. 6.

A likely compact/UCH ii region is seen at the tip of a pillar structure
oriented towards the source E4 thus suggesting evidence of triggered star
formation under the RDI framework. In addition, the detected dust clumps are
investigated to probe the collect and collapse model of triggered star
formation. The estimated dynamical time scales are seen to be smaller by a
factor of $\sim$10 compared to the fragmentation timescale of the clumps thus
clearly negating the collect and collapse mechanism at work.

  7. 7.

The ALMA 1.4 mm dust continuum map probes the dust clumps at higher resolution
and reveal the presence of 17 compact dust cores with masses and radii in the
range of $0.3-62.3~{}M_{\odot}$ and 0.01 – 0.03 pc, respectively. The largest
and the most massive core is located within Clump 3. The estimated core masses
are consistent with thermal Jeans fragmentation and support the competitive
accretion and hierarchical fragmentation scenario.

  8. 8.

Four ‘super-Jeans’ fragments are detected and are suitable candidates for
forming high-mass stars and their mass and radius estimates satisfy the
various threshold defined in literature for the potential high-mass star-
forming cores.

## Acknowledgements

We would like to thank the referee for comments and suggestions which helped
in improving the quality of the manuscript. We thank the staff of the GMRT
that made the radio observations possible. GMRT is run by the National Centre
for Radio Astrophysics of the Tata Institute of Fundamental Research. The
authors would like to thank Dr. Alvaro Sánchez-Monge for providing the FITS
image of the radio maps presented in Sánchez-Monge et al. (2013). CHIC
acknowledges the support of the Department of Atomic Energy, Government of
India, under the project 12-R&D-TFR-5.02-0700. This work is based [in part] on
observations made with the Spitzer Space Telescope, which is operated by the
Jet Propulsion Laboratory, California Institute of Technology under a contract
with NASA. This publication also made use of data products from Herschel (ESA
space observatory). This publication makes use of data products from the Two
Micron All Sky Survey, which is a joint project of the University of
Massachusetts and the Infrared Processing and Analysis Center/California
Institute of Technology, funded by the NASA and the NSF. This work makes use
of the ATLASGAL data, which is a collaboration between the Max-Planck-
Gesellschaft, the European Southern Observatory (ESO) and the Universidad de
Chile. This paper makes use of the following ALMA data:
ADS/JAO.ALMA#2016.1.00191.S. ALMA is a partnership of ESO (representing its
member states), NSF (USA) and NINS (Japan), together with NRC (Canada), MOST
and ASIAA (Taiwan), and KASI (Republic of Korea), in cooperation with the
Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and
NAOJ. This research has made use of the SIMBAD database, operated at CDS,
Strasbourg, France.

## Data Availability

The original data underlying this article will be shared on reasonable request
to the corresponding author.

## References

  * Anderson et al. (2012) Anderson L. D., et al., 2012, A&A, 542, A10
  * Arnal et al. (2008) Arnal E. M., Duronea N. U., Testori J. C., 2008, A&A, 486, 807
  * Arthur & Hoare (2006) Arthur S. J., Hoare M. G., 2006, ApJS, 165, 283
  * Beard et al. (1969) Beard M., Thomas B. M., Day G. A., 1969, Australian Journal of Physics Astrophysical Supplement, 11, 27
  * Beltrán et al. (2006) Beltrán M. T., Brand J., Cesaroni R., Fontani F., Pezzuto S., Testi L., Molinari S., 2006, A&A, 447, 221
  * Bertoldi & McKee (1992) Bertoldi F., McKee C. F., 1992, ApJ, 395, 140
  * Bessell & Brett (1988) Bessell M. S., Brett J. M., 1988, PASP, 100, 1134
  * Bik et al. (2005) Bik A., Kaper L., Hanson M. M., Smits M., 2005, A&A, 440, 121
  * Billot et al. (2010) Billot N., Noriega-Crespo A., Carey S., Guieu S., Shenoy S., Paladini R., Latter W., 2010, ApJ, 712, 797
  * Bronfman et al. (1996) Bronfman L., Nyman L.-A., May J., 1996, A&AS, 115, 81
  * Caswell & Haynes (1987) Caswell J. L., Haynes R. F., 1987, A&A, 171, 261
  * Cheng et al. (2018) Cheng Y., Tan J. C., Liu M., Kong S., Lim W., Andersen M., Da Rio N., 2018, ApJ, 853, 160
  * Churchwell et al. (2006) Churchwell E., et al., 2006, ApJ, 649, 759
  * Contreras et al. (2016) Contreras Y., Garay G., Rathborne J. M., Sanhueza P., 2016, MNRAS, 456, 2041
  * Cyganowski et al. (2003) Cyganowski C. J., Reid M. J., Fish V. L., Ho P. T. P., 2003, ApJ, 596, 344
  * Das et al. (2017) Das S. R., Tej A., Vig S., Liu H.-L., Liu T., Ishwara Chandra C. H., Ghosh S. K., 2017, MNRAS, 472, 4750
  * Das et al. (2018) Das S. R., Tej A., Vig S., Liu T., Ghosh S. K., Chandra C. H. I., 2018, A&A, 612, A36
  * Dyson & Williams (1980) Dyson J. E., Williams D. A., 1980, Physics of the interstellar medium
  * Elmegreen & Lada (1977) Elmegreen B. G., Lada C. J., 1977, ApJ, 214, 725
  * Fazio et al. (2004) Fazio G. G., et al., 2004, ApJS, 154, 10
  * Foster et al. (2011) Foster J. B., et al., 2011, ApJS, 197, 25
  * Griffin et al. (2010) Griffin M. J., et al., 2010, A&A, 518, L3
  * Haslam et al. (1982) Haslam C. G. T., Salter C. J., Stoffel H., Wilson W. E., 1982, A&AS, 47, 1
  * Heiderman et al. (2010) Heiderman A., Evans II N. J., Allen L. E., Huard T., Heyer M., 2010, ApJ, 723, 1019
  * Hildebrand (1983) Hildebrand R. H., 1983, QJRAS, 24, 267
  * Hoare et al. (1991) Hoare M. G., Roche P. F., Glencross W. M., 1991, MNRAS, 251, 584
  * Isequilla et al. (2019) Isequilla N. L., Fernández-López M., Benaglia P., Ishwara-Chandra C. H., del Palacio S., 2019, A&A, 627, A58
  * Issac et al. (2019) Issac N., Tej A., Liu T., Varricatt W., Vig S., Ishwara Chandra C. H., Schultheis M., 2019, MNRAS, 485, 1775
  * Jackson et al. (2013) Jackson J. M., et al., 2013, Publ. Astron. Soc. Australia, 30, e057
  * Ji et al. (2012) Ji W.-G., Zhou J.-J., Esimbek J., Wu Y.-F., Wu G., Tang X.-D., 2012, A&A, 544, A39
  * Kauffmann & Pillai (2010) Kauffmann J., Pillai T., 2010, ApJ, 723, L7
  * Kauffmann et al. (2008) Kauffmann J., Bertoldi F., Bourke T. L., Evans II N. J., Lee C. W., 2008, A&A, 487, 993
  * Kauffmann et al. (2013) Kauffmann J., Pillai T., Goldsmith P. F., 2013, ApJ, 779, 185
  * Koornneef (1983) Koornneef J., 1983, A&A, 500, 247
  * Krumholz & McKee (2008) Krumholz M. R., McKee C. F., 2008, Nature, 451, 1082
  * Kurtz (2002) Kurtz S., 2002, in Crowther P., ed., Astronomical Society of the Pacific Conference Series Vol. 267, Hot Star Workshop III: The Earliest Phases of Massive Star Birth. p. 81 (arXiv:astro-ph/0111351)
  * Lada & Adams (1992) Lada C. J., Adams F. C., 1992, ApJ, 393, 278
  * Lada et al. (2010) Lada C. J., Lombardi M., Alves J. F., 2010, ApJ, 724, 687
  * Lefloch & Lazareff (1994) Lefloch B., Lazareff B., 1994, A&A, 289, 559
  * Mac Low et al. (1991) Mac Low M.-M., van Buren D., Wood D. O. S., Churchwell E., 1991, ApJ, 369, 395
  * Marcote et al. (2015) Marcote B., Ribó M., Paredes J. M., Ishwara-Chandra C. H., 2015, MNRAS, 451, 59
  * Martín-Hernández et al. (2005) Martín-Hernández N. L., Vermeij R., van der Hulst J. M., 2005, A&A, 433, 205
  * Martins et al. (2005) Martins F., Schaerer D., Hillier D. J., 2005, A&A, 436, 1049
  * Marton et al. (2017) Marton G., et al., 2017, arXiv e-prints, p. arXiv:1705.05693
  * Menon et al. (2020) Menon S. H., Federrath C., Kuiper R., 2020, MNRAS, 493, 4643
  * Meyer et al. (1997) Meyer M. R., Calvet N., Hillenbrand L. A., 1997, AJ, 114, 288
  * Molinari et al. (2008) Molinari S., Pezzuto S., Cesaroni R., Brand J., Faustini F., Testi L., 2008, A&A, 481, 345
  * Ortega et al. (2013) Ortega M. E., Paron S., Giacani E., Rubio M., Dubner G., 2013, A&A, 556, A105
  * Paladini et al. (2012) Paladini R., et al., 2012, ApJ, 760, 149
  * Parker et al. (2005) Parker Q. A., et al., 2005, MNRAS, 362, 689
  * Paron et al. (2011) Paron S., Petriella A., Ortega M. E., 2011, A&A, 525, A132
  * Paron et al. (2017) Paron S., Celis Peña M., Ortega M. E., Fariña C., Petriella A., Rubio M., Ashley R. P., 2017, MNRAS, 470, 4662
  * Poglitsch et al. (2010) Poglitsch A., et al., 2010, A&A, 518, L2
  * Rieke & Lebofsky (1985) Rieke G. H., Lebofsky M. J., 1985, ApJ, 288, 618
  * Rodgers et al. (1960) Rodgers A. W., Campbell C. T., Whiteoak J. B., 1960, MNRAS, 121, 103
  * Roger et al. (1999) Roger R. S., Costain C. H., Landecker T. L., Swerdlyk C. M., 1999, A&AS, 137, 7
  * Roman-Lopes & Abraham (2006) Roman-Lopes A., Abraham Z., 2006, AJ, 131, 951
  * Sánchez-Monge et al. (2013) Sánchez-Monge Á., Beltrán M. T., Cesaroni R., Fontani F., Brand J., Molinari S., Testi L., Burton M., 2013, A&A, 550, A21
  * Sanhueza et al. (2019) Sanhueza P., et al., 2019, ApJ, 886, 102
  * Schmiedeke et al. (2016) Schmiedeke A., et al., 2016, A&A, 588, A143
  * Sewilo et al. (2004) Sewilo M., Watson C., Araya E., Churchwell E., Hofner P., Kurtz S., 2004, ApJS, 154, 553
  * Spitzer (1978) Spitzer L., 1978, Physical processes in the interstellar medium, doi:10.1002/9783527617722.
  * Swarup (1991) Swarup G., 1991, Giant metrewave radio telescope (GMRT). pp 376–380
  * Tang et al. (2019) Tang Y.-W., Koch P. M., Peretto N., Novak G., Duarte-Cabral A., Chapman N. L., Hsieh P.-Y., Yen H.-W., 2019, ApJ, 878, 10
  * Tapia et al. (2014a) Tapia M., Persi P., Roth M., Elia D., Molinari S., Saldaño H. P., Gómez M., 2014a, MNRAS, 437, 606
  * Tapia et al. (2014b) Tapia M., Persi P., Roth M., Elia D., Molinari S., Saldaño H. P., Gómez M., 2014b, MNRAS, 438, 2716
  * Tej et al. (2006) Tej A., Ojha D. K., Ghosh S. K., Kulkarni V. K., Verma R. P., Vig S., Prabhu T. P., 2006, A&A, 452, 203
  * Tenorio-Tagle (1979) Tenorio-Tagle G., 1979, A&A, 71, 59
  * Urquhart et al. (2014) Urquhart J. S., et al., 2014, MNRAS, 443, 1555
  * Veena et al. (2016) Veena V. S., Vig S., Tej A., Varricatt W. P., Ghosh S. K., Chandrasekhar T., Ashok N. M., 2016, MNRAS, 456, 2425
  * Walsh et al. (1997) Walsh A. J., Hyland A. R., Robinson G., Burton M. G., 1997, MNRAS, 291, 261
  * Wang et al. (2014) Wang K., et al., 2014, MNRAS, 439, 3275
  * Watson et al. (2008) Watson C., et al., 2008, ApJ, 681, 1341
  * Watson et al. (2010) Watson C., Hanspal U., Mengistu A., 2010, ApJ, 716, 1478
  * Whitworth et al. (1994) Whitworth A. P., Bhattal A. S., Chapman S. J., Disney M. J., Turner J. A., 1994, MNRAS, 268, 291
  * Williams & McKee (1997) Williams J. P., McKee C. F., 1997, ApJ, 476, 166
  * Wright et al. (1994) Wright A. E., Griffith M. R., Burke B. F., Ekers R. D., 1994, ApJS, 91, 111
  * Yang et al. (2021) Yang A. Y., et al., 2021, A&A, 645, A110
  * van Buren & Mac Low (1992) van Buren D., Mac Low M.-M., 1992, ApJ, 394, 534


2022

[1]J. Eduardo Méndez-Delgado

[1]Astronomisches Rechen-Institut, Zentrum für Astronomie der Universität
Heidelberg, Mönchhofstraße 12-14, Heidelberg, D-69120, Baden-Württemberg,
Germany

2]Instituto de Astrofísica de Canarias, (IAC), Vía Láctea, 1, San Cristóbal de
La Laguna, E-38205, Santa Cruz de Tenerife, Spain

3]Departamento de Astrofísica, Universidad de La Laguna, Astrofísico Francisco
Sánchez, s/n., San Cristóbal de La Laguna, E-38206, Santa Cruz de Tenerife,
Spain

4]Instituto de Astronomía, Universidad Nacional Autónoma de México, Apartado
Postal 70-264, Coyoacán, 04510, Mexico City, Mexico

# Temperature inhomogeneities in Mrk 71 can not be discarded

César Esteban  Jorge García-Rojas  Kathryn Kreckel  Manuel Peimbert  *  [  [
[

In a very recent work, Chen:2023  claim that the scenario of temperature
inhomogeneities proposed by Peimbert:1967  ($t^{2}>0$) is not able to explain
the O2+/H+ abundance discrepancy observed between the calculations based on
the optical [O III] collisional excited lines (CELs) and the O II
recombination lines (RLs) in the star forming galaxy Mrk 71. In this work, we
show that conclusions of Chen:2023  depend on several assumptions on the
absolute flux calibration, reddening correction and the adopted electron
density. In fact, using the data of Chen:2023  in a different way and even
considering their 1$\sigma$ uncertainties, it is possible to reach the
opposite conclusion, consistent with $t^{2}=0.097^{+0.008}_{-0.009}$.
Therefore, the existence of temperature inhomogeneities causing the O2+/H+
abundance discrepancy in Mrk 71 can not be ruled out.

Chen:2023  tested the presence of temperature inhomogeneities in the star
forming galaxy Mrk 71. To carry out their analysis, Chen:2023  used an optical
spectrum from the Keck Cosmic Web Imager (KCWI) at the W. M. Keck Observatory
as well as IR spectra from the Far Infrared Field-Imaging Line Spectrometer
(FIFI-LS) at the Stratospheric Observatory for Infrared Astronomy (SOFIA) and
from the Photodetector Array Camera and Spectrometer (PACS) at the Herschel
Space Observatory.

By comparing the different H I line flux ratios with the theoretical
predictions, they infer a reddening constant $c(\text{H}\beta)=0.09\pm 0.04$,
considering the reddening curve of Cardelli:1989  with $R_{V}=3.1$. They
derive the electron density ($n_{e}$) of the gas with three indicators: [O II]
$\lambda 3726/\lambda 3729$, [O III] $\lambda 52\mu\text{m}/\lambda
88\mu\text{m}$ and the O II V1 RL multiplet. On the other hand, they derive
the $T_{e}$ by considering [O III] $\lambda 4363/\lambda 4959$ as well as [O
III] $\lambda 4959/\lambda 52\mu\text{m}$ and $\lambda 4959/\lambda
88\mu\text{m}$. From the comparison of the O2+/H+ abundance derived both with
optical [O III] CELs and O II RLs and assuming that the abundance discrepancy
(AD) is produced by temperature variations, they infer a $t^{2}\sim 0.1$ (see
Eq. (12) from Peimbert:1967 ). This result would imply that both the derived
temperature from [O III] $\lambda 4959/\lambda 52\mu\text{m}$ and $\lambda
4959/\lambda 88\mu\text{m}$ should be $\sim 3000\text{ K}$ lower than what is
obtained from [O III] $\lambda 4363/\lambda 4959$. However, Chen:2023  found a
good consistency between their calculations of $T_{e}$ based on [O III]
$\lambda 4959/\lambda 52\mu\text{m}$, $\lambda 4959/\lambda 88\mu\text{m}$ and
$\lambda 4363/\lambda 4959$ and therefore they claim the absence of
significant temperature fluctuations in Mrk 71.

The results obtained by Chen:2023  are highly dependent on the accuracy of the
absolute flux calibration between the three instruments, since there are no H
I detections in the IR data that could be used to normalize the spectra.
Considering that the KCWI observations were taken under non-photometric
conditions Chen:2023 , that the FIFI-LS observations present telluric features
Sutter:2022  and that the PACS observations were carried out in the “un-
chopped” mode and show detector response variations Fadda:2016 , the absolute
flux calibration between the three different kinds of data is not
straightforward. In fact, the comparison of [C II] $\lambda 158\mu\text{m}$,
detected both in FIFI-LS and PACS reveals a difference of $\sim 15\%$ between
the flux calibrated data of both instruments even after the PACS detector
response variations correction. Possible systematic differences between the
optical and IR-spectra are not analyzed or quantified by Chen:2023 .

The difference between the FIFI-LS and PACS spectra implies the existence of a
systematic bias in the flux of at least one of the [O III] IR CELs used.
Dividing the flux difference of $\sim 15\%$ quadratically and including it in
the uncertainty bars does not properly treat the systematic error, as it
impacts differently [O III] $\lambda 4959/\lambda 52\mu\text{m}$ than [O III]
$\lambda 4959/\lambda 88\mu\text{m}$, given their different $n_{\rm
e}$-dependence. Considering that both [O III] IR CELs have similar fluxes, a
more robust way to reduce the impact of the flux bias is to use the sum of
their fluxes instead, comparing $T_{\rm e}$ derived from [O III] $\lambda
4959/\lambda\lambda 52+88\mu\text{m}$ with the value obtained using $\lambda
4363/\lambda 4959$. The flux systematic difference also calls into question
the density derived from [O III] $\lambda 52\mu\text{m}/\lambda 88\mu\text{m}$
with the reported absolute fluxes.

In Fig. 1 we present the resulting plasma diagnostics considering the line
fluxes and uncertainties reported by Chen:2023  under three values of
$c(\text{H}\beta)$, all consistent within the reported 1$\sigma$. In all
cases, the reddening curve of Cardelli:1989  with $R_{V}=3.1$ was used. The
atomic data used were the default ones from PyNeb Luridiana:2015  in its
version 1.1.16, the same ones assumed by Chen:2023 . As shown in Fig. 1,
depending on the $c(\text{H}\beta)$ and the $n_{e}$ adopted, it is possible to
be consistent either with the absence of temperature inhomogeneities (as
Chen:2023  concluded) or the opposite case, predicted by
$t^{2}=0.097^{+0.008}_{-0.009}$.

The adoption of a density value close to $n_{e}$([O II] $\lambda 3726/\lambda
3729$)=$160\pm 10\text{ cm}^{-3}$ instead of the available $n_{e}$(O
II)=$310\pm 50\text{ cm}^{-3}$ was not justified by Chen:2023 . Both values
could be considered typical within the range of densities found in previous
studies of Mrk 71 Gonzalez-Delgado:1994 ; Esteban:2002 ; Esteban:2009 . The
most evident problem with the $n_{e}$ value derived from [O II] $\lambda
3726/\lambda 3729$ by Chen:2023  is that panels (c) and (d) of their Fig. 2
show that the spectral resolution of KCWI is insufficient to have at least a
partial separation of the [O II] doublet. Therefore, this density value
strongly depends on the assumptions imposed on the Gaussian deblend required
to measure the lines separately with uncertainty bars of $\sim 0.6\%$.

Figure 1: Using the data presented by Chen:2023  it is possible to show either
the absence of temperature inhomogeneities or the opposite case. Each panel
shows a plasma diagnostic plot considering a slightly different reddening
constant, $c(\text{H}\beta)$, all consistent within the 1$\sigma$
uncertainties. According to the $t^{2}$-paradigm proposed by Peimbert:1967 ,
in the presence of temperature inhomogeneities, [O III] $\lambda
4959/\lambda\lambda 52+88\mu\text{m}$ (green band) should be lower than [O
III] $\lambda 4363/\lambda 4959$ (red band), matching the value predicted by
the O II RLs (black band). Depending on $c(\text{H}\beta)$ and the adopted
electron density ($n_{e}$), it is possible to argue either against the
existence of temperature fluctuations (as concluded by Chen:2023 ) or the
opposite case.

In order to reinforce their arguments on the absence of strong density and
temperature fluctuations in Mrk 71, Chen:2023  mention that in a sample of
regions, among which Mrk 71 does not appear, Mingozzi:2022  found differences
of only $\sim 1000\text{ K}$ between [O III] $\lambda 1666/\lambda 5007$ and
[O III] and $\lambda 4363/\lambda 5007$ and that this is not sufficient for
the temperature fluctuations scenario to explain the observed AD factor.
However, such an assertion can be proven to be incorrect. Considering Eq. 15
from Peimbert:1967 : $T_{e}(\text{[O\,III]}\lambda 1666/\lambda
5007)-T_{e}(\text{[O\,III]}\lambda 4363/\lambda 5007)[\text{K}]=12330\times
t^{2}$. For the most common value of $t^{2}\sim 0.04$ found for the star
forming regions garcia_rojas:2007 ; Peimbert:2012 , the temperature difference
would be of $\sim 500\text{ K}$. A very extreme value of
$t^{2}=0.097^{+0.008}_{-0.009}$, would imply a difference of $1200\pm
110\text{ K}$. This exercise demonstrates that the temperature differences
found by Mingozzi:2022  can comprise typical and extreme values of $t^{2}$.

We conclude that based on the data and the analysis presented by Chen:2023 ,
one cannot be conclusive on the presence or absence of temperature
inhomogeneities in Mrk 71, since both interpretations are possible even within
their estimated 1$\sigma$ uncertainties. To be conclusive in this regard, it
is necessary to consider the possible presence of density variations that
could introduce systematic biases on the $n_{e}$ diagnostics even if the line
intensity ratios are well measured Rubin:1989 . $n_{e}$-biases could affect
determinations of O2+/H+ based on [O III] $\lambda\lambda 52+88\mu\text{m}$ in
a much higher extent than those based on O II V1. Observational evidence of
temperature and density inhomogeneities in star-forming regions (including Mrk
71) is presented by mendez2023a  and mendez2023b , respectively.

Acknowledgments JEM-D thanks the help provided by V. Gomez-Llanos in managing
the assignment of colors in the PyNeb plasma diagnostics and to O. V. Egorov
for fruitful discussions.

Authors’ contributions JEM-D lead the analysis and writing of the manuscript.
CE, JG-R, KK and MP provided critical feedback and modified the text.

Conflict of interest/Competing interests

The authors declare that they have no competing financial interests.

Data availability

All the data discussed here was presented by Chen:2023 .

Code availability Our results use the PyNeb code, publicly available on
GitHub. https://github.com/Morisset/PyNeb_devel

Funding

JEM-D and KK gratefully acknowledge funding from the Deutsche
Forschungsgemeinschaft (DFG, German Research Foundation) in the form of an
Emmy Noether Research Group (grant number KR4598/2-1, PI Kreckel). CE and JG-R
acknowledge support from the Agencia Estatal de Investigación del Ministerio
de Ciencia e Innovación (AEI-MCINN) under grant Espectroscopía de campo
integral de regiones H II locales. Modelos para el estudio de regiones H II
extragalácticas with reference 10.13039/501100011033 and support under grant
P/308614 financed by funds transferred from the Spanish Ministry of Science,
Innovation and Universities, charged to the General State Budgets and with
funds transferred from the General Budgets of the Autonomous Community of the
Canary Islands by the MCIU. JG-R acknowledges support from an Advanced
Fellowship under the Severo Ochoa excellence program CEX2019-000920-S and
financial support from the Canarian Agency for Research, Innovation and
Information Society (ACIISI), of the Canary Islands Government, and the
European Regional Development Fund (ERDF), under grant with reference
ProID2021010074.

Additional Information

Correspondence should be addressed to JEM-D<EMAIL_ADDRESS>
## References

  *   * (1) Chen, Y. _et al._ Accurate Oxygen Abundance of Interstellar Gas in Mrk 71 from Optical and Infrared Spectra.  _arXiv e-prints_ arXiv:2304.09898 (2023).  10.48550/arXiv.2304.09898, arXiv:2304.09898 [astro-ph.GA].
  * (2) Peimbert, M.  Temperature Determinations of H II Regions.  _ApJ_ 150, 825 (1967).  10.1086/149385 .
  * (3) Cardelli, J. A., Clayton, G. C. & Mathis, J. S.  The Relationship between Infrared, Optical, and Ultraviolet Extinction.  _ApJ_ 345, 245 (1989).  10.1086/167900 .
  * (4) Sutter, J. & Fadda, D.  [C II] Map of the Molecular Ring and Arms of the Spiral Galaxy NGC 7331.  _ApJ_ 926 (1), 82 (2022).  10.3847/1538-4357/ac4252, arXiv:2112.05706 [astro-ph.GA].
  * (5) Fadda, D., Jacobson, J. D. & Appleton, P. N.  Transient effects in Herschel/PACS spectroscopy.  _A &A_ 594, A90 (2016).  10.1051/0004-6361/201527612, arXiv:1601.07729 [astro-ph.IM].
  * (6) Luridiana, V., Morisset, C. & Shaw, R. A.  PyNeb: a new tool for analyzing emission lines. I. Code description and validation of results.  _A &A_ 573, A42 (2015).  10.1051/0004-6361/201323152, arXiv:1410.6662 [astro-ph.IM].
  * (7) Gonzalez-Delgado, R. M. _et al._ Violent Star Formation in NGC 2363.  _ApJ_ 437, 239 (1994).  10.1086/174992 .
  * (8) Esteban, C., Peimbert, M., Torres-Peimbert, S. & Rodríguez, M.  Optical Recombination Lines of Heavy Elements in Giant Extragalactic H II Regions.  _ApJ_ 581 (1), 241–257 (2002).  10.1086/344104, arXiv:astro-ph/0208313 [astro-ph].
  * (9) Esteban, C. _et al._ Keck HIRES Spectroscopy of Extragalactic H II Regions: C and O Abundances from Recombination Lines.  _ApJ_ 700 (1), 654–678 (2009).  10.1088/0004-637X/700/1/654, arXiv:0905.2532 [astro-ph.CO].
  * (10) Mingozzi, M. _et al._ CLASSY IV. Exploring UV Diagnostics of the Interstellar Medium in Local High-z Analogs at the Dawn of the JWST Era.  _ApJ_ 939 (2), 110 (2022).  10.3847/1538-4357/ac952c, arXiv:2209.09047 [astro-ph.GA].
  * (11) García-Rojas, J. & Esteban, C.  On the Abundance Discrepancy Problem in H II Regions.  _ApJ_ 670 (1), 457–470 (2007).  10.1086/521871, arXiv:0707.3518 [astro-ph].
  * (12) Peimbert, A., Peña-Guerrero, M. A. & Peimbert, M.  A Classification of H II Regions Based on Oxygen and Helium Lines: The Cases of TOL 2146-391 and TOL 0357-3915.  _ApJ_ 753 (1), 39 (2012).  10.1088/0004-637X/753/1/39, arXiv:1204.4507 [astro-ph.CO].
  * (13) Rubin, R. H.  The Effect of Density Variations on Elemental Abundance Ratios in Gaseous Nebulae.  _ApJS_ 69, 897 (1989).  10.1086/191330 .
  * (14) Méndez-Delgado, J. E., Esteban, C., García-Rojas, J., Kreckel, K. & Peimbert, M.  Temperature inhomogeneities cause the abundance discrepancy in H II regions.  _Nature_ 618 (7964), 249–251 (2023).  10.1038/s41586-023-05956-2, arXiv:2305.11578 [astro-ph.GA].
  * (15) Méndez-Delgado, J. E. _et al._ Density biases and temperature relations for DESIRED H II regions.  _MNRAS_ 523 (2), 2952–2973 (2023).  10.1093/mnras/stad1569, arXiv:2305.13136 [astro-ph.GA].