11institutetext: Instituto de Astrofísica de La Plata (UNLP - CONICET), La
Plata, Argentina. 11email<EMAIL_ADDRESS>22institutetext:
Instituto de Astronomía y Física del Espacio (UBA, CONICET), CC 67, Suc. 28,
1428 Buenos Aires, Argentina 33institutetext: Departamento de Astronomía,
Universidad de Chile, Casilla 36, Santiago de Chile 44institutetext:
Observatório do Valongo, Universidade Federal do Rio de Janeiro, Ladeira Pedro
Antônio, 43, Rio de Janeiro - RJ, 20080-090, Brazil 55institutetext: Facultad
de Ciencias Astronómicas y Geofísicas, Universidad Nacional de La Plata, Paseo
del Bosque s/n, 1900 La Plata, Argentina 66institutetext: Instituto Argentino
de Radioastronomía (CCT-La Plata, CONICET; CICPBA), C.C. No. 5, 1894,Villa
Elisa, Argentina

#  Physical characterization of S169: A prototypical IR bubble associated with
the massive star-forming region IRAS 12326-6245

N. U. Duronea Member of the Carrera del Investigador Científico of CONICET,
Argentina.11 S. Cichowolski⋆ 22 L. Bronfman 33 E. Mendoza 44 R. Finger 33 L.
A. Suad⋆ 22 M. Corti⋆ 5566 E. M. Reynoso⋆ 22

(Received July 31; accepted December 4)

###### Abstract

Context.

Aims.  With the aim of studying the physical properties of Galactic IR bubbles
and to explore their impact in massive star formation, we present a study of
the IR bubble S169, associated with the massive star forming region IRAS
12326-6245.

Methods.  We used CO(2-1),13CO(2-1), C18O(2-1), HCN(3-2), and HCO+(3-2) line
data obtained with the APEX telescope using the on-the-fly full sampling
technique to study the properties of the molecular gas in the nebula and the
IRAS source . To analyze the properties and distribution of the dust, we made
use of images obtained from the IRAC-GLIMPSE, Herschel, and ATLASGAL archives.
The properties of the ionized gas in the nebula were studied using radio
continuum and H$\alpha$ images obtained from the SUMSS survey and SuperCOSMOS
database, respectively. In our search for stellar and protostellar objects in
the region, we used point source calalogs obtained from the MSX, WISE,
GLIMPSE, 2MASS, AAVSO, ASCC-2.5V3, and GAIA databases.

Results. The new APEX observations allowed us to identify three molecular
components, each one associated with different regions of the nebula, namely:
at $-$39 km s-1(component A), $-$25 km s-1(component B), and $-$17 km
s-1(component C). Component A is shown to be the most dense and clumpy. Six
molecular condensations (MC1 to MC6) were identified in this component, with
MC3 (the densest and more massive one) being the molecular counterpart of IRAS
12326-6245. For this source, we estimated an H2 column density up to
8$\times$1023 cm-2. An LTE analysis of the high density tracer lines HCO+(3-2)
and HCN(3-2) on this source, assuming 50 and 150 K, respectively, indicates
column densities of $N$(HCO+) = (5.2 $\pm$ 0.1) $\times$ 1013 cm-2 and
$N$(HCN) = (1.9 $\pm$ 0.5) $\times$ 1014 cm-2. To explain the morphology and
velocity of components A, B, and C , we propose a simple model consisting of a
partially complete semisphere-like structure expanding at $\sim$ 12 km s-1.
The introduction of this model has led to a discussion about the distance to
both S169 and IRAS 12326-6245, which was estimated to be $\sim$ 2 kpc. Several
candidate YSOs were identified, projected mostly onto the molecular
condensations MC3, MC4, and MC5, which indicates that the star-formation
process is very active at the borders of the nebula. A comparison between
observable and modeled parameters was not enough to discern whether the
collect-and-collapse mechanism is acting at the edge of S169. However, other
processes such as radiative-driven implosion or even a combination of both
mechanisms, namely, collect-and-collapse and radiative-driven implosion, could
be acting simultaneously in the region.

###### Key Words.:

ISM: molecules, ISM: IR dust bubbles, ISM: individual object: S169,
ISM:individual object: IRAS 12326-6245 .

##  1 Introduction

Massive stars modify the interstellar medium (ISM), where they are born
through their stellar winds and UV photons, creating interstellar bubbles and
Hii regions. These nebulae are detected in the optical and radio continuum
ranges, showing the presence of ionized gas, and at infrared (IR), millimeter,
and submillimeter wavelengths, indicating the presence of both dust (at
different temperatures) and molecules.

Milky Way surveys at IR wavelengths have provided a plethora of examples of
bubbles/Hii regions, where their physical properties can be obtained on
pc/sub-pc scales. Making use of the 8.0 $\mu$m Spitzer-GLIMPSE survey,
Benjamin et al. (2003) and Churchwell et al. (2006, 2007) identified more than
600 candidates for interstellar IR dust bubbles between Galactic longitudes of
$-$60∘ to 60∘. Follow-up observations of the stellar and prestellar
populations around some of these bubbles have begun to link them with sites of
recent triggered massive star formation (e.g., Zavagno et al. 2006; Watson et
al. 2008; Deharveng et al. 2008, 2009; Samal et al. 2014; Kendrew et al. 2016;
Duronea et al. 2017). These kinds of studies provide important insights into
the evolution of interstellar bubbles, the characteristics of the interstellar
medium where they evolve, the role of massive stars on favoring or suppressing
new generations of stars, and the physical conditions under which massive
stars may be induced to form.

Figure 1: Composite image of S169 and its environs showing the (IRAC-GLIMPSE)
emissions at 8.0 (red) and 5.8 $\mu$m (green). The superCOSMOS H${\alpha}$
emission is in blue. White contours depict the continuum emission at 870
$\mu$m (ATLASGAL) arising from IRAS 12326-6245.

As part of a project aimed at characterizing and studying the physical
properties of Galactic IR bubbles and their surroundings, and to better
understand their influence in massive star formation, we selected a southern
field containing S169 (Churchwell et al., 2006), which is a prototypical IR
bubble-shaped nebula of about 8 arcmin in diameter, centered at RA, Dec.
(J2000)= (12h35m29$\aas@@fstack{s}$2,
$-$62∘58′03$\aas@@fstack{\prime\prime}$8). Figure 1 shows a composite image of
S169 in 8.0 and 5.8 $\mu$m (IRAC-GLIMPSE) and the H${\alpha}$ (superCOSMOS)
emissions. The presence of emission in the 8-$\mu$m band, mostly dominated by
the emission of polycyclic aromatic hydrocarbons molecules (PAHs) at 7.7 and
8.6 $\mu$m, indicates the existence of a photodissociation region (PDR)
probably created by a nearby stellar source. Since PAHs molecules are
destroyed inside the ionized gas of an Hii region (Povich et al., 2007), but
fluorescing when irradiated with weak ultraviolet radiation, they indicate the
limits of the ionization front and delineate the boundaries of the bubble
nebula, tracing the distribution of the parental molecular gas in which S169
is developing. This scenario is also supported by the presence of H${\alpha}$
emission at the center of the bubble nebula. A distinctive feature in the
emission of S169 is the presence of two bright borders at the southern region
of the IR nebula, which are discernible in Fig. 1 in yellow tonalities. The
spatial distribution of the IR emission gives a good hint of the distribution
of the molecular gas and the location of the densest regions, very likely
associated with the two bright borders described before. As can be seen from
Fig.1, the IR emission at 8.0 $\mu$m toward the west of the cavity is much
weaker and composed of a faint arc-like filament approximately at RA =
12h35m10s.

The nebula S169 is placed near the IR source IRAS 12326-6245, which has
characteristics of a young stellar object (YSO), according to the photometric
criteria from Junkes et al. (1992). Substantial evidence suggests that IRAS
12326-6245 is a dense molecular core where massive star formation is taking
place, namely: a) this source was observed in the CS(2-1) line at a velocity
of $-$39.4 km s-1 (Bronfman et al., 1996; Osterloh et al., 1997). It was also
observed by Zinchenko et al. (2000) in several lines of HNCO at frequencies
between 290 and 806 GHz. This indicates undoubtedly the existence of dense
molecular gas associated with the source; b) emission from dense dust has been
previously detected at 1.2mm as a compact source (Faúndez et al., 2004; Hill
et al., 2005; Miettinen et al., 2006). The source is also a strong emitter at
870 $\mu$m (Dedes et al. 2011; König et al. 2017; see Fig. 1); c) CH3OH maser
emission, a common phenomenon in massive star forming regions, was reported by
Caswell (2009) at a velocity of $-$39.8 km s-1. OH and H2O maser emission was
also detected at the source (Caswell, 1998; MacLeod et al., 1998); d) a
variety of complex organic molecules have been detected in the source (Araya
et al., 2005; Dedes et al., 2011). These kinds of molecules are mostly formed
in the surface of dust grains and later evaporate into the gas phase. This is
characteristic of “hot molecular cores” (HMC; Kurtz et al. 2000), which are
better known to be associated with early stages of massive star formation; e)
molecular line maps indicate the presence of a very powerful bipolar outflow
at the center of the source (Henning et al., 2000).

In this work, we present a complete multifrequency analysis of the ISM around
S169 using high spatial resolution molecular observations, which are necessary
to establish the excitation conditions in the molecular environs, as well as
to estimate their physical properties and to investigate the kinematics of the
region. The analysis of the molecular observations are complemented with an
analysis of a multifrequency archival dataset, which altogether contribute to
our understanding of the impact of stellar feedback and how the ionized gas
interacts with its molecular environment. We also investigate the status of
the star-formation processes in the region. In that context, the proximity
between S169 and IRAS 12326-6245 is crucial for the analysis since it will
allow us to better understand the high-mass star-formation process around IR
dust bubbles. An analysis of dense molecular line tracers and submillimeter
emission in IRAS 12326-6245 is also presented.

The source IRAS 12326-6245 is believed to be placed at a distance of 4.4 kpc.
Such distance was first estimated by Osterloh et al. (1997) and consistently
adopted afterwards by several authors (e.g., Henning et al. 2000; Faúndez et
al. 2004; Dedes et al. 2011). The distance to IRAS 12326-6245 is also matter
of some debate in this work.

##  2 Observations and databases

###  2.1 Molecular observations

The molecular observations presented in this paper were made in December 2016
with the Atacama Pathfinder EXperiment111This publication is based on data
acquired with the Atacama Pathfinder Experiment (APEX) under program ID
C-098.F-9801. APEX is a collaboration between the Max-Planck-Institut fur
Radioastronomie, the European Southern Observatory, and the Onsala Space
Observatory (APEX) telescope (Güsten et al., 2006) at Llano de Chajnantor
(Chilean Andes). As the front end for the observations, we used the APEX-1
receiver of the Swedish Heterodyne Facility Instrument (SHeFI; Vassilev et al.
2008). The back end for all observations was the eXtended bandwidth Fast
Fourier Transform Spectrometer2 (XFFTS2), with a 2.5 GHz bandwidth divided
into 32768 channels. The observed transitions and basic observational
parameters are summarized in Table 1. Calibration was done by the chopper-
wheel technique and the output intensity scale given by the system is $T_{\rm
A}$, which represents the antenna temperature corrected for atmospheric
attenuation. The observed intensities were converted to the main-beam
brightness temperature scale by $T_{\rm mb}$ = $T_{\rm A}$/$\eta_{\rm mb}$,
where $\eta_{\rm mb}$ is the main beam efficiency. For the APEX-1 receiver we
adopted $\eta_{\rm mb}$ = 0.75.

Observations were made using the on-the-fly (OTF) mode with a step size of
10′′ (full sampling) in two orthogonal scan directions along RA and
Dec.(J2000). For the lines CO(2-1), 13CO(2-1), and C18O(2-1) the observed
region was of $\sim$14${}^{\prime}\times$14′ in size centered on RA, Dec.
(J2000)= (12h35m44s, $-$63∘00′08″). For the lines HCN(3-2) and HCO+(3-2) the
observed region was $\sim$4${}^{\prime}\times$4′ in size centered on the
position of the source IRAS 12326-6245 at RA, Dec. (J2000)= (12h35m33s,
$-$63∘02′56″). The data reduction was carried out with the CLASS90 package of
the IRAM GILDAS software222http://www.iram.fr/IRAMFR/GILDAS.

Table 1: Observational parameters for the observed transitions.  molecular | Frequency | Beam | Velocity | rms
---|---|---|---|---
transition | (GHz) | (′′) | resolution | noise
|  |  | (km s-1) | (K)
CO(2-1) | 230.538000 |  $\sim$27 | 0.3 |  $\sim$0.3
13CO(2-1) | 220.398677 |  $\sim$28 | 0.3 |  $\sim$0.3
C18O(2-1) | 219.560357 |  $\sim$28 | 0.3 |  $\sim$0.25
HCN(3-2) | 265.886180 |  $\sim$23 | 0.3 |  $\sim$0.15
HCO+(3-2) | 267.557526 |  $\sim$23 | 0.3 |  $\sim$0.15

###  2.2 Archival data

Infrared data: a) Images of ATLASGAL at 870 $\mu$m (345 GHz) (Schuller et al.,
2009). This survey covers the inner Galactic plane, $l$ = 300∘ to 60∘, $|b|$
$\leq$ 1.∘5, with the root mean square (rms) noise in the range of 0.05 - 0.07
Jy beam-1. The beam size at 870 $\mu$m is 19$\aas@@fstack{\prime\prime}$2\. b)
Images from the Herschel333 Herschel is an ESA space observatory with science
instruments provided by European-led Principal Investigator consortia and with
important participation from NASA
(http://www.cosmos.esa.int/web/herschel/science-archive) Infrared GALactic
(Hi-GAL) plane survey key program (Molinari et al., 2010). We used images from
the photometric array camera and spectrometer (PACS) survey at 70 and 100
$\mu$m, with FWHM of 5$\aas@@fstack{\prime\prime}$5 and 11″, respectively, and
from the spectral and photometric imaging receiver (SPIRE) at 350 $\mu$m with
a FHWM of 25″. c) Images from Spitzer at 5.8 and 8.0 $\mu$m from the Galactic
Legacy Infrared Mid-Plane Survey Extraordinaire (Spitzer-GLIMPSE, Benjamin et
al. 2003), retrieved from the Spitzer Science
Center444http://scs.spitzer.caltech.edu. The images have a spatial resolution
of $\sim$ 2″.

Radio continuum data: Mosaics obtained from the Sydney University Molonglo Sky
Survey555http://www.astrop.physics.usyd.edu.au/mosaics (SUMSS; Bock et al.,
1999). This widefield radio imaging survey covers the southern sky at 843 MHz.
The resolution is $43^{\prime\prime}\times 43^{\prime\prime}cosec(\delta)$ and
the rms noise level is $\sim 1$ mJy beam-1.

Hi data: Data cubes from the Southern Galactic Plane Survey (SGPS; McClure-
Griffiths et al., 2005). This survey includes the regions limited by Galactic
longitude $253^{\circ}\leq l\leq 358^{\circ}$ (SGPS-1) and $5^{\circ}\leq
l\leq 20^{\circ}$ (SGPS-2) and Galactic latitude $|b|\leq
1$$.\\!\\!{}^{\circ}$5\. Interferometric observations obtained with the
Australia Telescope Compact Array (ATCA) are combined with single dish data
from the 64 m Parkes radiotelescope. The final products have an angular
resolution of $\sim$ 2′ and a rms sensitivity of $\sim 1$ mJy beam-1.

Narrow band H$\alpha$ data: Retrieved from the SuperCOSMOS H$\alpha$
Survey666http://www-wfau.roe.ac.uk/sss/halpha/hapixel.html (SHS). The images
have a spatial resolution of $\sim$1′′ (Parker et al., 2005).

To investigate ionizing stars candidates and YSO candidates in the region of
the nebula, we made use of the IRAS Point Source
Catalog777https://irsa.ipac.caltech.edu/Missions/iras.html (Beichman et al.,
1988), the MSX Infrared Point Source
Catalog888https://irsa.ipac.caltech.edu/Missions/msx.html (Egan et al., 2003),
the WISE All-Sky Source
Catalog999https://irsa.ipac.caltech.edu/Missions/wise.html (Wright et al.,
2010), the GLIMPSE point source
catalog101010https://irsa.ipac.caltech.edu/Missions/spitzer.html (Benjamin et
al., 2003), the Two Micron All Sky Survey
(2MASS)111111https://irsa.ipac.caltech.edu/Missions/2mass.html (Skrutskie et
al., 2006), the American Association of Variable Star Observers (AAVSO)
Photometric All-Sky SurveyA (APASS)121212https://www.aavso.org/apass (Henden
et al., 2010), the ASCC-2.5 V3: All-sky Compiled Catalog of 2.5 million stars
(Kharchenko & Roeser, 2009), and the Astrometric catalog GAIA Data Release 2
131313https://www.cosmos.esa.int/web/gaia/home (DR2; Gaia Collaboration et al.
(2016, 2018)).

Figure 2:  Total averaged spectra of CO(2-1) (red), 13CO(2-1) (green), and
C18O(2-1) (blue) obtained within a region of $\sim$14${}^{\prime}\times$14′
centered on RA, Dec. (J2000)= (12h35m44s, $-$63∘00′08″).  Figure 3:  Spatial
distribution of the C18O(2-1) emission (left panel), 13CO(2-1) emission
(middle panel), and CO(2-1) emission (right panel) in the velocity interval
from $\sim$ –42 km s-1 to –37 km s-1 (component A). In the left panel, the
contours are 0.35 ($\sim$ 3 rms), 0.7, 1.15, 2, 3, 5, and 7 K km s-1. In the
middle panel the contour levels are 0.5 ($\sim$ 6 rms), 1, 1.5, 2, 2.5, 3, 4,
5, 7, 9, 11, 13, and 15 K km s-1. In the right panel the contour levels are
1.5 ($\sim$ 20 rms), 2, 2.5, 3, 4, 5, 6, 8, 10, 14, 18, 22, and 26 K km s-1
($T_{\rm mb}$). Identified molecular condensations are indicated in yellow
over the 13CO(2-1) emission.

##  3 Molecular and dust emission

###  3.1 Emission of CO and its isotopologs

Figure 2 portrays the total averaged spectra of CO(2-1), 13CO(2-1), and
C18O(2-1) in a region of $\sim$14${}^{\prime}\times$14′ in a position that is
approximately at the center of the bubble nebula. The CO emission clearly
depicts a strong narrow velocity component peaking at $\sim$ –39 km s-1 and
two weaker broad velocity components peaking at $\sim$ –25 km s-1 and $\sim$
–17 km s-1. For the sake of the analysis, these components will be hereafter
referred to as component A, component B, and component C, respectively.

The peak velocity of component A, which is very likely dominated by the
emission of IRAS 12326-6245, is almost coincident with the peak velocity of
reported non-masing and masing emission lines towards the IRAS source
(Bronfman et al., 1996; Zinchenko et al., 2000; Henning et al., 2000; Dedes et
al., 2011; Araya et al., 2005; Caswell, 1998, 2009; MacLeod et al., 1998).
This component is also detected in 13CO and C18O emissions, which indicates
that it is composed of high density molecular gas. Components B and C, on the
other hand, are barely detected in the 13CO emission and not detected in the
C18O emission. As opposed to the case of the velocity of component A, neither
masing nor non-masing emission lines were reported for IRAS 12326-6245 at the
velocity of components B and C.

###  3.2 Component A

In Fig. 3, we show the spatial distribution of the CO(2-1), 13CO(2-1), and
C18O(2-1) lines in the velocity interval from $\sim$ $-$42 km s-1 to $-$37 km
s-1. As expected according to the IR emission distribution (see Fig. 1), the
bulk of the molecular emission appears concentrated toward the eastern and
southern borders of the IR nebula, with no significant molecular emission
detected toward the western and northeastern regions. The morphology and
location of component A, H$\alpha$ emission (see Fig. 1), and radio continuum
emission (see Sect. 4) suggest that the ionized gas is expanding against the
molecular cloud in the eastern and southern borders. The molecular emission at
the southern border of the bubble seems to be tracing the densest areas, which
are probably sculpted by the action of the Hii region. The morphology and
location of component A with respect to the near IR emission is similar to
many other in Galactic IR bubbles found in the literature (e.g., Zavagno et
al. 2006; Deharveng et al. 2008, 2009; Anderson et al. 2015; Liu et al. 2015;
Cappa et al. 2016; Duronea et al. 2017; Devine et al. 2018). For a better
visualization and comparison with the IR emission, in panel a of Fig. 4 we
show the emission of 13CO(2-1) of component A superimposed on the IRAC-GLIMPSE
emissions at 8.0 and 5.8 $\mu$m.

Figure 4: Panel a: Emission distribution of the 13CO(2-1) line in the total
velocity range from –37 to –42 km s-1 (white contours and blue color scale)
superimposed on the IRAC-GLIMPSE images at 8.0 $\mu$m and 5.8 $\mu$m (red and
green color scales, respectively). The contours levels are the same than those
shown in the middle panel of Fig.3. Panels b to i: Channel maps of the
13CO(2-1) line emission in velocity intervals of 0.7 km s-1 (white contours)
in the velocity range from –36.7 km s-1to –42.5 km s-1 superimposed on the 8.0
$\mu$m and 5.8 $\mu$m (IRAC-GLIMPSE) emissions (red and green color scales)
and 870 $\mu$m ATLASGAL emission (blue color scale) inside the 3 rms limit.
Contours levels are 0.9 ($\sim$ 5.5 rms), 2.6, 4.4, 6.2, 7.9, 10, 12, 16, and
20 K km s-1. The velocity interval is indicated in the top right corner of
each panel.

From Fig. 3, it can also be seen that the CO emission of component A is more
extended than that of 13CO and C18O. This indicates that 13CO and C18O lines,
optically thinner than that of CO (see below), are actually tracing the
distribution of the densest molecular gas surrounding the bubble nebula. In
order to perform a more detailed analysis of the dense molecular gas in
component A, we have roughly identified six molecular condensations in the
13CO(2-1) emission using the C18O(2-1) emission (which traces the denser gas)
as a reference. The identification was done using the 1.5 K km s-1 contour
level in the 13CO(2-1) emission. The condensations are labeled in Fig. 3 as
MC1, MC2, MC3, MC4, MC5, and MC6. These condensations are the places where
star formation is likely to be taking place (see Sect. 5). In panels b to i of
Fig. 4, we show the channel maps of 13CO in the velocity range from $-$36.7 km
s-1 to $-$42.5 km s-1 in intervals of 0.7 km s-1. In order to compare the
molecular emission distribution with the warm and cold dust emission, the 13CO
emission was superimposed on the 8.0 and 5.8 $\mu$m (IRAC-GLIMPSE) and the 870
$\mu$m (ATLASGAL) maps. Figure 4 shows that MC3, located at RA, Dec. (J2000)=
(12h35m35s, $-$63∘02′25″), is the only molecular condensation that is present
in the entire velocity range. The position and size of MC3 perfectly match
with those of the bright spot seen at 870 $\mu$m. This condensation certainly
represents the carbon monoxide counterpart of IRAS 12326-6245, which reveals
the physical association between S169 and the IRAS source. In the velocity
interval from $-$36.7 to $-$38.2 km s-1 condensation MC1 becomes noticeable,
reaching its peak temperature in the velocity interval from $-$38.9 to $-$39.6
km s-1, close to a 8.0 $\mu$m source seen at RA, Dec. (J2000)= (12h35m24s,
$-$63∘05′42″). It can be detected till a velocity of $-$40.4 km s-1. For the
case of condensation MC2, it can be first noticed in the velocity interval
from $-$36.7 to $-$38.9 km s-1, where its emission merges with MC1. Its peak
emission is observed in the velocity range from $-$37.4 to $-$38.2 km s-1 and
is coincident with another 8.0 $\mu$m source seen at RA, Dec. (J2000)=
(12h35m29s, $-$63∘04′25″). Condensation MC4 becomes detectable in the velocity
interval from $-$36.7 to $-$37.4 km s-1, reaching to its maximum peak emission
in the velocity interval from $-$38.2 to $-$38.9 km s-1. It can be noticed
till a velocity of $-$39.6 km s-1 although its emission appears blended with
MC2. Regarding condensations MC5 and MC6, they become detectable in the
velocity range from $-$36.7 to $-$37.4 km s-1. Their emission distribution is
quite irregular and achieve their peak temperatures in the velocity range from
$-$38.2 to $-$38.9 km s-1. Condensation MC4 is detectable till a velocity of
$-$39.6 km s-1, while MC5 till a velocity of $-$40.4 km s-1.

Figure 4 also shows a strong morphological correlation between the brightest
parts of the PDR (described in Sect. 1) and the molecular condensations MC3,
MC4, and MC5, especially at velocities between –38.2 to –38.9 km s-1. In
particular, it is noticeable how the bright emission feature detected at 5.8
and 8.0 $\mu$m borders the area of MC3, which are likely to be exposed to
stellar radiation.

To estimate the physical properties of the 13CO condensations, we obtained six
13CO emission maps (not shown here) integrated in the velocity interval in
which each condensation is detected. The deconvolved effective radius of the
condensations derived from the 13CO line, $R_{\rm D}^{\rm{}^{13}CO}$, is
calculated as:

$R_{\rm D}^{\rm{}^{13}CO}=\sqrt{R^{2}_{\rm eff}-({\rm HPBW}/2)^{2}},$ (1)

where $R_{\rm eff}$ is the effective radii of the condensation ($R_{\rm eff}$
= $\sqrt{A_{\rm cond}/\pi}$), being $A_{\rm cond}$ the area of the
condensation, and HPBW is the half-power beam width of the instrument. To
estimate $R_{\rm D}^{\rm{}^{13}CO}$ and other parameters, we adopted a
distance of 2.03${}^{+0.77}_{-0.61}$ kpc (see Sect. 6.2). Then we obtained the
effective radii of 0.40, 0.27, 0.94, 0.49, 0.75, and 0.86 pc for MC1, MC2,
MC3, MC4, MC5, and MC6, respectively.

Assuming that all rotational levels are thermalized with the same excitation
temperature (LTE) and that the emission is optically thick, we derived the
excitation temperature of each molecular condensation, $T_{\rm exc}$, from the
CO(2-1) line using

$T_{\rm peak}=T_{12}^{*}\left[\left(e^{\frac{T_{12}^{*}}{T_{\rm
exc}}}-1\right)^{-1}-\left(e^{\frac{T_{12}^{*}}{T_{\rm
bg}}}-1\right)^{-1}\right],$ (2)

where $T_{\rm 12}^{*}$ = $h\nu_{\rm 12}/k$, being $\nu_{\rm 12}$ the frequency
of the 12CO(2-1) line, and $T_{\rm bg}=2.7$ K. To obtain the peak main beam
temperature of the CO(2-1), 13CO(2-1), and C18O(2-1) lines $(T_{\rm peak}),$
we used Gaussian fits over the spectra profiles obtained integrating in the
direction of the peak emission of the 13CO condensations (in all the cases
coincident with those of the C18O emission) over an area equivalent to the
area of the beam. It is worth to point out that a single Gaussian component
was considered for MC3, which shows self-absorption in the spectrum of CO(2-1)
at the velocity of the peak emission of 13CO(2-1) and C18O(2-1) (Henning et
al., 2000; Dedes et al., 2011). Excitation temperatures derived with Eq. 2 are
then 21.3 K, 33.8 K, 47.8 K, 21.2 K, and 16.0 K, and 16.8 K for MC1, MC2, MC3,
MC4, MC5, and MC6, respectively.

Table 2: Observed parameters obtained from the emission lines and physical properties derived for the molecular condensations MC1 to MC6 | $T_{\rm peak}$ | $\int T_{\rm mb}\ dv$ |  $\Delta$v | $\tau$ |  $N(^{13}$CO) |  $N$(C18O) | $N({\rm H_{2}})$ | $M({\rm H_{2}})$
---|---|---|---|---|---|---|---|---
| (K) | (K km s-1) | (km s-1) |  | (cm-2) | (cm-2) | (cm-2) | ($M_{\odot}$)
| CO 13CO C18O |  13CO C18O | CO 13CO | CO 13CO C18O |  |  |  |
MC1 | 16.1 6.8 0.8 | 15.9 0.7 | 2.75 1.49 | 16.51 0.54 0.05 | 2.8$\times$1016 | 1.0$\times$1015 | (2.0/0.6)$\times$1022 | 300
MC2 | 28.4 10.0 1.3 | 7.8 1.0 | 2.85 1.91 | 16.32 0.43 0.04 | 1.6$\times$1016 | 1.6$\times$1015 | (1.2/0.9)$\times$1022 | 90
MC3 | 42.2 23.3 7.1 | 32.8 8.1 | 7.51 5.85 | 36.2 0.83 0.19 | 9.8$\times$1016 | 1.8$\times$1016 | (1.0/0.7)$\times$1023 | 5500
MC4 | 15.9 7.4 2.1 | 9.9 1.2 | 2.83 2.16 | 26.38 0.61 0.13 | 1.8$\times$1016 | 1.1$\times$1015 | (1.3/0.7)$\times$1022 | 280
MC5 | 10.9 6.9 1.6 | 6.1 0.8 | 3.18 1.59 | 27.76 0.98 0.15 | 1.2$\times$1016 | 1.1$\times$1015 | (8.5/6.5)$\times$1021 | 420
MC6 | 11.7 8.4 1.3 | 7.1 1.1 | 2.55 1.65 | 44.36 1.21 0.11 | 1.6$\times$1016 | 1.5$\times$1015 | (1.2/0.9)$\times$1022 | 720

The optical depths ($\tau$) of the 13CO(2-1) and C18O(2-1) lines, denoted in
the next equations as $\tau^{13}$ and $\tau^{18}$, respectively, were obtained
assuming that the excitation temperature is the same for CO(2-1), 13CO(2-1),
and C18O(2-1), and using the expressions

$\tau^{13}=-{\rm ln}\left[1-\frac{T_{\rm
peak}{\rm{}^{13}CO}}{T_{13}^{*}}\left[\left(e^{\frac{T_{13}^{*}}{T_{\rm
exc}}}-1\right)^{-1}-\left(e^{\frac{T_{13}^{*}}{T_{\rm
bg}}}-1\right)^{-1}\right]^{-1}\right]$ (3)

and

$\tau^{18}=-{\rm ln}\left[1-\frac{T_{\rm peak}{\rm
C^{18}O}}{T_{18}^{*}}\left[\left(e^{\frac{T_{18}^{*}}{T_{\rm
exc}}}-1\right)^{-1}-\left(e^{\frac{T_{18}^{*}}{T_{\rm
bg}}}-1\right)^{-1}\right]^{-1}\right],$ (4)

where $T_{\rm 13}^{*}$ = $h\nu_{\rm 13}/k$, $T_{\rm 18}^{*}$ = $h\nu_{\rm
18}/k$, with $\nu_{\rm 13}$ and $\nu_{\rm 18}$ the frequencies of the
13CO(2-1) and C18O(2-1) lines, respectively. We can also estimate the optical
depth of the CO(2-1) line from the 13CO(2-1) line using

$\quad\tau^{12}=\ \left[\frac{\nu^{13}}{\nu^{12}}\right]^{2}\
\times\left[\frac{\Delta{\rm v}^{13}}{\Delta{\rm v}^{12}}\right]\
\times\left[\frac{\rm CO}{{}^{13}{\rm CO}}\right]\ \tau^{13},$ (5)

where CO/13CO is the isotope ratio (assumed to be $\sim$ 62; Langer & Penzias
1993); $\Delta{\rm v}^{13}$ and $\Delta{\rm v}^{12}$ are defined as the full
width half maximum (FWHM) of the spectra of the 13CO and CO lines,
respectively, which are derived by using a single Gaussian fitting (FWHM = 2
$\times$ $\sqrt{2\ ln2}$ $\times$ $\sigma_{\rm gauss}$).

In LTE, the 13CO and C18O column densities can be estimated from the 13CO(2-1)
and C18O(2-1) line using

$N({\rm{}^{13}CO})=3.23\times 10^{14}\left[\frac{e^{\frac{T_{13}^{*}}{T_{\rm
exc}}}}{1-e^{-\frac{T_{13}^{*}}{T_{\rm exc}}}}\right]T_{\rm exc}\int\tau^{13}\
dv\ \ \textrm{ (cm${}^{-2}$)}$ (6)

and

$N({\rm C^{18}O})=3.21\times 10^{14}\left[\frac{e^{\frac{T_{18}^{*}}{T_{\rm
exc}}}}{1-e^{-\frac{T_{18}^{*}}{T_{\rm exc}}}}\right]T_{\rm exc}\int\tau^{18}\
dv\ \ \textrm{ (cm${}^{-2}$)}.$ (7)

The integral in Eqs. 6 and 7 can be approximated by

$\ \ T_{\rm exc}\int{\tau\ dv\approx\ \frac{\tau}{1-e^{(-\tau)}}\int{T_{\rm
mb}}}\ \ d{\rm v}.$ (8)

This approximation helps to eliminate to some extent optical depth effects and
is good within 15% for $\tau<$ 2 (Rohlfs & Wilson 2004). Bearing in mind the
values of $\tau^{\rm 13}$ obtained, the approximation is appropriate for our
region. Estimations of the H2 column density, $N$(H2), were obtained from both
the ${}^{13}{\rm CO}$ and C18O column densities, and adopting abundances
$X$[13CO] = 7.1$\times$105 and $X$[C18O] = 5.9$\times$106 (Frerking et al.,
1982). For the optical depth and column density, we estimated uncertainties of
about 20 $\%$ and 25 $\%$, respectively, arising mostly from calibration
uncertainties of about 20 $\%$ (Dumke & Mac-Auliffe, 2010).

The total hydrogen mass of each clump was calculated using

$\ \ {\rm M(H_{2})}\ =\ (m_{\rm sun})^{-1}\ \mu\ m_{\rm H}\ \Omega\ {\it
N}(\rm H_{2})\ {\it d}^{2}\quad\quad\quad\textrm{(M${}_{\odot}$)},$ (9)

where $m_{\rm sun}$ is the solar mass ($\sim$2$\times$1033 g), $\mu$ is the
mean molecular weight, which is assumed to be equal to 2.76 after allowing for
a relative helium abundance of 25% by mass (Yamaguchi et al., 1999), $m_{\rm
H}$ is the hydrogen atom mass ($\sim$1.67$\times$10-24 g), $\Omega$ is the
solid angle of the 13CO emission expressed in sr, $d$ is the adopted distance
expressed in cm, and $N$(H2) is the H2 column density obtained from
$N({\rm{}^{13}CO})$. Uncertainties in molecular masses are about 60 $\%$, and
originate mainly from a distance uncertainty of about 30 $\%$ (see Sect. 6.2).
The derived physical parameters are presented in Table 2.

In order to perform a study of the molecular gas in IRAS 12326-6245, we also
present an analysis for MC3 using the high-density tracer lines HCN(3-2) and
HCO+(3-2). In Fig. 5, we show the emission of these lines in the velocity
intervals from $-$47.6 to $-$29.2 km s-1 and $-$44.8 to $-$32.5 km s-1,
respectively. Their emissions show a good morphological correspondence with
that at 870 $\mu$m (see Fig. 4) and seem to be delineating an obscured region
in the 8.0 and 5.8 $\mu$m emissions. They are certainly the HCN and HCO+
counterparts of the molecular condensation MC3. In order to estimate the HCN
and HCO+ column densities, we used LTE models numerically approximated with
the Markov chain Monte Carlo (MCMC) method (Foreman-Mackey et al., 2013). The
lines in LTE were simulated and compared with the observations (e.g., Comito
et al. 2005; Möller et al. 2017). The solutions presented were obtained from
the averaged spectra considering different offsets since the individual
solutions for each one of them were similar. That is a consequence of the
homogeneous emission traced by the species analyzed here. In Fig. 5, we show
the average spectra (in antenna temperature) of HCN(3-2) and HCO+(3-2). The
results presented here were obtained by taking into account those reported in
IRAS 12326-6245 by Dedes et al. (2011), who assumed temperature values of 150
K and 50 K for the gas conditions of species in an apparent hot core region
and extended envelope, respectively. For the sake of the analysis, we also
included the optically thin lines of 13CO(2-1) and C18O(2-1). Then, assuming
an excitation temperature of $T_{\rm exc}$=50 K for the HCO+(3-2) line, we
obtained a column density of $N$(HCO+) = (5.2 $\pm$ 0.1) $\times$ 1013 cm-2.
For the 13CO(2-1) and C18O(2-1) lines, also adopting $T_{\rm exc}$=50 K, we
obtained $N$(13CO) = (2.01 $\pm$ 0.04) $\times$ 1017 cm-2 and $N$(C18O) = (3.9
$\pm$ 0.1) $\times$ 1016 cm-2. For the case of the HCN(3-2) line, LTE models
suggest a physical condition related with the hot core region. Based on the
hypothesis stating that $T_{\rm exc}$=150 K, we derived $N$(HCN) = (1.9 $\pm$
0.5) $\times$ 1014 cm-2. Concerning the uncertainties, the results presented
here are approximations that should be considered (at least) within the
calibration uncertainty since several assumptions were adopted for single
transitions. In order to inspect the quality of the fits, in Fig. 5 we also
exhibit the residuals after subtracting the emission lines from their LTE
models.

Figure 5: Upper panel: Spectral line and LTE model of the HCN(3-2) line
obtained from LTE-MCMC calculations assuming $T_{\rm exc}$=150 K (see text),
with a resulting density of $N$=(1.9 $\pm$ 0.5) $\times$ 1014 cm-2. The blue
line indicates the residual obtained from the subtraction between the
observation and LTE model. In colors: Spatial distribution of the HCN(3-2)
emission in the velocity interval from $-$47.6 to $-$29.2 km s-1(color blue
and white contours) superimposed on the IRAC 8 and 5.8 $\mu$m emissions (red
and green colors). Lower panel: Same as above for the HCO+(3-2) line, assuming
$T_{\rm exc}$=50 K, with a resulting density of $N$=(5.2 $\pm$ 0.1) $\times$
1013 cm-2. The HCO+(3-2) emission distribution is in the velocity interval
from $-$44.8 to $-$32.5 km s-1. The contour levels for the HCN(3-2) are 0.25
($\sim$5 rms), 0.85, 2.2, and 4.0 K km s-1; and for the HCO+(3-2) line, they
are 0.28 ($\sim$ 5 rms), 0.8, 2.4, and 5 K km s-1

###  3.3 Components B and C

As mentioned in Sect. 3.1, two weaker broad velocity molecular components
peaking approximately at $-$25 km s-1 (component B) and $-$17 km s-1
(component C) were also detected in S169. These components are hardly detected
in the 13CO emission and not detected at all in the C18O emission, which
certainly indicates that they are mostly composed by low-density gas.

Figure 6:  Upper panel: Spatial distribution of the CO emission (blue color
scale and white contours) in the velocity interval from $-$28.9 km s-1 to
$-$18.9 km s-1 (component B) superimposed on the IRAC 8 $\mu$m and 5.8 $\mu$m
emissions (red and green color scales). The contour levels start from 0.95 K
km s-1 ($\sim$ 9 rms) in steps of 0.5 K km s-1. Lower panel: Spatial
distribution of the CO emission in the velocity interval from $-$18.2 km s-1
to $-$12.7 km s-1 (component C). The contour levels start from 0.85 K km s-1
($\sim$ 7 rms) in steps of 0.4 K km s-1.

In order to analyze their spatial distribution, we constructed two CO emission
maps integrated in the velocity intervals from $-$28.9 to $-$18.9 km s-1
(component B) and from $-$18.2 to $-$12.7 km s-1 (component C). They are shown
in Fig. 6 superimposed onto the 8.0 and 5.8 $\mu$m emissions. For the case of
component B, three molecular features are seen projected towards different
regions of the nebula. The brightest feature (hereafter, feature 1) is seen
approximately centered at RA, Dec. (J2000)= (12h36m34s, $-$62∘55′13″)
perfectly delineating the northeastern and northern borders of the IR nebula.
As can be seen from the figure, the brightest regions of the molecular gas
appear projected onto the faintest regions of the IR emission, which suggests
that a PDR was formed over the surface of feature 1 (viewed from the side).
Another feature is seen centered approximately at RA, Dec. (J2000)=
(12h36m40s, $-$63∘06′00″) (feature 2). This feature appears bordering the
southeastern edge of the nebula, also projected onto the faintest regions of
the IR nebula. The third feature is seen projected along the center of the
nebula approximately at Dec=$-$62∘58′00″ and is still noticeable at more
negative velocities. Since this feature shows no morphological correspondence
with any region of the IR nebula, a physical link with S169 cannot be
suggested. Probably, it is an unconnected molecular structure located behind
the nebula since no absorption features can be noticed along the brightest
regions of the IR emission. Regarding component C, only one molecular
structure (feature 3) stands out, peaking approximately at RA, Dec. (J2000)=
(12h35m05s, $-$62∘58′00″). This feature shows a good morphological
correspondence with the faint IR arclike filament seen at the western section
of the nebula at RA = 12h35m10s. This correspondence is also well-displayed in
Fig. 10. Another structure has been noted at the northeastern region of the
nebula, however it represents the remains of feature 1 at slightly more
positive velocities.

Table 3: Physical properties estimated for the CO feature 1, feature 2, and feature 3 | $\Omega$ | $\int{T_{\rm mb}({\rm CO})}\ d{\rm v}$ | $N({\rm H_{2}})$ | $M(\rm H_{2})$
---|---|---|---|---
Feature | (10-6 sr) | (K km s-1) | (1020cm-2) | ($M_{\odot}$)
1 | 2.4 | 14.9 | 2.4 | 52
2 | 1.3 | 12.3 | 2.0 | 24
3 | 2.5 | 7.9 | 1.3 | 29

Since no discernible emission of the 13CO(2-1) line is detected for features
1, 2, and 3, in order to estimate their column density and mass, we use the
relation between the H2 integrated column density and the CO integrated
emission:

$N({\rm H_{2}})\ =\ X\ \times\int{T_{\rm mb}({\rm CO})}\ d{\rm v},$ (10)

where $X$ is an empirical factor that has been shown to be roughly constant
for the 12CO(1-0) line in Galactic molecular clouds and lies in the range of
(1 - 3) $\times$ 1020 cm-2 (K km s-1)-1, as estimated by the virial theorem
and $\gamma$-ray emission (Bloemen et al., 1986; Solomon et al., 1987; Digel
et al., 1996; Strong & Mattox, 1996). In this paper, we adopt $X$ = 1.6
$\times$ 1020 cm-2 (K km s-1)-1 (Hunter et al., 1997). The integrated emission
$\int{T_{\rm mb}({\rm CO})}\ d{\rm v}$ was calculated adopting an area
determined by the first contour level indicated in Fig. 6. For the column
density, we estimated an uncertainty of about 50 $\%$, arising mostly from the
factor $X$ and calibration uncertainties. The mass of the features were then
estimated using Eq. 9. The obtained parameters are shown in Table 3. It is
worthwhile to point out that the area and mass obtained for features 1 and 2
are likely lower limits since these features seem to be part of a larger
structure (extending beyond the area covered in our APEX observations), which
probably corresponds to the parental molecular cloud in which the IR bubble
was formed.

###  3.4 Submillimeter dust emission

In the upper panel of Fig. 7, we present a three-color composite image of S169
as seen by Herschel images. The 70 and 160 $\mu$m emissions (red and green
color scales, respectively) underline the emission of the warmer dust close to
the ionized gas. As can be seen from the figure, the emission at 70 $\mu$m has
a good morphological correspondence with the IRAC 8 $\mu$m emission (black
contours), although the latter exhibits a sharper edge towards the center of
the nebula (see also Fig.1), which is likely due to the destruction of the PAH
molecules by the ionization front. The 70 $\mu$m emission, in contrast, is
more diffuse (partly because of the lower angular resolution) and extends
towards the center of the nebula, which indicates that it is also tracing the
warm dust that still remains mixed with the ionized gas. The same feature can
be observed in the IRAC-GLIMPSE 24 $\mu$m emission (not shown here). The 350
$\mu$m emission (blue color scale) seems to be underlining the emission from
the dust in the outer parts of the nebula. It is very likely that the emission
distribution at this wavelength is mostly tracing the distribution of the cold
dust immersed in the parental molecular cloud placed at the eastern border of
the nebula, which is partially disclosed by features 1 and 2 (see Fig. 6).

Figure 7: Upper panel: Three-color composite image of S169: Herschel PACS 70
and 160 $\mu$m emission are in red and green, respectively. Herschel SPIRE 350
$\mu$m emission is in blue. Black contours underline de IRAC-GLIMPSE 8 $\mu$m
emission at 35, 55, 75, 105, 150, 170, 300, and 500 MJy sr-1. Middle panel:
Dust temperature map derived from the 70 and 160 $\mu$m emission. The color
temperature scale is on the right. Lower panel: Column density map obtained
from the 160 $\mu$m emission. Dashed black contours indicate the location of
molecular concentrations MC1 to MC6 for orientation.

To study the distribution of the dust temperature in the nebula, we
constructed a temperature map from the ratio of the observed fluxes in two
Herschel bands. Since a high level of radiative feedback from powering massive
stars is usually observed in Hii regions/IR bubbles, it is reasonable to
assume that the molecular gas and dust are not too cold. Then we used the
ratio of the 70 $\mu$m versus 160 $\mu$m maps, which are better suited to
measure color temperatures up to $\sim$ 80 K. Furthermore, the angular
resolution of the temperature map obtained using these bands is the highest
resolution achievable from the Herschel maps. The 70 $\mu$m map was smoothed
down to the angular resolution of the 160 $\mu$m data. Then the temperature
map was constructed as the inverse function of the ratio 70 to 160 $\mu$m maps
that is, $T_{\rm c}=f^{-1}_{(T)}$. Assuming a dust emissivity following a
power law of $\kappa_{\nu}$ $\propto$ $\nu^{\beta}$, with ${\beta}$ as the
spectral index of the thermal dust emission (assumed to be $\beta$ = 2), in
the optically thin thermal dust emission regime, $f_{(T)}$ takes the
parametric form:

$\qquad
f_{(T)}=\frac{S_{70}}{S_{160}}=\frac{B(70,T)}{B(160,T)}\left(\frac{70}{160}\right)^{\beta},$
(11)

where $B(70,T)$ and $B(160,T)$ make up the blackbody Planck function for a
temperature $T$ at wavelengths 70 and 160 $\mu$m, respectively. The assumption
that the emission is optically thin is further justified from the derived
column densities in the nebula (see Preibisch et al. 2012 for details). From
the obtained map (middle panel of Fig. 7), dust temperatures between $\sim$ 20
and 30 K can be discerned. These temperatures are commonly observed in the
closest regions of IR bubble nebulae (e.g., Anderson et al. 2012; Duronea et
al. 2015; Figueira et al. 2017). Warmer dust is seen in the region of the
denser molecular condensations (MC1, MC2, MC3, and MC4) and the border of the
PDR, particularly on two spots at RA, Dec. (J2000)= (12h35m24s, $-$63∘00′56″)
and RA, Dec. (J2000)= (12h35m35s, $-$63∘02′29″) where the temperature achieves
$\sim$ 30 K. The first spot is coincident with a bright rim placed over the
southern region of the PDR and a peak in the radio continuum emission (see
Fig.8 in Sect. 4), which suggests external heating of the dust. The second
spot is coincident with the molecular condensation MC3 (IRAS12326-6245) which
is very likely to indicate that the dust is also internally heated. Not
surprisingly, lower temperatures (12 - 18 K) are seen in the region of feature
3 and behind the PDR, coincident with the position of features 1 and 2.

With the temperature map, we converted the surface brightness map at 160
$\mu$m into a beam averaged column density map using (Hildebrand, 1983):

$\qquad N({\rm H_{2}})\ =\ R\ \ \frac{S_{\rm 160}}{\Omega_{\rm beam}\
\kappa_{160}\ \mu\ m_{\rm H}\ B(160,T_{\rm dust})},$ (12)

where $\Omega_{\rm beam}$ is the beam solid angle ($\pi$ $\theta_{\rm
HPBW}^{2}$ / 4 ln 2), $R$ is the gas-to-dust ratio (assumed to be 100), and
$\kappa_{160}$ is the dust opacity per unit mass at 160 $\mu$m assumed to be
40.5 cm2 g-1 (Ossenkopf & Henning, 1994). The resulting map (presented in the
lower panel of Fig. 7) gives an overall view of the density distribution in
the whole region of the nebula. Clearly, denser regions are observed behind
the PDR in the region of component A, especially towards IRAS12326-6245, where
the column density achieves peak values up to $\sim$ 4$\times$1022 cm-2.

As shown in Figs. 1 and 4, the source IRAS 12326-6245 is very bright in the
continuum emission at 870 $\mu$m (no emission at this wavelength is detected
above the 3 rms limit outside the IRAS source). The emission at 870 $\mu$m is
usually optically thin and dominated by the thermal emission from dust
contained in dense material (e.g., dense molecular cores or filaments). Then
we used the image of ATLASGAL to more confidently derive the beam averaged
density in IRAS 12326-6245 using Eq. 12. For the calculations, we adopted
$\kappa_{870}$ = 1.85 cm2 g-1, a value that is representative for relatively
dense molecular clouds (Ossenkopf & Henning, 1994; Henning et al., 1995) and
$T_{\rm dust}$ = 30 K and we obtained a column density of $N({\rm H_{2})}$=8.2
$\times$ 1023 cm-2. This value is almost two orders of magnitude higher than
that observed in the column density map obtained with the emission at 160
$\mu$m, which might indicate that the emission is optically thick at that
wavelength. Furthermore, the temperature used to derive $N({\rm H_{2})}$ in
IRAS 12326-6245 (30 K) might have been underestimated since, as mentioned
above, its emission at 160 $\mu$m might not be optically thin at such high
density. This could led to an overestimation of the column density in the IRAS
source. The column density obtained from the 870 $\mu$m emission is also
somewhat higher than those derived for MC3 in Sect. 3.2 from the 13CO(2-1) and
C18O(2-1) emissions. We do keep in mind, however, that the column density
values derived from continuum IR and carbon monoxide emissions are strongly
dependent on the adopted abundances (dust and carbon monoxide-to H2).

##  4 Ionized gas

In Fig. 8, we show the radio continuum emission map at 843 MHz overlaid on the
8.0 and 5.8 $\mu$m emissions. Three bright structures stand out in the
direction of the nebula. The brightest one (S843 = 985 mJy), which does not
have a discernible counterpart in the H$\alpha$ emission, is observed towards
the norther border of the nebula at RA, Dec. (J2000)=(12h35m39s,
$-$62∘54′26″). Using a radio continuum image from the SGPS (not shown here),
we found that this compact source has an integrated flux density of
S${}_{1384}=0.63$ mJy. Since the spectral index turns out to be $\alpha=-0.9$,
we conclude that this source is extragalactic and will not be taken into
further consideration.

An arc-shaped structure can be discerned towards the center of the nebula,
which may be described by two main components: (i) a southern bright component
with its outer edge closely delineating the PDR depicted by the 8.0 $\mu$m
emission; and (ii) a weaker tail about 5′-long extending from the northern end
of the previous component towards the northwest. A number of artifacts in the
radio continuum emission are obviously produced by sidelobes and grating rings
from the bright extragalactic source, such as a void at RA, Dec.
(J2000)$\approx$(12h35m45s, $-$62∘56′33″) and two radial strips connecting the
compact source and the weak tail. Hence, we believe that the morphology of the
arc-shaped structure is probably distorted in the region closer to the
extragalactic source. However, the morphology of the unaffected part
(southernmost component) suggests that the ionized gas traced by the radio
continuum emission is likely to be expanding against component A and features
1 and 2, supporting the same conclusion derived in Sect. 1 based on the
H$\alpha$ emission.

Figure 8:  Radio continuum emission distribution at 843 MHz (blue color and
white contours), superimposed on the 8 and 5.8 $\mu$m emission (red and green
color). Contours levels are 3.1 ($\sim$ 3 rms), 4.5, 6.3, 7.6, 12, 20, 30, 40,
100, 200, 300, and 400 mJy beam-1.

Using the integrated radio continuum flux obtained for the whole arc-shaped
structure (i.e., the southern bright component and weaker tail), S843 = 308
mJy, we obtained its electron density and ionized mass using (Panagia &
Walmsley, 1978):

$n_{\rm e}=3.113\times 10^{2}\ S_{843}^{0.5}\ T_{\rm e}^{0.25}\ d^{-0.5}\
b(\nu,T_{\rm e})^{-0.5}\ \theta_{R}^{-1.5}\ \ {\rm cm}^{-3}$ (13)

and

$M_{\rm ion}=0.7934\ S^{0.5}_{843}\ T^{0.25}_{\rm e}\ d^{2.5}\ b(\nu,T_{\rm
e})^{-0.5}\ \theta^{1.5}_{R}\ \ (1+Y)^{-1}\ \ {\rm M_{\odot}},$ (14)

where $T_{\rm e}$ is the electron temperature (assumed to be 1 $\times$ 104 K)
in units of 104 K, $\theta_{\rm R}$ is the angular radius of the source in arc
minutes (assumed to be 4′), $b(\nu,T_{\rm e})$ = 1 $+$
0.3195$\times$log($T_{\rm e}$/104 K) - 0.213$\times$log($\nu$/1 GHz), $d$ is
in kpc, and $Y$ is the abundance ratio by number of He+ to H+. Then we
obtained $n_{\rm e}$ = 15 cm-3 and $M_{\rm ion}$ = 17 M⊙ . It is worthwhile
pointing out that the presence of the extragalactic source projected on S169
prevents the estimation of the whole radio continuum emission related to the
structure.

The number of ionizing Lyman continuum photons needed to sustain the current
level of ionization in the arc-shaped structure can be calculated using
(Chaisson, 1976):

$N_{\rm Lyc}=0.76\times\ 10^{47}\ T^{-0.45}_{\rm e}\ S_{843}\ \nu^{0.1}d^{2}\
\ {\rm s}^{-1},$ (15)

where $T_{e}$ is in units of 104 K, $S_{843}$ is in Jy, $\nu$ is in GHz, and
$d$ is in kpc. Then, in considering distance uncertainties, we obtained for
this source a range of $N_{\rm Lyc}$ = 0.5 - 2.0 $\times$ 1047 s-1. We keep in
mind that these values are likely lower limits since about 25 - 50 $\%$ of the
UV photons are absorbed by interstellar dust in the Hii region (Inoue, 2001).
Furthermore, Watson et al. (2008), on the basis of studies of several bubbles,
concluded that $N_{\rm Lyc}$ estimated with this method is lower than expected
by about a factor of 2. Then for the arc-shaped structure, we estimate a range
of $N_{\rm Lyc}$ $\approx$ (1 - 4) $\times$ 1047 s-1. Adopting ionizing photon
rates extracted from Martins et al. (2005), we estimate the spectral type of
the ionizing star of S169 to be at most O9.5V. Alternatively, a handful of
later B-type stars could also be responsible for powering the Hii region. The
third structure detected in the 843 MHz emission is a small bright spot
discernible in the southern region of the nebula at RA, Dec. (J2000)=
(12h35m35s, $-$63∘02′29″), which is certainly the radio continuum counterpart
of IRAS 12326-6245. Using the flux density obtained for this source
($S_{843}$=12 mJy), we obtained $n_{\rm e}$ = 70 cm-3, $M_{\rm ion}$ = 0.3 M⊙
, and $N_{\rm Lyc}$= (0.4 - 1.4) $\times$ 1046 s-1. It is worthwhile pointing
out that for the previous estimations, we assumed optically thin emission at
843 MHz; since that assumption cannot be verified, estimations obtained with
the Eqs. 13, 14, and 15 should be taken as the lower limits. In fact, for the
case of the radio continuum emission associated with IRAS 12326-6245, Dedes et
al. (2011) used cm observations obtained by Urquhart et al. (2007) to derive a
Lyman continuum flux of 8.2$\times$1048 s-1, which is almost three orders of
magnitude larger than the value obtained with Eq. 15 for the source.

In a broad sense, the distribution of the ionized gas inside the cavity
suggests an electron density gradient, with the densest material near the
border of the IR nebula. This configuration suggests that the nebula is
bounded by density to the west, and bounded by ionization to the east, north,
and south. Furthermore, the structure detected to the west of the nebula along
RA $\sim$ 12h34m44s, which appears connected with the eastern arc-shaped
structure within the 3 rms emission limit, might be indicative of a
”champagne-flow” scenario (Tenorio-Tagle, 1979).

##  5 Stellar and protostellar content in the nebula

###  5.1 Ionizing star candidates

With the aim of identifying possible ionizing stars in the nebula, we first
inspected IR catalogs in a region of $\sim$ 7′ in radius centered on RA, Dec.
(J2000)$\approx$ (12h35m28s, $-$62∘58′04″). Making use of GLIMPSE data, we
identified Class III candidates, which are usually referred to as main-
sequence (MS) young stars using color-color diagrams (Allen et al., 2004).
Then we cross-correlated these sources with sources selected from the 2MASS
catalog with photometric uncertainty, $K_{s}<$ 0.05 mag (Ducati et al., 2001),
and we obtained $\sim$ 600 Class III sources with reliable $JHK$ photometry.
Since none of these sources were found to have spectral classification, we
used the parameter $Q_{NIR}$ = $(J-H)-1.7(H-K_{s})$ following Borissova et al.
(2011) to separate sources between early (O-B) MS stars (-0.1$<Q_{NIR}<$0.15)
and stars that reveal IR excess ($Q_{NIR}<$ –0.1) and are likely preMS stars
or YSOs, or both. Then we obtained 60 MS candidates.

Next, we inspected visual catalogs in the same region. A search in the APASS
and ASCC-2.5 V3 catalogs yielded 228 and 10 sources, respectively, both with V
and B magnitudes. We also carried out a search in the astrometric catalog GAIA
DR2 of more than 6000 sources with parallax data in the region of S169.
Thereafter, we cross-correlated all sources found with the APASS, GAIA DR2,
and ASCC-2.5 catalogs, obtaining only 90 sources in common between the first
two catalogs, which includes the 10 sources of the ASCC-2.5 catalog. We then
estimated the distance of visual sources employing the color-magnitude ($B-V$
vs $V$) diagram (see Fig. 9), calibrated with the optical ($BV$) (Schmidt-
Kaler, Th. 1982) to adjust the theoretical MS by a distance of $\sim$ 2 kpc
(see Sect. 6.2) and color excess E$(B-V)$ = 0.8 magnitudes. Thus, we were able
to identify the stars with photometric distance close to the nebula. Among
these sources, only 19 had distances between 1.7 kpc and 2.3 kpc that were
estimated from the parallax catalog from GAIA.

In the next step, we cross-correlated the 60 MS candidates obtained from the
IR data with the sources previously studied with a color magnitude diagram and
we found 47 sources in common. In order to better determine their spectral
types, we inspected their loci in a $(J-H)$ vs. $(H-K)$ diagram adopting the
MS calibration given by Schmidt-Kaler (Th. 1982), Cousins (1978), and
Koornneef (1983), the absorption ratios ($r_{x}$ = $A_{x}/A_{v}$) given by
Rieke & Lebofsky (1985), van der Hulst curve 15, and the color transformations
from 2MASS to Koorneef data of Carpenter et al. (2001).

Finally, we kept only those stars that are projected onto the cavity of the
bubble ($\sim$ 3′ in radius centered on RA, Dec. (J2000)$\approx$(12h35m28s,
$-$62∘58′04″)). This yielded only 10 sources with spectral types O9 (2
sources), B1 (1 source) and B2 or later (7 sources), all of them probable MS
stars; it is worthwhile pointing out that only 5 of these sources belong to
the group of 19 sources with distances between 1.7 kpc and 2.3 kpc obtained
from GAIA. The positions of these sources in the MC diagram and in the cavity
are shown in Fig. 9 and Fig. 10, respectively.

Figure 9: Color magnitude diagram of stars searched in a region centered in
S169. Black dots are young spectral type stars with different reddening values
(see text in Sect. 5). Magenta, green, and yellow crosses represent the O9V,
B1V, and B2V (or later) stars, all of them ionizing candidate stars displayed
with the same colors in Fig. 10. Green and black curves represent the MS
shifted according to the distance modulus adopted with and without reddening,
respectively. Red line indicates the normal reddening path (Rv = 3.1)

###  5.2 Identification of candidate YSOs

Making use of the MSX, WISE and Spitzer point source catalogs, we look for the
primary tracers of stellar formation activity onto the molecular clouds
related to S169.

Figure 10:  MSX, WISE, and Spitzer candidate YSOs projected onto the IRAC-
GLIMPSE image at 8.0 and 5.8 $\mu$m. Cyan diamonds indicate the MSX sources,
blue circles and blue squares indicate WISE Class I and Class II sources,
respectively, and green circles and green squares indicate the Spitzer Class I
and Class II sources, respectively. The numerical references of YSOs are based
on those from Table 4. The size of the symbols do not necessary match the
angular size of the sources. Ionizing candidate stars identified in Sect 5.1
are also indicated by crosses in magenta (O9V), green (B1V), and yellow (B2V
or later). White contours indicate the 13CO emission of component A (first
contour level $\sim$ 8 rms) and the CO emission of feature 1 (first contour
level $\sim$13 rms) , feature 2 (first contour level $\sim$12 rms), and
feature 3 (first contour level $\sim$10 rms).

To this aim, MSX sources were selected satisfying a flux quality Q $>$1 in the
four bands. WISE sources with photometric flux uncertainties $>$ 0.2 mag and
signal-to-noise ratio (S/N) $<$ 7 in the W1, W2, W3, and W4 bands, were
rejected. Finally, we kept Spitzer sources with photometric uncertainties $<$
0.2 mag in the four IRAC bands. Then, within a 14′ box centered at RA, Dec.
(J2000)= (12h35m44s, $-$63∘00′08″), we found a total of 2 MSX, 400 WISE, and
810 Spitzer sources fulfilling the selection criteria mentioned above. To
identify candidate YSOs we adopted the classification scheme described in
Lumsden et al. (2002), Koenig et al. (2012), and Gutermuth et al. (2009) for
the MSX, WISE, and Spitzer sources, respectively. Several sources were found
to qualify the above criteria, which are listed in Table 4. The table presents
the designation of the sources, their coordinates, flux densities, information
concerning the type of the protostellar object, matching with another source
in the table, and coincidence with a molecular or IR structure.

Among the MSX sources, we found one MYSO candidate, which coincides with the
center of IRAS 12326-6245, and one CHII region candidate, which coincides with
a WISE Class II candidate (see below). Before attempting to identify the
candidate YSOs from the listed WISE and Spitzer sources, we selected the non-
YSO sources with excess infrared emission, such as PAH emitting galaxies,
broad line active galactic nuclei (AGNs), unresolved knots of shock emission,
and PAH emission features. Then, a total of 89 and 51 WISE and Spitzer
sources, respectively, were dropped from the lists. Among the remaining 311
WISE and 759 Spitzer sources, 12 (5 WISE and 7 Spitzer) were identified as
Class I sources (i.e., sources with IR emission arising mainly from a dense
infalling envelope, including flat spectrum objects) and 60 (43 WISE and 17
Spitzer) as Class II sources (i.e., preMS stars with optically thick disks).

Table 4: YSO candidates obtained from MSX, WISE, and Spitzer catalogs.  MSX
sources
---
$\\#$ | Designation | $\alpha$ | $\delta$ | F8 | F12 | F14 | F21 | Notes
|  | [h m s] | [$\circ$ $\prime$ $\prime\prime$] | [Jy] | [Jy] | [Jy] | [Jy] | (class, match with other YSO and
|  |  |  |  |  |  |  | with 13CO/CO or IR structure)
1 | G301.1364-00.2249 | 12 35 35.2 | –63 02 30 | 1.41 | 5.30 | 17.92 | 120.23 | (MYSO) MC3, WISE $\\#$6
2 | G301.1300-00.2354 | 12 35 31.5 | –63 03 06 | 2.216 | 2.702 | 1.477 | 8.930 | (CHII) MC3, WISE $\\#$3
WISE sources
$\\#$ | Designation | $\alpha$ | $\delta$ | W1 | W2 | W3 | W4 |
|  | [h m s] | [$\circ$ $\prime$ $\prime\prime$] | [mag] | [mag] | [mag] | [mag] |
3 | J123531.48-630306.4 | 12 35 31.5 | –63 03 06.4 | 8.759 | 7.615 | 3.207 | -0.535 | (Class I), MC3, MSX $\\#$2
4 | J123529.88-630417.1 | 12 35 29.9 | –63 04 17.1 | 11.847 | 10.812 | 5.958 | 2.858 | (Class I), MC2
5 | J123553.78-630012.6 | 12 35 53.8 | –63 00 12.6 | 9.87 | 8.769 | 5.207 | 1.367 | (Class I), (+)
6 | J123535.13-630230.3 | 12 35 35.1 | –63 02 30.3 | 9.49 | 6.629 | 1.011 | -3.604 | (Class I), MC3, MSX $\\#$1
7 | J123534.52-630147.8 | 12 35 34.5 | –63 01 47.8 | 11.804 | 10.715 | 5.25 | 2.989 | (Class I), MC3, (*)
8 | J123529.71-630425.9 | 12 35 29.7 | –63 04 25.9 | 10.829 | 9.987 | 6.386 | 4.376 | (Class II), MC2
9 | J123528.94-630432.6 | 12 35 28.9 | –63 04 32.6 | 10.492 | 9.912 | 6.113 | 3.56 | (Class II), MC2
10 | J123531.19-630328.3 | 12 35 31.2 | –63 03 28.3 | 10.324 | 9.984 | 6.269 | 3.648 | (Class II), MC3
11 | J123538.44-630318.5 | 12 35 38.5 | –63 03 18.5 | 13.115 | 12.557 | 6.902 | 3.138 | (Class II), MC3
12 | J123537.22-630300.1 | 12 35 37.2 | –63 03 00.1 | 11.906 | 10.999 | 8.069 | 1.563 | (Class II), MC3
13 | J123534.55-630311.1 | 12 35 34.6 | –63 03 11.1 | 11.405 | 10.717 | 5.906 | 3.999 | (Class II), MC3
14 | J123628.80-630500.5 | 12 36 28.8 | –63 05 00.5 | 10.379 | 9.698 | 6.464 | 4.797 | (Class II), Feature 2, Spitzer $\\#$61
15 | J123606.78-630458.5 | 12 36 06.8 | –63 04 58.5 | 10.812 | 10.396 | 9.114 | 6.107 | (Class II)
16 | J123553.95-625717.3 | 12 35 53.9 | \- -62 57 17.3 | 12.834 | 12.427 | 5.898 | 3.722 | (Class II)
17 | J123556.74-625804.3 | 12 35 56.7 | –62 58 04.3 | 13.196 | 12.637 | 6.007 | 4.376 | (Class II)
18 | J123556.58-630058.1 | 12 35 56.6 | –63 00 58.1 | 12.681 | 12.309 | 6.877 | 5.598 | (Class II), MC5, (*)
19 | J123559.25-630043.5 | 12 35 59.3 | –63 00 43.5 | 12.784 | 12.065 | 6.416 | 4.826 | (Class II), MC5, (*)
20 | J123605.26-630000.2 | 12 36 05.3 | –63 00 00.2 | 12.44 | 11.992 | 7.689 | 5.894 | (Class II), MC5, (*)
21 | J123608.20-625952.4 | 12 36 08.2 | –62 59 52.4 | 13.621 | 12.973 | 7.051 | 5.06 | (Class II), MC5, (*)
22 | J123544.39-630151.3 | 12 35 44.4 | –63 01 51.3 | 12.362 | 11.598 | 7.181 | 4.462 | (Class II), MC3, (*)
23 | J123535.61-630129.4 | 12 35 35.6 | –63 01 29.4 | 11.64 | 11.000 | 5.101 | 3.019 | (Class II), MC3, (+)
24 | J123533.44-630130.5 | 12 35 33.4 | –63 01 30.5 | 11.504 | 10.981 | 5.308 | 2.927 | (Class II), MC3, (+)
25 | J123539.82-630127.8 | 12 35 39.8 | –63 01 27.8 | 11.651 | 11.14 | 6.441 | 3.786 | (Class II), MC3, (+)
26 | J123542.44-630148.3 | 12 35 42.3 | –63 01 48.3 | 11.042 | 10.237 | 6.363 | 2.654 | (Class II), MC3, (*), Spitzer $\\#$67
27 | J123538.39-630139.4 | 12 35 38.4 | –63 01 39.4 | 10.882 | 10.354 | 5.752 | 2.963 | (Class II), MC3, (+)
28 | J123525.27-630114.3 | 12 35 25.3 | –63 01 14.3 | 9.593 | 8.952 | 4.275 | 2.262 | (Class II), MC3, (+)
29 | J123528.77-630113.8 | 12 35 28.8 | –63 01 13.8 | 10.472 | 9.535 | 4.4 | 2.781 | (Class II), MC3, (+)
30 | J123522.97-630048.0 | 12 35 23.0 | –63 00 48.0 | 11.545 | 10.94 | 6.364 | 1.844 | (Class II), (+)
31 | J123526.73-630156.8 | 12 35 26.7 | –63 01 56.8 | 11.679 | 10.821 | 5.627 | 2.865 | (Class II), MC3
32 | J123521.24-630158.0 | 12 35 21.2 | –63 01 58.0 | 11.983 | 11.316 | 5.168 | 2.302 | (Class II), MC3, (+)
33 | J123525.41-630204.4 | 12 35 25.4 | –63 02 04.4 | 11.457 | 10.729 | 5.582 | 3.669 | (Class II), MC3
34 | J123524.65-630218.0 | 12 35 24.7 | –63 02 18.0 | 11.548 | 10.785 | 5.839 | 3.296 | (Class II), MC3
35 | J123525.28-630126.8 | 12 35 25.3 | –63 01 26.8 | 10.764 | 10.072 | 5.26 | 4.748 | (Class II), MC3, (+)
36 | J123552.92-625938.6 | 12 35 52.9 | –62 59 38.6 | 13.013 | 12.391 | 6.458 | 4.429 | (Class II), (+)
37 | J123546.50-630010.4 | 12 35 46.5 | –63 00 10.4 | 13.483 | 12.56 | 6.118 | 4.25 | (Class II), (+)
38 | J123551.36-630013.5 | 12 35 51.4 | –63 00 13.5 | 12.246 | 11.293 | 5.803 | 4.145 | (Class II), (+)
39 | J123550.35-630001.0 | 12 35 50.4 | –63 00 01.0 | 12.141 | 11.6 | 6.098 | 4.645 | (Class II), (+)
40 | J123541.99-630100.8 | 12 35 42.0 | –63 01 00.8 | 11.681 | 10.913 | 4.941 | 3.329 | (Class II), (+)
41 | J123543.50-630002.3 | 12 35 43.5 | –63 00 02.3 | 13.039 | 12.656 | 7.897 | 4.866 | (Class II), (+)
42 | J123543.44-630013.6 | 12 35 43.4 | –63 00 13.6 | 12.344 | 12.003 | 6.657 | 4.019 | (Class II), (+)
43 | J123542.11-630037.9 | 12 35 42.1 | –63 00 37.9 | 12.257 | 11.573 | 5.277 | 3.368 | (Class II), (+)
44 | J123543.69-630050.6 | 12 35 43.7 | –63 00 50.6 | 11.771 | 11.174 | 5.177 | 3.43 | (Class II), (+)
45 | J123545.43-630023.2 | 12 35 45.4 | –63 00 23.2 | 11.848 | 11.265 | 5.344 | 3.169 | (Class II), (+)
46 | J123544.72-630038.4 | 12 35 44.7 | –63 00 38.4 | 12.206 | 11.608 | 5.446 | 4.213 | (Class II), (+)
47 | J123543.60-630027.6 | 12 35 43.6 | –63 00 27.6 | 12.647 | 12.25 | 6.496 | 5.909 | (Class II), (+)
48 | J123544.53-625720.1 | 12 35 44.5 | –62 57 20.1 | 7.801 | 6.65 | 5.019 | 4.336 | (Class II), Spitzer $\\#$52
49 | J123454.60-630009.5 | 12 34 54.6 | –63 00 09.5 | 11.662 | 11.15 | 8.66 | 6.146 | (Class II), Feature 3, Spitzer $\\#$72
50 | J123558.47-625603.9 | 12 35 58.5 | –62 56 03.9 | 10.983 | 10.688 | 7.725 | 5.376 | (Class II)
Table 4:  Continued.  Spitzer sources
---
$\\#$ | Designation | $\alpha$ | $\delta$ | 3.6 $\mu$m | 4.5 $\mu$m | 5.8 $\mu$m | 8.0 $\mu$m |
|  | [h m s] | [$\circ$ $\prime$ $\prime\prime$] | [mag] | [mag] | [mag] | [mag] |
51 | G301.1340-00.0894 | 12 35 38.3 | –62 54 21.5 | 12.44 | 11.532 | 10.577 | 9.693 | (Class I)
52 | G301.1489-00.1383 | 12 35 44.5 | –62 57 20.5 | 8.62 | 7.55 | 6.595 | 5.873 | (Class I), WISE $\\#$48
53 | G301.1866-00.1701 | 12 36 03.4 | –62 59 23.0 | 12.424 | 11.565 | 10.773 | 9.521 | (Class I)
54 | G301.1955-00.2012 | 12 36 07.1 | –63 01 16.6 | 11.595 | 10.529 | 9.709 | 8.791 | (Class I), MC5, (*)
55 | G301.1766-00.2110 | 12 35 56.8 | –63 01 47.8 | 14.545 | 12.909 | 11.945 | 12.002 | (Class I), MC4
56 | G301.2000-00.1534 | 12 36 11.0 | –62 58 25.8 | 13.885 | 12.788 | 11.921 | 10.916 | (Class I)
57 | G301.1338-00.2711 | 12 35 32.2 | –63 05 14.4 | 12.721 | 11.54 | 10.612 | 9.959 | (Class I), MC1
58 | G301.1402-00.2298 | 12 35 37.0 | –63 02 47.4 | 12.521 | 11.843 | 11.118 | 10.65 | (Class II), MC3
59 | G301.1395-00.2303 | 12 35 36.6 | –63 02 49.0 | 11.273 | 10.604 | 10.012 | 9.482 | (Class II), MC3
60 | G301.1523-00.2264 | 12 35 43.5 | –63 02 37.8 | 12.498 | 12.146 | 11.893 | 10.887 | (Class II), MC3
61 | G301.2400-00.2609 | 12 36 28.8 | –63 05 00.6 | 10.204 | 9.729 | 9.231 | 7.869 | (Class II), Feature 2, WISE $\\#$14
62 | G301.1973-00.2637 | 12 36 06.1 | –63 05 01.6 | 11.215 | 10.808 | 10.571 | 10.131 | (Class II)
63 | G301.2169-00.2211 | 12 36 17.8 | –63 02 32.7 | 13.615 | 12.915 | 12.217 | 11.599 | (Class II), MC6
64 | G301.1745-00.2102 | 12 35 55.7 | –63 01 44.5 | 13.464 | 12.670 | 12.148 | 10.998 | (Class II)
65 | G301.2153-00.1838 | 12 36 18.1 | –63 00 18.3 | 7.683 | 7.482 | 7.007 | 6.772 | (Class II), MC5, (*)
66 | G301.1311-00.2271 | 12 35 32.3 | –63 02 35.7 | 12.992 | 11.898 | 11.354 | 10.504 | (Class II), MC3
67 | G301.1500-00.2124 | 12 35 42.7 | –63 01 47.0 | 11.332 | 10.95 | 10.353 | 9.342 | (Class II), MC3, (*), WISE $\\#$26
68 | G301.1402-00.2155 | 12 35 37.5 | –63 01 56.0 | 11.273 | 10.637 | 10.065 | 9.522 | (Class II), MC3, (*)
69 | G301.1432-00.2154 | 12 35 39.9 | –63 01 56.3 | 12.104 | 11.591 | 11.068 | 10.594 | (Class II), MC3, (*)
70 | G301.1090-00.1906 | 12 35 21.8 | –63 00 19.6 | 13.625 | 12.923 | 12.308 | 11.642 | (Class II)
71 | G301.1250-00.2201 | 12 35 29.3 | –63 02 09.2 | 10.166 | 9.538 | 8.947 | 8.425 | (Class II), MC3
72 | G301.0569-00.1911 | 12 34 54.3 | –63 00 09.5 | 12.334 | 11.901 | 11.645 | 10.713 | (Class II), Feature 3
73 | G301.0578-00.1910 | 12 34 54.8 | –63 00 09.4 | 12.095 | 11.671 | 11.336 | 10.556 | (Class II), Feature 3, WISE $\\#$49
74 | G301.1525-00.0712 | 12 35 48.6 | –62 53 20.2 | 12.940 | 12.224 | 11.714 | 11.15 | (Class II)

The location of the identified candidate YSOs is shown in Fig. 10, where the
8.0 and 5.8 $\mu$m emissions together with the emission of both the 13CO and
CO structures (i.e, component A and features 1, 2, and 3) are also displayed.
As expected, the bulk of the candidate YSOs are seen projected onto the
densest molecular gas (component A). A noticeable feature is the concentration
of sources along the two brightest PDR areas which, as mentioned in Sect. 3.2,
delineates the borders of the molecular condensations. Namely, the WISE
sources $\\#$5, $\\#$23, $\\#$24, $\\#$25, $\\#$27, $\\#$28, $\\#$29, $\\#$30,
$\\#$32, and $\\#$35 to $\\#$47, seem to be related to the PDR closer to the
ionized gas bordering the molecular emission of MC3 and MC4 (noted as (+) in
the last column of Table 4), while WISE sources $\\#$7, $\\#$18, $\\#$19,
$\\#$20, $\\#$21, $\\#$22, and $\\#$26, and Spitzer sources $\\#$54, $\\#$65,
$\\#$67, $\\#$68, and $\\#$69 lie along the PDR that lies farther away from
the ionized gas and related to the densest parts of MC3, MC4, and MC5 (noted
as (*) in Table 4). Not surprisingly, another appreciable concentration of
candidate YSOs, namely WISE sources $\\#$3, $\\#$6, $\\#$10, $\\#$11, $\\#$12,
and $\\#$13, Spitzer sources $\\#$58, $\\#$59, $\\#$60, $\\#$66, and $\\#$71,
and MSX sources $\\#$1 and $\\#$2 can be noticed in the region of the
molecular concentration MC3.

##  6 Discussion

###  6.1 A model for S169 and its molecular environment

Since IR bubbles are believed to be born within dense molecular clouds, the
classical scenario predicts that, in a uniform medium, the molecular gas
around the bubble should expand spherically. Under that assumption, a
molecular shell with a central velocity V0 and an expansion velocity Vexp
should depict in a position-position diagram a ”disk-ring” pattern when
observed at different velocities. At V0 (corresponding to the systemic
velocity of the bubble), the shell should attain its maximum diameter;
whereas, at extreme velocities (either negative or positive with respect to
V0), the molecular emission should resemble a disk. At intermediate
velocities, the radius of the ring shrinks as the extreme velocity is
approached. Although these features were never observed all together in the
molecular gas around IR bubbles, the behavior of the molecular gas around S169
certainly clashes with that of the classical IR bubble model.

In order to address the observed characteristics of the molecular gas
associated with S169, we present a simple model taking into account its
morphology, relative position with respect to the IR and H$\alpha$ emissions
and the velocity intervals as constraints. Needless to say, we assume that
component A, feature 1, feature 2, and feature 3 are physically associated
with S169 – a well-grounded conclusion given their good morphological
correspondence with the IR nebula.

Figure 11:  Simple sketch of the proposed model to explain the morphology and
velocity of the molecular gas associated with S169. The right ascension axis,
as visualized by the observer, is indicated as reference (the declination axis
is perpendicular to the plane of the sketch, increasing towards the reader).

Hence, we propose that component A and features 1, 2, and 3 could be explained
in terms of an expanding, partially complete semispherical structure. A simple
sketch is presented in Fig. 11, where the right ascension axis is indicated
for the sake of clarity (the declination axis is then perpendicular to the
sketch, increasing towards the reader). The expansion of the molecular gas is
then revealed by the velocity of the different molecular structures. The
molecular gas approaching in the direction of Earth is depicted by the
molecular structure at more negative velocity (component A) with an
approaching velocity of Vappr $\approx$ $-$39 km s-1, while the molecular gas
receding is depicted by the structure at more positive velocity (feature 3)
with a receding velocity of Vrec $\approx$ $-$17 km s-1. Molecular structures
at intermediate velocities (features 1 and 2) thus represent the molecular gas
at V0 $\approx$ $-$25 km s-1 expanding approximately in perpendicular
directions with respect to the direction of Earth. A glance at Fig. 6 suggests
that the expansion velocity of features 1 and 2 might have components in the
direction of the declination axis, which is positive for the case of feature 1
and negative for the case of feature 2 (this cannot be depicted in the two-
dimension sketch presented in Fig. 11).

Therefore, with this simple model we are able to explain not only the
morphology and kinematics of the molecular gas, but also the characteristics
of the ionized gas (traced by the radio continuum emission), which seems to be
expanding against the molecular gas towards the east, north, and south, while
it is expanding freely to the west, probably in a sort of a champagne-flow
effect.

According to the model, a rough estimation of the expansion velocity of the
molecular gas around S169 may be simply obtained by:

$\quad{\rm V}_{\rm exp}=\frac{|{\rm V}_{\rm appr}|-|{\rm V}_{\rm
rec}|}{2\times\ {\rm cos}(\theta)},$ (16)

with $\theta$ as the angle between the line of sight and the direction of Vexp
of component A and feature 2. Based on the morphology of the IR nebula, a
conservative value $\theta$=25∘ can be adopted, and the obtained expansion
velocity turns out to be Vexp $\sim$ 12 km s-1. Having the expansion velocity,
the dynamical age of the nebula, $t_{\rm dyn}$, can be derived considering a
radius of 2.3$\pm$0.8 pc ($\sim$ 4′ at a distance of 2.03${}^{+0.77}_{-0.61}$
kpc), which turns out to be in the range of (1.2 - 2.6) $\times$ 105 yr.

###  6.2 Distance to S169 and IRAS 12326-6245

As noted in Sect. 1, the adopted distance in the literature for IRAS
12326-6245 is 4.4 kpc. This value was first estimated by Osterloh et al.
(1997), who performed a study of a number of southern IRAS sources using IR
continuum and millimeter line data. The authors adopted distances that were
taken from the literature where possible (either directly or by kinematic and
spatial association with another source of known distance). Otherwise,
kinematical distances were determined using the central velocity of the
CS(2-1) or CO(2-1) emission lines and the Galactic rotation curve of Clemens
(1985). For the case of IRAS 12326-6245, a kinematical distance could not be
determined by the authors since the central velocity of the CS and CO lines
($-$39.4 km s-1 and $-$39.7 km s-1, respectively) are forbidden according to
the Galactic rotation model. Then the authors adopted for IRAS 12326-6245 the
distance to the source IRAS 1283-6128 (4.4 kpc), which is kinematically close
to within 0.5 km s-1 as well as spatially to within 2∘.

Figure 12: Left panel. HI emission/absorption profile towards the bright
southern component of the radio continuum arc-shaped feature. Right panel.
Same for the extragalactic source at the north of the nebula. The error
envelopes are plotted with blue dotted lines. Vertical red lines indicate the
velocity of the tangent point.

In Sect. 3.1, we offer a confirmation of the physical association between S169
and IRAS 12326-6245. Furthermore, according to the model proposed in Sect.
6.1, the source IRAS 12326-6245 (embedded in component A) is expanding at a
velocity coarsely estimated as Vexp $\sim$ 12 km s-1 with respect to the
systemic velocity of the bubble (V${}_{0}\approx-$25 km s-1). Assuming that
our model is well-suited to explain the characteristics of the molecular gas
around the bubble nebula, this leads to the unavoidable conclusion that the
velocity reported in the literature for IRAS 12326-6245 ($\sim$ $-$39 km s-1)
does not represent the velocity of the ISM in the surroundings of both, the
IRAS source and S169. This is understandable considering that previous studies
of IRAS 12326-6245 were not performed in the whole context of the nebula S169,
but rather focused only on the IRAS source. Instead, we believe that the
velocity of features 1 and 2 ($\sim$ $-$25 km s-1), corresponding
approximately to the systemic velocity of the bubble nebula, is more suitable
for determining the kinematical distance for both, S169 and IRAS 12326-6245.
Furthermore, in a study of the rotation curve in the Southern Galaxy, Alvarez
et al. (1990) found an excess anomally of $+$12.2 km s-1 in the terminal
velocities in the Galactic longitude range between $l$ = 280 and $l$ = 312,
with respect to the trend in the rest of the IV Galactic quadrant. This excess
includes the region studied here since we have our source at Galactic
coordinates $(l,b)$ = (301.134, –0.225). Next, by adopting a velocity of $-$25
km s-1, and using the rotation curve of Reid et al. (2014) with the Monte
Carlo method141414https://www.treywenger.com/kd/index.php (see Wenger et al.
2018 for details) and a conservative velocity uncertainty of 1 km s-1, two
kinematical distances are obtained, namely: a near kinematical distance of
2.03${}^{+0.77}_{-0.61}$ kpc, and a far kinematical distance of
6.59${}^{+0.62}_{-0.77}$ kpc.

To resolve the twofold ambiguity in kinematic distance, we made use of the Hi
data from the SGPS constructing absorption profiles towards the bright
southern component of the radio continuum arc-shaped feature and, for
comparison, towards the extragalactic source (see Sect. 4). The radio
continuum data were convolved to a beam of 2′ to make them compatible with the
resolution of the Hi data. The expected off-source profile was computed using
a bilinear interpolation method as described in Reynoso et al. (2017). The
results are shown in Fig. 12. From the emission spectra, the tangent point
lies at $\sim-50$ km s-1 (indicated with red vertical lines in Fig. 12), while
emission at the outer Galaxy is observed as far as $\sim+110$ km s-1. The
extragalactic source (right panel) exhibits, as expected, absorption features
up to $\sim-45$ km s-1, coincident with the tangent point within a few km s-1,
a difference that can be ascribed to the ISM turbulence. Another clear
absorption feature at $\sim$ +110 km s-1 confirms the extragalactic nature of
this source. In contrast, the profile corresponding to the arc-shaped feature
(left panel) displays the last significant absorption feature at about
$\sim-30$ km s-1. Considering the absence of absorption in a gap of about 20
km s-1 between this velocity and the tangent point, it is very unlikely that
S169 is located at the far side of the Solar circle. Based on these results,
we conclude that the near kinematical distance (2.03${}^{+0.77}_{-0.61}$ kpc)
is the most adequate for S169. The detection of H$\alpha$ emission at the
center of the nebula (see Fig.1), which is likely the optical counterpart of
the ionized nebula, gives more support to this conclusion since the absorption
that could be produced by the ISM along the far kinematical distance ($\sim$
6.6 kpc) would make the H$\alpha$ emission much fainter or even undetectable.
Furthermore, the whole molecular cloud that contains IRAS 12326-6245 has been
identified to be in the near side of the Carina Spiral Arm by Cohen et al.
(1985).

###  6.3 Possible triggered star formation

The spatial distribution of the candidate YSOs along the borders of the IR
nebula and the densest part of the molecular gas (see Sect. 5.2) is indicative
of a triggered star-forming process such as collect-and-collapse (C$\&$C;
Elmegreen & Lada 1977) or radiative driven implosion (RDI; Lefloch & Lazareff
1994) might be acting on the nebula.

In order to test the C$\&$C mechanism, we use the classical model of Whitworth
et al. (1994) for expanding Hii regions. Then, we estimated the fragmentation
time, $t_{\rm frag}$, and the radius of the Hii region when the fragmentation
occurs, $R_{\rm frag}$, as:

$\quad t_{\rm frag}\ \ =\ 1.56\ {a_{0.2}}^{7/11}\ n_{3}^{-5/11}\
(N_{49})^{-1/11}\ \ 10^{6}\ \ {\rm yr},$ (17) $\quad R_{\rm frag}\ \ =\ 5.8\
{a_{0.2}}^{4/11}\ n_{3}^{-6/11}\ (N_{49})^{1/11}\ \ {\rm pc},$ (18)

with $a_{0.2}$ as the isothermal sound speed in the compressed layer in units
of 0.2 km s-1 ($a_{s}/0.2$ km s-1), $n_{3}$ the surrounding homogeneous
infinite medium into which the Hii region expands, in units of 103 cm-3
($n_{0}/10^{3}$ cm-3), and $N_{49}$ the number of ionizing Lyman continuum
photons, in units of 1049 s-1 ($N^{*}_{\rm Lyc}$/1049 s-1). We adopt $a_{s}$ =
0.2 km s-1 for the collected layer (likely a lower limit since both turbulence
and extra heating from intense subLyman continuum photons leaking from the Hii
region could increase this value) and $N^{*}_{\rm Lyc}$ = 2.5 $\times$ 1047
s-1 (see Sect. 4). To estimate the initial density, $n_{0}$, we averaged the
total mass of the molecular gas around S169 ($\sim$ (7.4 $\pm$ 4.4) $\times$
103 M⊙ ) over a sphere of 2.3 $\pm$ 0.8 pc in radius, which yields a density
in the range of $n_{0}$ = (0.35 - 12) $\times$ 103 cm-3. Thus, we obtain
$t_{\rm frag}$ in the range between 0.5 and 2.2 Myr and $R_{\rm frag}$ between
1.1 and 7.4 pc.

A comparison between $R_{\rm frag}$ and the present radius of the nebula
suggests (within errors) that the C$\&$C process may be acting in S169. On the
contrary, the estimated fragmentation time is higher than the $t_{\rm dyn}$
derived for the region in Sect. 6.1. Thus, there is no conclusive evidence to
support that protostellar objects at the borders of S169 could have been
formed as a result of the C$\&$C mechanism. We ought to keep in mind, however,
that the model of Whitworth et al. (1994) is valid for evolution of the Hii
regions only in a uniform molecular environment, which is certainly not the
case with S169.

Regarding the possibility that the RDI process is taking place in certain
regions of the nebula, a more detailed study is required, with better angular
resolutions of both the molecular and radio continuum data. Nevertheless, we
can clearly conclude that the presence of the intense PDR emission detected as
a bright rim of the densest molecular feature suggests that the RDI process
could be responsible for the formation of some candidate YSOs. In fact, an
alternative scenario for the star formation in S 169 could be that proposed by
Walch et al. (2015), which considers the action of massive stars on a fractal
molecular cloud resulting in a hybrid mechanism that combines elements of C&C
and RDI and which has the capability to generate shell structures with dense
clumps within them. If this were the case for S169, the denser region where
IRAS 12326-6245 is located could be explained by considering an initial non-
uniform molecular cloud with a significant increase in density and, thus, the
RDI process could be the main action responsible for the formation of the
candidate YSOs located behind this border and, in particular, for the IR
source.

##  7 Summary

As part of a broader project aimed at characterizing and studying the physical
properties of Galactic IR bubbles and their surroundings, we present a
multiwavelength analysis of the southern IR dust bubble S169 associated with
the massive star forming core IRAS 12326-6245. To analyze the characteristics
of the molecular gas in the whole nebula, we use CO(2-1), 13CO(2-1), C18O(2-1)
line data, while for the IRAS source we also use HCN(3-2), and HCO+(3-2) data,
all obtained with the APEX telescope. To study the characteristics of the dust
and ionized gas, as well as to investigate the presence of stellar and
protostellar objects in the nebula, we use archival data.

The analysis of the CO(2-1), 13CO(2-1), and C18O(2-1) data allows us to
identify the molecular gas linked to the nebula. We report three molecular
components at $\sim$ $-$39, $-$25, and $-$17 km s-1 (components A, B, and C),
which are morphologically correlated with different regions of the IR nebula.
The component at $-$39 km s-1 (component A) is associated with the brightest
part of the IR nebula (at the southern region) and is the only one detected in
13CO(2-1) and C18O(2-1) emissions. For this component, we have identified six
molecular condensations (namely MC1, MC2, MC3, MC4, MC5, and MC6) based on the
13CO and C18O emissions. These condensations have masses between $\sim$ 90 and
5500 M⊙ , and H2 column densities between $\sim$ 6$\times$1021 and
1$\times$1023 cm-2. The densest one (MC3) is the molecular counterpart of IRAS
12326-6245, which demonstrates the physical association between the nebula and
the IRAS source. An LTE analysis of the HCO+(3-2) and HCN(3-2) lines on this
source, assuming 50 and 150 K, respectively, indicates column densities of
$N$(HCO+) = (5.2 $\pm$ 0.1) $\times$ 1013 cm-2 and $N$(HCN) = (1.9 $\pm$ 0.5)
$\times$ 1014 cm-2. For this condensation, an H2 column density up to $\sim$
8$\times$1023 cm-2 is obtained from the emission at 870 $\mu$m.

The molecular components at $-$25 and $-$17 km s-1 (components B and C) seem
to be associated with the faintest and more external regions of the IR nebula.
Unlike component A, components B and C are not detected in 13CO and C18O
emissions, which indicates that they are composed of low-density gas. This can
be confirmed from Herschel images at 70, 160, and 350 $\mu$m.

With the purpose of explaining the spatial distribution and velocity of
components A, B, and C, we propose a very simple model, which consists of a
partially complete semisphere structure expanding at $\sim$ 12 km s-1.
According to the model, component A represents the molecular gas approaching
at a velocity of $-$39 km s-1, while component C (feature 3) the molecular gas
receding at a velocity of $-$17 km s-1. The systemic velocity of the molecular
gas associated with the bubble is then $-$25 km s-1, which is the velocity of
component B (features 1 and 2). The distribution in the radio continuum
emission at 843 MHz suggests an Hii region bounded by ionization to the east,
north and south, and by density to the west. This appears to be in line with
the model proposed for the molecular gas. The model proposed for the molecular
gas brings an additional discussion about the distance of IRAS 12326-6245 (and
consequently S169) since the radial velocity adopted in the literature for the
IRAS source ($\sim$ –39 km s-1) would not be representative of the velocity of
the ISM in its surroundings. Instead, we believe that the systemic velocity of
the bubble (–25 km s-1) is more adequate for determining its kinematical
distance. Then, using the Galactic rotation model and HI absorption profiles,
we determine a kinematical distance of 2.03${}^{+0.77}_{-0.61}$ kpc for S169
and IRAS 12326-6245.

Using point source catalogs, we identify10 ionizing stars candidates projected
onto the cavity. They have spectral types between O9V and B2, which are
necessary to sustain the current level of ionization of the Hii region. We
also identify a number of candidate YSOs projected mostly onto component A,
more precisely: MC3, MC4, and MC5, confirming that active star formation has
developed along the borders of the bubble. After comparing the fragmentation
time and fragmentation radius with the age and current radius of the nebula,
we cannot assert that the C$\&$C process is acting in the collected layers of
gas at the edge of the bubble. We do keep in mind, however, the limitations of
the models applied to S169 and the possibility that other triggering star-
forming processes, such as RDI or a combination of both, RDI and C$\&$C, could
be acting in the region.

In summary, the infrared bubble S169 is a relatively young HII region that has
profoundly affected its surroundings, creating a molecular shell that
continues to possess a large expanding motion and where several condensations
host candidate protostellar objects in their dense interiors, with the source
IRAS 12326-6245 standing out as the most striking among them.

###### Acknowledgements.

We would like to thank the anonymous referee for his/her helpful comments and
suggestions that led to the improvement of this paper. N.U.D and M.A.C
acknowledge support from UNLP, PPID G005 and CONICET grant PIP
112$-$201701$-$00507 (Argentina). L.B. and R.F. acknowledges support from
CONICYT project Basal AFB-170002. E.M. acknowledges support from the Brazilian
agency CNPq (grant 150465/2019-0). L.A.S acknowledges support from UNLP PPID
G005 grant. S.C., L.A.S. and E.M.R. are partially funded by CONICET grant PIP
112-201701-00604 (Argentina).

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