# SOFIA FEEDBACK survey: PDR diagnostics of stellar feedback in different
regions of RCW 49

M. Tiwari  University of Maryland, Department of Astronomy, College Park, MD
20742-2421, USA  Max-Planck Institute for Radioastronomy, Auf dem Hügel, 53121
Bonn, Germany  M. Wolfire  University of Maryland, Department of Astronomy,
College Park, MD 20742-2421, USA  M. W. Pound  University of Maryland,
Department of Astronomy, College Park, MD 20742-2421, USA  E. Tarantino
University of Maryland, Department of Astronomy, College Park, MD 20742-2421,
USA  R. Karim  University of Maryland, Department of Astronomy, College Park,
MD 20742-2421, USA  L. Bonne  SOFIA Science Center, USRA, NASA Ames Research
Center, M.S. N232-12, Moffett Field, CA 94035, USA  C. Buchbender  I. Physik.
Institut, University of Cologne, Zülpicher Str. 77, 50937 Cologne, Germany  R.
Güsten  Max-Planck Institute for Radioastronomy, Auf dem Hügel, 53121 Bonn,
Germany  C. Guevara  I. Physik. Institut, University of Cologne, Zülpicher
Str. 77, 50937 Cologne, Germany  S. Kabanovic  I. Physik. Institut, University
of Cologne, Zülpicher Str. 77, 50937 Cologne, Germany  Ü. Kavak  SOFIA Science
Center, USRA, NASA Ames Research Center, M.S. N232-12, Moffett Field, CA
94035, USA  M. Mertens  I. Physik. Institut, University of Cologne, Zülpicher
Str. 77, 50937 Cologne, Germany  N. Schneider  I. Physik. Institut, University
of Cologne, Zülpicher Str. 77, 50937 Cologne, Germany  R. Simon  I. Physik.
Institut, University of Cologne, Zülpicher Str. 77, 50937 Cologne, Germany  J.
Stutzki  I. Physik. Institut, University of Cologne, Zülpicher Str. 77, 50937
Cologne, Germany  A. G. G. M. Tielens  University of Maryland, Department of
Astronomy, College Park, MD 20742-2421, USA  Leiden Observatory, Leiden
University, PO Box 9513, 2300 RA Leiden, Netherlands.

(Received March 25, 2022; Accepted August 16, 2022)

###### Abstract

We quantified the effects of stellar feedback in RCW 49 by determining the
physical conditions in different regions using the [C II] 158 $\mu$m and [O I]
63 $\mu$m observations from SOFIA, the 12CO (3-2) observations from APEX and
the H2 line observations from Spitzer telescopes. Large maps of RCW 49 were
observed with the SOFIA and APEX telescopes, while the Spitzer observations
were only available towards three small areas. From our qualitative analysis,
we found that the H2 0-0 S(2) emission line probes denser gas compared to the
H2 0-0 S(1) line. In four regions (“northern cloud”, “pillar”, ”ridge”, and
“shell”), we compared our observations with the updated PDR Toolbox models and
derived the integrated far-ultraviolet flux between 6–13.6 eV ($G_{\rm 0}$), H
nucleus density ($n$), temperatures and pressures. We found the ridge to have
the highest $G_{\rm 0}$ (2.4 $\times$ 103 Habing units), while the northern
cloud has the lowest $G_{\rm 0}$ (5 $\times$ 102 Habing units). This is a
direct consequence of the location of these regions with respect to the Wd2
cluster. The ridge also has a high density (6.4 $\times$ 103 cm-3), which is
consistent with its ongoing star formation. Among the Spitzer positions, we
found the one closest to the Wd2 cluster to be the densest, suggesting an
early phase of star formation. Furthermore, the Spitzer position that overlaps
with the shell was found to have the highest $G_{\rm 0}$ and we expect this to
be a result of its proximity to an O9V star.

ISM — Photodissociation regions — Molecular clouds

##  1 Introduction

Stellar feedback is one of the most important ingredients in the evolution of
the interstellar medium (ISM). In particular, massive stars play a key role in
regulating different feedback mechanisms (for recent observational studies see
Pabst et al. 2019, Luisi et al. 2021, Tiwari et al. 2021 and Kavak et al.
2022). They inject immense amounts of energy into their surroundings through
stellar winds and through extreme ultra-violet (EUV, $h\nu$ $>$ 13.6 eV) and
far-UV (FUV, 6 $<$ $h\nu$ $<$ 13.6 eV) photons. The EUV photons ionise gas (H
$\to$ H+) in the vicinity of a star giving rise to H II regions. The FUV
photons heat the gas via the photoelectric effect on small grains and PAHs,
dissociate molecules, and ionise atoms e.g., C $\to$ C+, giving rise to photo-
dissociation regions (PDRs). This mechanical and radiative energy input
catalyses various physical and chemical processes, which shape the dense and
diffuse gas into molecular clouds and into structures like shells and pillars.
These structures vary in their physical conditions based on their location
with respect to the illuminating source and on the geometry of the surrounding
medium.

Determination of physical conditions in these molecular clouds, shells and
pillars is done through various chemical species. Neutral atomic gas in the
ISM cools mainly at far-infrared wavelengths through [C II] (158 $\mu$m) and
[O I] (63 and 145 $\mu$m) line emission and probes regions predominantly at
$A_{\rm v}$ $<$ 3\. Molecular gas cooling is dominated by 12CO rotational
emission which probes regions at larger $A_{\rm v}$ (for details see
Hollenbach & Tielens 1999; Wolfire et al. 2003, 2022).

RCW 49 is one of the most luminous massive star-forming regions in our Galaxy.
It is located at a distance of 4.16 $\pm$ 0.33 kpc (Vargas Álvarez et al.
2013, Zeidler et al. 2015 and also see discussion in Tiwari et al. 2021),
close to the tangent of the Carina arm at $l$ = 284$\degree$.3, $b$ = -
0$\degree$.3. RCW 49 hosts a compact stellar cluster, Westerlund 2 (Wd2),
consisting of 37 OB stars and 30 early-type OB star candidates surrounding it
(Ascenso et al. 2007; Tsujimoto et al. 2007; Rauw et al. 2011; Mohr-Smith et
al. 2015; Zeidler et al. 2015). There are two Wolf-Rayet stars also associated
with RCW 49: WR20a, which is a part of the Wd2 cluster and is suggested to be
among the most massive binaries in the Galaxy (Rauw et al., 2005), and WR20b,
which is a few arcminutes away south-east of Wd2. Winds and radiation from
these stars are responsible for sculpting the gas and dust into dense
molecular clouds, shells and pillars in RCW 49.

RCW 49 has been studied previously at radio, sub-millimeter, IR and X-ray
wavelengths. Whiteoak & Uchida (1997) and Benaglia et al. (2013) studied radio
data (at 0.843, 1.38, 2.38, 5.5, 9 GHz) and reported two shells towards RCW
49. Furukawa et al. (2009) obtained 12CO (2-1) data with NANTEN2 telescope
towards RCW 49. With an angular resolution of 1.5$\arcmin$, they identified
two large-scale molecular clouds in velocity ranges of -11 to 9 km s-1 and 11
to 21 km s-1, and suggested that their collision led to the formation of the
Wd2 cluster. Ohama et al. (2010) analysed 12CO (1-0), 12CO (2-1), and 13CO
(2-1) to estimate temperature and density distributions for the two clouds
identified by Furukawa et al. (2009). Churchwell et al. (2004) studied RCW 49
at mid-IR wavelengths (Infrared Array Camera data at 3.6, 4.5, 5.8, and 8.0
$\mu$m) and identified different regions based on dust emission with respect
to the angular radius from Wd2. Moreover, diffuse X-ray observations (0.5–7
keV) probed the hot plasma distributed around Wd2 (Townsley et al., 2019).

Although the studies mentioned above unveiled the rich region of RCW 49, the
observations used for the analysis lacked either spatial or spectral
resolution, which hindered precise investigation of the morphology, energetics
and physical conditions. The new high resolution [C II] 158 $\mu$m and the [O
I] 63 $\mu$m observations towards RCW 49 were taken as part of the
Stratospheric Observatory For Infrared Astronomy (SOFIA, Young et al. 2012)
legacy program FEEDBACK111https://feedback.astro.umd.edu/ (Schneider et al.,
2020) and the 12CO observations were taken with the Atacama Pathfinder
Experiment (APEX222APEX is a collaboration between the Max-Planck-Institut für
Radio-astronomie, Onsala Space Observatory, and the European Southern
Observatory., Güsten et al. 2006). In our previous work, we characterised the
expanding shell of RCW 49 using the [C II] , 12CO and 13CO observations
(Tiwari et al., 2021). With the goal of quantifying radiative and mechanical
input by massive stars into their surroundings, similar studies were performed
to understand the stellar feedback mechanisms in the ISM through the FEEDBACK
program (first scientific results in Luisi et al. 2021; Tiwari et al. 2021;
Kabanovic et al. 2021; Beuther et al. 2022).

In this paper, we quantify the effects of stellar feedback on different
regions of RCW 49 by determining their physical conditions. We want to
understand the contribution of stellar feedback in the evolution of these
regions in terms of morphology and future star formation. We compare our
observations (obtained through the FEEDBACK program and through the Spitzer
telescope) with PDR models to derive the FUV incident flux ($G_{\rm 0}$),
density (n), surface temperature ($T_{\rm surf}$) and pressures ($p$) in
different regions of RCW 49. Moreover, through this paper, we aim at providing
a suitable PDR analysis strategy that can be used to determine the physical
conditions in the ISM.

##  2 Observations

###  2.1 SOFIA

The [C II] and [O I] observations were observed during three flights in June
of 2019 using upGREAT333German Receiver for Astronomy at Terahertz. (up)GREAT
is a development by the MPI für Radioastronomie and the KOSMA/Universität zu
Köln, in cooperation with the DLR Institut für Optische Sensorsysteme.
(Risacher et al., 2018). The upGREAT receiver can observe both [C II] and [O
I] lines simultaneously by using a dual 7 pixel low-frequency array (LFA) that
was tuned to the [C II] line, and in parallel a seven pixel high frequency
array (HFA) that was tuned to the [O I] 63 $\mu$m line. The observing mode was
driven by the [C II] line sensitivity, thus, limiting the S/N of the [O I]
line. Moreover, the [O I] data was under sampled (for more details see
Schneider et al. 2020).

A Fast Fourier Transform Spectrometer (FFTS) with 4 GHz instantaneous
bandwidth and a frequency resolution of 0.244 MHz (Klein et al., 2012) served
as backend. Thus, both the [C II] and [O I] data have an original resolution
of 0.04 km s-1. The data set was then re-binned to a spectral resolution of
0.2 km s-1. The instrument and telescope optics determine the intrinsic half-
power beam widths: 14.1$\arcsec$ for the [C II] line and 6.3$\arcsec$ for the
[O I] line. All spectra are presented on a main beam brightness temperature
($T_{\rm mb}$) scale. The main beam efficiency ($\eta_{\rm mb}$) values for [C
II] and [O I] are 0.65 and 0.69, respectively. Further observational details
are given Tiwari et al. (2021).

Figure 1: A red blue green (RBG) image of RCW 49 is shown in the upper left.
Red corresponds to the [C II] integrated intensity emission within the
velocity range of -30 to 30 km s-1. Green corresponds to the GLIMPSE (Benjamin
et al., 2003) 8 $\mu$m emission and blue corresponds to the ATLASGAL (Schuller
et al., 2009) 870 $\mu$m continuum. Different interesting regions towards
which the PDR analysis is done are shown and outlined in white boxes. The [C
II] and 12CO emission towards the northern cloud, ridge, shell and pillar are
shown in the middle and right panels. The red crosses mark the lines of sight
for which the spectra are presented in Fig. 3. In the lower left panel, the 8
$\mu$m emission image is overlaid with [C II] emission contours highlighting
these regions. The contour colors are same to the ones in the [C II] maps
shown in the middle panels. The Wd2 cluster, the WR20b star and the O9V star
are shown marked in pink in all panels.

###  2.2 APEX

The $J$ = 3-2 transition of 12CO and 13CO was observed towards RCW 49 using
the LAsMA array on the APEX telescope (Güsten et al., 2006). The beam width is
18.2$\arcsec$ at 345.8 GHz. Advanced FFTS (Klein et al. 2012) are used as
backends with a bandwidth of 2 $\times$ 4 GHz and a native spectral resolution
of 61 kHz. All spectra are calibrated in $T_{\rm mb}$ with $\eta_{\rm mb}$ =
0.68 at 345.8 GHz. After a linear baseline subtraction, all data were binned
to 0.2 km s-1. The angular resolution of the maps is $\sim$ 20$\arcsec$. All
other data sets were convolved to this resolution for the PDR analysis
described later in this paper. For more observational details, see (Tiwari et
al., 2021).

###  2.3 Spitzer

The H2 observations were carried out using the Infrared Spectrograph (IRS,
Houck et al. 2004) on board the Spitzer Space Telescope (Werner et al., 2004).
The H2 0-0 S(1) and H2 0-0 S(2) line data presented in this paper are part of
the Spitzer program PID 20012 (PI M. Wolfire). The program observed seven
fields (Castellanos et al., 2014) and we analysed three (p1, p3 and p7) of
them as shown in Fig. 1. The AOR keys of each of them are 13812992, 13813504,
and 13814528. The spectrograph observes in four modes depending on the
wavelengths and resolution (Houck et al., 2004): short wavelengths with low
resolution (SL), long wavelengths with low resolution (LL), short wavelengths
with high resolution (SH), and long wavelengths with high resolution (LH). The
S(1) and S(2) lines are at 17.03 and 12.28 $\mu$m observed with the SH mode.
The pixel size is 2.26$\arcsec$, the slit size is 4.7$\arcsec$ $\times$
11.3$\arcsec$, and the mapped regions are $\sim$ $50\arcsec\times 60\arcsec$
for p1, $\sim$ $50\arcsec\times 50\arcsec$ for p3 and $\sim$ $40\arcsec\times
40\arcsec$ for p7.

Observation data cubes (along with uncertainty cubes) were generated with the
CUbe Builder for IRS Spectra Maps
(CUBISM444https://irsa.ipac.caltech.edu/data/SPITZER/docs/
dataanalysistools/tools/cubism/, Smith et al. 2007a), which is a tool used for
constructing data cubes of the mapping mode spectra taken with the Spitzer IRS
spectrograph. We then used
PAHFIT555http://tir.astro.utoledo.edu/jdsmith/research/pahfit.php (Smith et
al., 2007b) on the data cubes, which decomposes the IRS spectra in order to
give us intensity values at each pixel such that we have H2 intensity and
uncertainty maps.

##  3 Source selection in RCW 49

Figure 1 shows RCW 49 marked with seven different regions (in white boxes),
which are spatially and in some cases are spectrally distinct. These regions
are at different distances to the Wd2 cluster and some have very distinct
morphology, making them ideal to study the impact of stellar feedback in RCW
49. In this section we introduced these different regions by describing their
physical and chemical characteristics.

We have selected the “northern cloud” as representative of the cloud whose
collision with another cloud may have triggered the formation of the Wd2
cluster (Furukawa et al., 2009). This cloud can be spatially disentangled
through its velocity range (2 to 8 km s-1). The “ridge” position was selected
as part of the other cloud that is thought to have been involved in the cloud
collision that gave rise to the formation of Wd2 (Furukawa et al., 2009). Part
of the ridge is associated with the inner dust ring identified in the Spitzer
study of RCW 49 (Churchwell et al., 2004). The velocity of the ridge is in the
range 16 to 22 km s-1. We have selected a region in the “shell” that was
identified by Tiwari et al. (2021) as expanding toward us at 12 km s-1. This
particular position is well separated from the ridge and pillar positions
included in our analysis. The shell spans the velocity range -25 to 0 km s-1
and is most intense in -12 to 0 km s-1 range. We have selected the “pillar” as
it may mark the radiative interaction of Wd2 with the ridge. However, while
the pillar points back toward Wd2, we note that the Wolf-Rayet star, WR20b is
in close proximity and may have played a key role in its formation and
determining its physical conditions. While the above regions where selected
based on morphological considerations of RCW 49, we have added three positions
selected from the Spitzer IRS spectral survey also covered by the [C II] , [O
I] and 12CO (3-2) surveys. Position “p1” (Fig. 1, upper left panel), marks the
transition from the cavity surrounding Wd2 to the bright [C II] and 8 $\mu$m
emission. Position “p3” overlaps with the eastern part of the ridge. Position
“p7” is at the northern tip of the expanding shell identified by Tiwari et al.
(2021).

##  4 Spatial distribution of observed species

The northern cloud has bright [C II] and 12CO emission but appears to be less
prominent in GLIMPSE 8 $\mu$m emission (Fig. 1). This is likely due to the
high spectral resolution of the [C II] and 12CO (3–2) observations, which
allows us to disentangle the emission from this component (from 2 to 8 km s-1)
from the surroundings. Figure 1 shows bright [C II] , 12CO and 8 $\mu$m
emission towards the ridge. The shell is outlined very well by the [C II] and
12CO emission as seen in Fig. 1. However, unlike the [C II] emission
distribution, the 12CO emission distribution is fragmented towards the shell.
This is also seen in the 13CO emission map towards this region (Tiwari et al.,
2021). For the 8 $\mu$m emission, similar to the northern cloud, one cannot
disentangle the emission solely from the shell due to its lack of velocity
resolution. Bright [C II] , 12CO and 8 $\mu$m emission is observed towards the
pillar.

The [C II] and 12CO (3-2) emission maps overlaid with H2 0-0 S(1) and H2 0-0
S(2) emission contours respectively, towards these positions are shown in Fig.
2. Due to the lower velocity resolution $\Delta v\sim 500$ km ${\rm s^{-1}}$
in the IR line data, we compare their spatial distribution with the [C II] and
12CO maps where their emission is integrated over the entire velocity range.
This ensures a fair comparison between the IR, FIR and submm wavelength data
presented here. In general, the H2 0-0 S(1) emission distribution seems to be
related more to the [C II] emission distribution than to the 12CO emission
distribution, and the H2 0-0 S(2) emission distribution seems to be related
more to the 12CO emission distribution than to the [C II] emission
distribution. This suggests that the H2 0-0 S(1) line traces less dense gas
compared to the H2 0-0 S(2) line. For p1, there are differences in the spatial
distribution of the emission between H2 0-0 S(1) and H2 0-0 S(2). In p3, the
H2, [C II] and 12CO emission is brighter in the south-west, which is in the
direction of the Wd2 cluster. In p7, the H2 emission is essentially tracing
the northern-most part of the shell. This structure is less detailed in the [C
II] and 12CO maps because of the lower angular resolution compared to the IR
data. The dependence of the H2 lines’ emission on the physical conditions can
be seen in PDR model contour plots666e.g.,
https://dustem.astro.umd.edu/models/wk2006/h200s2s1
z1web.html. Along with the H2 lines, several ionised gas tracers ([S IV] 10.5
$\mu$m, [Ne II] 12.8 $\mu$m, [Ne III] 15.6 $\mu$m, [S III] 18.7 $\mu$m) were
also mapped using the Spitzer telescope. Their corresponding IR spectra
averaged over the regions of p1, p3 and p7 are discussed and shown in Appendix
A and Fig. 10, respectively.

Figure 2: Emission maps of [C II] and 12CO (3–2) towards the regions p1, p3
and p7 (shown in Fig. 1). The [C II] maps are overlaid with H2 0-0 S(1)
intensity contours and the 12CO maps are overlaid with H2 0-0 S(2) intensity
contours. The contours are colored such that the darkest shade of red
corresponds to the highest intensity. The red crosses mark the lines of sight
for which the spectra are presented in Fig. 3.  Figure 3: Representative
spectra of [C II] , 12CO, 13CO and [O I] towards the northern cloud (nc),
ridge, shell, pillar, p1, p3 and p7. All data was convolved to the same
angular resolution of $\sim$ 20$\arcsec$ and the spectral resolution of the
spectra shown here is 1 km s-1.

##  5 Spectra

Figure 3 shows the representative spectra of the [C II] , 12CO, 13CO and [O I]
63 $\mu$m lines towards the seven regions analysed in detail in this work.

We have convolved the maps to the same spatial resolution and extracted data
for a single pixel. The similarity in peak position and profile of all tracers
for positions, the ridge, p1 and the two velocity components in p3, ensures
that we are tracing the same gas. However for the northern cloud, shell,
pillar and p7, the velocity peak positions of the [C II] and [O I] lines are
shifted by $\sim$ 1 km s-1 with respect to the 12CO line. These velocity
shifts are in accordance with expectations for the advection of gas through
the PDR (Tielens, 2005) and we surmise that these species trace gas in the
same PDR. In our source selection, we avoided clumps, which is vital for the
PDR analysis described in Sect. 6.1. Though not used in our PDR analaysis, we
also present spectra of 13CO towards the different regions. We can see that
the 13CO spectral profile follows that of 12CO, suggesting no self absorption
in 12CO from any cold foreground. Similarly, we did not observe any self
absorption features in the [O I] line profiles along the lines of sight where
it was detected. This is in contrast with most [O I] observations towards
other Galactic sources (Poglitsch et al. 1996; Liseau et al. 2006; Gerin et
al. 2015; Rosenberg et al. 2015; Schneider et al. 2018; Goldsmith et al.
2021).

We considered the intensities integrated over the specific velocity ranges
corresponding to their regions (same as shown in Fig. 1). For the northern
cloud and p1, the [O I] intensity values are upper limits to the [O I]
emission towards these regions. For [C II] , 12CO and [O I] , they are
estimated for the entire mapped regions by finding the rms noise in an
emission free velocity window. Their values are given in Table 1. Furthermore,
there are also calibration uncertainties, which could be up to 10% for 12CO
data (APEX, LASMA), $<$ 20% for the [C II] data (SOFIA, upGREAT) and $\geq$
20% for the [O I] data.

Figure 4: Overlay plots of the observed intensity ratios in the PDR Toolbox
modeled grid of $G_{\rm 0}$ and $n$ for the northern cloud, ridge, shell and
pillar. The shaded region corresponds to the error on the observations. Each
plot is overlaid with the contours of its corresponding corner plot.

Figure 5: Overlay plots of the observed intensity ratios in the PDR Toolbox
modeled grid of $G_{\rm 0}$ and $n$ for p1, p3 (upper and lower panels include
and exclude H2 data respectively.) and p7 in the upper panel. The shaded
region corresponds to the error on the observations. Each plot is overlaid
with the contours of its corresponding corner plot .

##  6 PDR diagnostics

We have used the PDR Toolbox777http://dustem.astro.umd.edu/ (Kaufman et al.
2006; Pound & Wolfire 2008, 2011), a Python-based software package which
employs state-of-the-art PDR models to determine the physical conditions in
PDRs (the far UV flux ($G_{\rm 0}$) and H nucleus density ($n$)) from
observations.

Table 1: Observed and model intensity∗ values in different regions of RCW 49.

Parameter | rms noise | northern cloud | ridge | shell | pillar | p1 | p3 | p3v1 | p3v2 | p7
---|---|---|---|---|---|---|---|---|---|---
R.A. (J2000) |  | 10:23:59.7 | 10:24:06.5 | 10:24:29.7 | 10:24:11.7 | 10:23:56.2 | 10:24:30.3 |  |  | 10:24:11.1
Dec. (J2000) |  | -57:38:52.7 | -57:46:17.2 | -57:48:12.4 | -57:49:06 | -57:43:22.7 | -57:45:02.5 |  |  | -57:40:32.3
$I$([C II] ) | 0.18 | 6.7 | 9.9 | 9.4 | 11.0 | 3.0 | 34.0 | 20.0 | 12.0 | 4.2
$I$(12CO (3-2)) | 0.003 | 0.032 | 0.029 | 0.017 | 0.017 | 0.028 | 0.076 | 0.047 | 0.029 | 0.014
$I$([O I] ) | 3.3 | 3.7 | 30.0 | 20.0 | 18.0 | 3.3 | 110.0 | 13.0 | 15.0 | 14.0
$I$(H2 0-0 S(1)) | 0.05 |  |  |  |  | 0.31 | 1.6 |  |  | 1.0
$I$(H2 0-0 S(2)) | 0.06 |  |  |  |  | 0.75 | 2.2 |  |  | 1.2
$I$(FIR) |  | 2400 | 12200 | 8200 | 5700 | 4200 | 12400 |  |  | 3800
$I_{\rm pre}$([C II] ) |  | 4.0 | 6.2 | 5.5 | 4.6 | 6.6 | 5.9 | 4.7 | 5.6 | 10.0
$I_{\rm pre}$(12CO (3-2)) |  | 0.023 | 0.021 | 0.011 | 0.0082 | 0.039 | 0.014 | 0.013 | 0.014 | 0.031
$I_{\rm pre}$([O I] ) |  | 11.0 | 24 | 14 | 8.7 | 67.4 | 17 | 11 | 17 | 110
$I_{\rm pre}$(H2 0-0 S(1)) |  |  |  |  |  | 0.79 | 0.34 |  |  | 0.220
$I_{\rm pre}$(H2 0-0 S(2)) |  |  |  |  |  | 0.16 | 0.022 |  |  | 0.75

Note. — * All intensities are in the units of 10-4 erg cm-2 s-1 sr-1. The
observed intensities were converted from K km s-1 ($W$) to erg cm-2 s-1 sr-1
($I$). For [C II] 158 $\mu$m, 12CO (3–2) and [O I] 63 $\mu$m, respectively,
the conversions are $I$ (erg s-1 cm-2 sr-1) = 7.05 $\times$ 10-6 $W$ (K km
s-1), 4.238 $\times$ 10-8 $W$ (K km s-1) and 1.1 $\times$ 10-4 $W$ (K km s-1).
$I$ and $I_{\rm pre}$ are the observed intensities and intensities predicted
by models, respectively. The columns p3, p3v1 and p3v2 correspond to the
intensities integrated over the velocity window of -30 to 30 km s-1, -15 to 5
km s-1 and 5 to 20 km s-1, respectively.

In fitting the observations, we used the “Wolfire-Kaufman 2020 (wk2020)” model
set. These are plane-parallel PDR models, with radiation that is incident on
one side normal to the layer, that solve for the gas temperature in thermal
equilibrium, and atomic and molecular abundances in steady state. This model
set has updated chemistry from that used in the Kaufman et al. (2006) models.
Most notably, the chemical rates, PAH chemistry, and collision rates discussed
in Hollenbach et al. (2012), Neufeld & Wolfire (2016), Kovalenko et al.
(2018), Tran et al. (2018), and Dagdigian (2019), plus photodissociation and
photoionization rates in Heays et al. (2017), and collisional excitation rates
for [O I] in Lique et al. (2018). The models are carried out to a depth
$A_{\rm v}=7$ with grain-surface chemistry turned off. We also adopt an
incident (primary) cosmic-ray ionization rate per hydrogen $=2.0\times
10^{-16}$ ${\rm s^{-1}}$ that decreases with depth as $1/N$ as suggested by
Neufeld & Wolfire (2017). The PDR Toolbox uses the face-on line intensities
that emerge from the PDR.

We used the [C II] , [O I] and 12CO (3-2) observations in the northern cloud,
ridge, shell and pillar. In addition to the above species, we used the H2 0-0
S(1) and H2 0-0 S(2) lines in p1, p3 and p7. The models take into account
optical depth effects in the PDR. Moreover, we do not see any signature of
self-absorption in our spectra (Fig. 3) from a cold foreground. We convolved
all the observations to the same angular resolution of 12CO, which is $\sim$
20$\arcsec$. The observed line intensities are summarized in Table 1. As
mentioned before, these values are toward specific lines of sight (see Fig. 1
and 2).

###  6.1 Overlay plots

Plotting overlays of intensity ratios of the observed species in the model
space is a good way to understand the phase space (Figs. 4 and 5) of allowable
solutions. In such plots, locations where many observational lines cross
indicate areas of ($n,G_{0}$) space that provide a good fit to the
observations. For p1 and p7, we used the [C II] , [O I] , and 12CO integrated
intensities of the -15 to -5 km s-1 and the -12 to 0 km s-1 components,
respectively, as these dominate the spectra (Fig. 3) and the H2 data. However,
for p3, both -15 to 5 and 5 to 20 km s-1 velocity components contribute
significantly to the emission. Thus, we carried out the PDR analysis for p3
for three cases: first, where we used the integrated intensities of [C II] ,
[O I] and 12CO in the entire velocity spread of -30 to 30 km s-1 together with
the H2 emission; second, where we only used the integrated intensities of [C
II] , [O I] and 12CO within the -15 to 5 km s-1 velocity range and third,
where again, we only used the integrated intensities of [C II] , [O I] and
12CO within the 5 to 20 km s-1 velocity range. Moreover, we included ratios of
the observed line intensities to the FIR continuum intensity. The FIR
intensities were determined using the formalism described in Goicoechea et al.
(2015), where the dust parameters (dust temperature and opacity) were
calculated from 70 and 160 $\mu$m data from the Herschel Space Archive (HSA),
obtained within the Hi-GAL Galactic plane survey (Molinari et al., 2010)
observed with the Photodetector Array Camera and Spectrometer (PACS; Poglitsch
et al. 2010). A detailed description of the method is given in (Tiwari et al.,
2021).

Table 2: Derived physical parameters in different regions of RCW 49.

Parameter | units | northern cloud | ridge | shell | pillar | p1 | p3 | p3v1 | p3v2 | p7
---|---|---|---|---|---|---|---|---|---|---
$n$ | 103 cm-3 | 8.6(0.5) | 6.4(0.35) | 3.1(0.2) | 2.2(0.1) | 27.0(2.0) | 4.3(0.1) | 3.6(0.1) | 4.5(0.2) | 14.0(1.0)
$G_{\rm 0}$ | 103 Habing units | 0.5(0.07) | 2.4(0.3) | 1.9(0.4) | 0.97(0.21) | 2.9(0.32) | 2.1(0.2) | 0.79(0.07) | 2.0(0.2) | 32.0(8.0)
$G_{\rm 0}$(FIR) | 103 Habing units | 0.92 | 4.7 | 3.2 | 2.2 | 1.7 | 4.8 |  |  | 1.5
$T_{\rm surf}$ | K | 160 | 246 | 250 | 223 | 244 | 242 | 196 | 242 | 441
$T_{\rm ex}$([C II] ) | K | 55.2 | 72.0 | 47.3 | 54.6 | 33.6 |  | 70.4 | 73.0 | 36.8
$p_{\rm th}$ |  $10^{6}$ K cm-3 | 1.4 | 1.6 | 0.77 | 0.5 | 6.6 | 1.0 | 0.71 | 1.1 | 6.2
$p_{\rm turb}$ |  $10^{6}$ K cm-3 | 9.4 | 3.0 | 3.5 | 1.8 | 25.0 | 12.0 | 3.9 | 1.8 | 19.0
$\Delta v^{2}_{\rm turb}$ | km2 s-2 | 38.5 | 16.7 | 39.9 | 29.3 | 32.1 | 101.4 | 38.1 | 14.2 | 48.4
$N_{\rm H}$ (160 $\mu$m) | 1022 cm-2 | 4.5 | 3.3 | 5.6 | 4.2 | 4.5 |  | 13.0 | 13.0 | 4.6
$N_{\rm H}$ (13CO) | 1022 cm-2 | 1.3 | 0.84 | 0.36 | 0.78 | 1.3 |  | 1.6 | 1.3 | 0.4

Note. — $n$ is hydrogen volume density and $G_{\rm 0}$ is far-UV flux
determined from the PDR ToolBox. Errors from fits are in parenthesis. $G_{\rm
0}$ (FIR) is the far-UV flux determined from the Herschel data. $T_{\rm
ex}$([C II] ) is the [C II] excitation temperature. $p_{\rm th}$ and $p_{\rm
turb}$ are the thermal and turbulent pressures. $\Delta v^{2}_{\rm turb}$ is
the turbulent velocity in km2 s-2. $N_{\rm H}$ (160 $\mu$m) and $N_{\rm H}$
(13CO) are hydrogen column densities estimated using the 160 $\mu$m Herschel
data and the 13CO APEX data. The columns p3, p3v1 and p3v2 correspond to the
intensities integrated over the velocity window of -30 to 30 km s-1, -15 to 5
km s-1 and 5 to 20 km s-1, respectively.

We used two complementary fitting techniques from the PDR Toolbox to estimate
the radiation field and H density, $G_{\rm 0}$ and $n$. The first is the
Levenberg-Marquardt least squares fit (LSQ) method, which finds “best-fit”
$G_{\rm 0}$ and $n$ by minimizing the $\chi^{2}$ of the observed intensity
ratios versus the predicted model intensity ratios, weighted by the inverse
square of the observational errors. The second method is a Markov Chain Monte
Carlo (MCMC) method to determine the posterior probability density functions
(PDFs) of $G_{\rm 0}$ and $n$. The MCMC results can been visualised through
corner plots (Foreman-Mackey, 2016), which display the PDFs of $G_{\rm 0}$ and
$n$ and the sampling distribution. As these distributions are approximately
Gaussian, we characterise them on the overlay plots as 1 sigma contours in the
estimated $G_{\rm 0}$ and $n$ values for each region (Figures 4 and 5). As the
estimated $n$ and $G_{\rm 0}$ from the two methods are consistent, we report
the values derived by the MCMC method.

Although all intensity ratios enter into the calculation of the physical
parameters, the ratios with the smallest error bar (the thinnest curve in
overlay plots) dominate the fits. It can be seen in Figs. 4 and 5 that the [C
II] /CO ratio plays the most important role in determining the physical
conditions for different regions. Also, in certain cases, two points of
intersection exist, but the [C II] /CO dominates the determination of the
physical conditions. The results are summarised in Table 2.

###  6.2 Temperatures and pressures

Using the derived $n$ and $G_{\rm 0}$ values for different regions, we
determined the PDR surface temperatures ($T_{\rm surf}$ from the $T_{\rm
surf}$ map in the PDR Toolbox), which are the model gas temperatures at
$A_{\rm v}$ = 0.01 and are characteristic of the [C II] and [O I] emitting
regions (Table 2). We also estimated the [C II] excitation temperatures
($T_{\rm ex}$) assuming an optical depth of $\sim$ 3, which should be
considered an upper limit for the entire mapped region of RCW 49 (Tiwari et
al., 2021). The calculated $T_{\rm ex}$ ([C II] ) are given in Table 2. It can
be seen that the $T_{\rm ex}$ ([C II] ) values are less than the $T_{\rm
surf}$ values. This could be because [C II] is actually optically thin towards
most of the lines of sight and a derivation assuming it to be optically thick
results in lower excitation temperatures. Alternatively, a low excitation
temperature may indicate a density below the critical density.

We also calculated thermal ($p_{\rm th}$) and turbulent ($p_{\rm turb}$)
pressures (Table 2) in the different regions of RCW 49 using:

$p_{\rm th}=nT_{\rm surf}\hskip 2.84526pt({\rm K\,cm^{-3}}),$ (1)

and,

$p_{\rm turb}=\mu mn\left[\frac{\Delta v^{2}_{\rm
turb}}{8\ln(2)k}\right]\hskip 2.84526pt({\rm K\,cm^{-3}}),$ (2)

where

$\Delta v^{2}_{\rm turb}=\Delta v^{2}_{\rm FWHM}-[8\ln(2)kT_{\rm surf}/m_{\rm
c}].$ (3)

Here, $n$ is density (values from Table 2), $T_{\rm surf}$ is surface
temperature (values in Table 2), $\mu$ = 1.3 is the mean molecular weight, $m$
is the hydrogen mass, $\Delta v_{\rm FWHM}$ is full-width at half-maximum line
width of [C II] emission spectra (shown in Fig. 3 and summarised in Table 2),
$k$ is the Boltzmann constant and $m_{\rm c}$ is the carbon atom’s mass. For
determining $p_{\rm turb}$, we assumed $\Delta v^{2}_{\rm turb}$ to be
dominated by turbulence. However, the broad line widths of [C II] could also
be affected by velocity crowding along a line of sight. Thus, the $p_{\rm
turb}$ estimated here should be considered as upper limits.

We can see that the turbulent pressures are higher than the thermal pressures
and thus dominate the dynamics of the gas in all regions of RCW 49. This is
similar to the findings towards the stellar feedback driven shells of RCW 49
and RCW 36 (Tiwari et al. 2021 and Bonne et al., submitted). When comparing
the $p_{\rm th}$ and $p_{\rm turb}$ values derived for the shell in this work
with that from Tiwari et al. (2021), we find that the values estimated here
are lower by up to a factor of $\sim$ 1.6. This difference is because in
Tiwari et al. (2021), we derived the physical parameters using the average
intensity values of [C II] and 12CO towards the entire shell of RCW 49, while
in this work, we used the observed intensity values of [C II] , [O I] , 12CO
and FIR towards a specific line-of-sight. Another reason is that in Tiwari et
al. (2021), the $G_{\rm 0}$ values towards the shell were estimated using the
synthetic spectra for the stellar population of RCW 49 using the PoWR stellar
atmosphere grids (Sander et al., 2015), while in this work we used the PDR
ToolBox to derive $G_{\rm 0}$ values. Also, the PDR models we used in this
work are newly updated (“wk2020”) compared to the “wk2006” models used in
Tiwari et al. (2021).

###  6.3 Hydrogen nucleus column densities

We estimated the dust column densities towards different regions of RCW 49
from the dust spectral energy distribution (see Tiwari et al. 2021 for
details). The derived dust optical depths are converted to $N_{\rm H}$ (Table
2) using the Draine (2003) $R_{\rm v}$ = 3.1 value of the dust extinction
cross section per hydrogen nucleus at 160 $\mu$m. We also calculated $N_{\rm
H}$ using the 13CO data towards the same lines of sight. We first estimated
the excitation temperature using the optically thick 12CO data, which can then
be used to determine 13CO optical depths. Finally, using the 13CO integrated
intensities along with the excitation temperatures and optical depths, we
estimated the 13CO column densities, which can be converted to H2 column
densities (Table 2), adopting a H2/13CO abundance of 6.1 $\times$ 105 (Milam
et al., 2005; Tielens, 2005). The column densities derived from the dust are
in general higher than the ones derived from 13CO. This may indicate the
presence of CO-dark gas along the line of sight. Alternatively, column density
estimation from 70 and 160 $\mu$m data can include foreground and background
contribution of up to 30% (Tiwari et al., 2021).

Figure 6: The $n$ and $G_{\rm 0}$ model phase-space for [C II] , [O I] , 12CO
and FIR intensities. left panel: modeled phase-space with respect to [C II]
/CO(3-2) vs. [C II] /FIR. Right panel: modeled phase-space with respect to [C
II] /[O I] vs. [O I] \+ [C II] /FIR. Data points are shown in colored
(different regions) markers. The solid and dashed colored curves depict
constant $n$ and $G_{\rm 0}$ values, respectively.

Figure 7: The $n$ and $G_{\rm 0}$ model phase-space for [C II] , [O I] , 12CO
and FIR intensities. left panel: modeled phase-space with respect to [C II]
/CO(3-2) vs. [C II] /FIR. Right panel: modeled phase-space with respect to [C
II] /[O I] vs. [O I] \+ [C II] /FIR. Data points are shown in colored
(different regions) markers. The solid and dashed colored curves depict
constant $n$ and $G_{\rm 0}$ values, respectively.

###  6.4 Phase-space diagrams

To examine the dispersion in physical properties across the mapped regions, we
extracted all the data points from the [C II] , CO, [O I] and FIR intensity
maps and overlaid their ratios on the $n$ and $G_{\rm 0}$ modeled phase-space
(Figs. 6 and 7). These specific diagnostic plots were selected for this
analysis because the line ratios involved separate out $n$ and $G_{0}$
reasonably well into orthogonal directions. Specifically, the [C II] /CO (3-2)
and the [C II] /FIR ratios are sensitive to $G_{\rm 0}$ while the [C II] /[O
I] and ([O I] +[C II] )/FIR probe $n$. These diagnostic plots allow us to
generalize the results of the detailed analysis of the specific points in
Sects. 6.1 to the whole region.

All the data were regridded to the same angular resolution of $\sim$
20$\arcsec$ and filtered such that only data with values $>3\sigma$ are
included. For [C II] and CO, most of the data points had values $>3\sigma$.
However, for [O I] , the number of such data points decreased owing to lower
sensitivity, thus, we used upper limits in certain cases.

The individual points in these regions cover a wide range in $n$ and $G_{\rm
0}$ (Table 3). The estimated range of $G_{\rm 0}$ values derived from ratios
using [C II] , 12CO and FIR intensities is similar to those derived using [C
II] , [O I] and FIR intensities. However, the estimated $n$ values differ in
certain cases. Specifically, in p1, the $n$ values in the left panel are
higher by almost two magnitudes compared to the right panel. There are also
significant differences in the ridge and p7. This is most likely because of
the lower sensitivity of the [O I] data towards these regions, resulting in a
bias toward higher density regions.

Table 3: Physical conditions estimated by phase-space diagrams.

Parameter | units | northern cloud | ridge | shell | pillar | p1 | p3v1 | p3v2 | p7
---|---|---|---|---|---|---|---|---|---
$n$ ([C II] -CO-FIR) | 103 cm-3 | 5-10 | 1-100 | 0.5-5 | 1-100 | 10-100 | 1-10 | 1-10 | 1-10
$n$ ([C II] -[O I] -FIR) | 103 cm-3 | 5-7 | 1-10 | 1-10 | 1-100 | 0.1-1 | 0.5-5 | 0.5-5 | 0.1-70
$G_{\rm 0}$ ([C II] -CO-FIR) | 103 Habing units | 0.5-0.7 | 1-50 | 0.5-1 | 1-10 | 0.8-6 | 0.6-0.7 | 1 | 0.7-1
$G_{\rm 0}$ ([C II] -[O I] -FIR) | 103 Habing units | 0.5-0.7 | 3-10 | 1 | 0.7-10 | 1-4 | 0.6-1 | 1 | 1

Note. — Parameters with [C II] -CO-FIR correspond to the values estimated from
the left panels of Figs. 6 and 7. Similarly, parameters with [C II] -[O I]
-FIR correspond to the values estimated from the right panels of Figs. 6 and
7.

##  7 Discussion

###  7.1 Overlay plots vs. phase-space diagrams

In Sect. 6.1, we used the LSQ and MCMC method to find the best point of
intersection of all the observed ratios in different positions of RCW 49,
while in Sect. 6.4, we used phase-space diagrams on a modeled grid of
different combinations of ratios of [C II] , [O I] 12CO and FIR intensities.
Although the phase-space diagrams are a good way to visualize the entire data
($\sim$ 30–230 [C II] and 12CO spectra and $\sim$ 3–70 [O I] spectra for
different regions) in $n$–$G_{\rm 0}$ space and also visualizes gradients in a
specific region, the estimations are biased based on the combination of ratios
used for specific species. For instance, the [C II] /CO(3-2) ratio is more
sensitive to $G_{\rm 0}$, while the [O I] /[C II] ratio is more sensitive to
$n$. Another point to note is that observed [O I] observations have lower
sensitivity than the other tracers and thus, the phase-space probed by the [O
I] /[C II] ratio will have larger errors compared to the phase-space probed by
the [C II] /CO(3-2) ratio. In contrast, deriving the physical parameters using
all the ratios (four in Fig. 4 and five in Fig. 5) simultaneously towards a
representative position in different regions (as done in Sect. 6.1) should be
considered a more accurate approach for PDR analysis. Another way to compare
the physical parameters derived from the overlay plots and the phase-space
diagrams is by identifying the specific data points used in the overlay plots
(from Fig. 4 and 5) in the phase-space diagrams. These figures are shown and
discussed in detail in Appendix B and in Figs. 11 and 12. We find that the
data points used to do the analysis of the overlay plots lie on the extreme
ends of the phase-space probed by the entire data sets of different regions,
which might imply that the results obtained from the overlay plots are not
characteristic of the entire regions. However, as mentioned in Sect. 5, these
specific lines of sight were chosen for the overlay plots to ensure that clump
emission is excluded, which can interfere with the PDR analysis and give
misleading results. Thus, we believe that the overlay plots (results
summarised in Table 2) towards these representative lines of sight are a
better way to estimate the physical parameters and we use them as a guide to
understand stellar feedback in RCW 49.

Figure 8: [C II] modeled intensity map as a function of $G_{\rm 0}$ and $n$.
The map is highlighted with the observed intensity contours as well as marked
with the model predicted intensity values for different regions in different
colors. The thickness of the contours correspond to their observational error
bars.

###  7.2 Observed versus model intensities

As described in Sect. 6.1, the PDR Toolbox derives $n$ and $G_{\rm 0}$ (given
in Table 2) from the ratios of [C II] , [O I] , CO(3-2), H2 0-0 S(1), H2 0-0
S(2) and FIR intensities. We will now compare the individual observed
intensity values of different species with those predicted by the models. For
this, we extracted the intensity values of [C II] , CO, [O I] , H2 0-0 S(1)
and H2 0-0 S(2) from the “wk2020” PDR Toolbox models for the derived $n$ and
$G_{\rm 0}$ values for different regions. We can see from the Table 1 that the
model predicted intensity values of different species differ from their
observed values. The differences are most significant for p3 and p7 positions.

These differences can be visualised in Fig. 8, where the [C II] modeled
intensity map is highlighted with the observed intensity contours along with
the predicted [C II] intensities for different regions in RCW 49. It is clear
that the observational errors do not explain the difference between the
predicted and observed [C II] intensity values.

We ascribe these differences to the effects of geometry. The current PDR
models assume a face-on geometry and we can check results for an edge-on
geometry (see Pabst et al., 2017). In this view, the radiation is incident
from the side, and the PDR layers are spread across the sky with the atomic
gas closest to the incident radiation and the molecular gas at greater
distance. We calculate the line intensities from the (face-on) local
emissivities integrated along the line of sight while accounting for the line
optical depths. At a typical radiation field and density found in Table 2
($G_{\rm 0}$ $\sim$ 3 $\times$ 103 Habing units and $n$ $\sim$ $\times$ 103
cm-3), the [C II] intensity is about a factor of 2 higher in the edge-on case
compared to face-on due to the increased column density along the line of
sight. The [O I] and 12CO lines are relatively unchanged since they are
already optically thick and increasing the column density does little to
increase the emitted line intensity. However, this would lead to disagreement
with the observed intensity ratios. Since the density is greater than the [C
II] critical density but less than the [O I] and 12CO critical densities, an
increase in the density by a factor of 2 does little to change [C II] but
increases [O I] by a factor of $\sim$ 2 and 12CO by a factor of $\sim$ 2-3.
Combining these two effects, the observed ratios can be explained by an edge-
on model with a factor of $\sim$ 2 higher density. Thus, the model fits to
some of the regions could be improved and geometry introduces an uncertainty
by a factor of ${\sim}2$ in density. Position p3 remains problematic with a
high observed [C II] intensity. Such high intensities $>10^{-3}$ erg ${\rm
cm^{-2}}$ ${\rm s^{-1}}$ ${\rm sr^{-1}}$ typically indicate a large
contribution from ionized gas along the line of sight (Seo et al., 2019).

Overall, we believe that geometrical effects are the main cause for the
difference between the predicted and the observed intensity values for various
species.

Figure 9: $p_{\rm th}$ vs $G_{\rm 0}$ for different regions. The red, green,
magenta and dark blue points correspond to regions studied by Young Owl et al.
(2002) (using FIR lines analysis), Joblin et al. (2018) (using high-$J$ 12CO
analysis), Pabst et al. (2021) (using low-$J$ 12CO and [C II] analysis) and
this work, respectively. p3v1 and p3v2 correspond to the two velocity
components of p3: -15 to 5 km s-1 and 5 to 20 km s-1, respectively. The black
line refers to the fit by Wu et al. (2018) using mid to high-$J$ 12CO and [C
I] analysis. The light blue solid and dashed lines are for the analytic
relation by Seo et al. (2019) and the relation divided by 3 to account for
pressure equipartition, respectively.

###  7.3 Relationship between $p_{\rm th}$ and $G_{\rm 0}$

We studied the relationship between $p_{\rm th}$ and $G_{\rm 0}$ for different
regions of RCW 49 (Table 2) and for various PDRs studied in the literature.
Figure 9 is adapted from Pabst et al. (2021, Fig. 20). We can see that all
studies follow the general trend, where $p_{\rm th}$ increases with $G_{\rm
0}$. However, studies by Young Owl et al. (2002) (red data points in Fig. 9)
and Joblin et al. (2018) (green data points in Fig. 9) are somewhat shifted
when compared to the regions of RCW 49 owing to higher $G_{\rm 0}$ and $p_{\rm
th}$ values, respectively. Young Owl et al. (2002) estimated higher values of
$G_{\rm 0}$ for their sources compared to the values determined by Salgado et
al. (2016), Joblin et al. (2018) and Pabst et al. (2021). Joblin et al. (2018)
used high-$J$ 12CO lines for their analysis and therefore probe high density
gas and/or clumps due to high critical densities of these lines. The relation
between $p_{\rm th}$ and $G_{\rm 0}$ as derived by Wu et al. (2018) (black
line in Fig. 9) also derived from high-$J$ 12CO lines analysis and is biased
towards high density clumps.

There is a simple relation (shown with solid light blue line in Fig. 9)
between $p_{\rm th}$ and $G_{\rm 0}$ for a PDR in pressure equilibrium within
a Stromgren sphere (Young Owl et al., 2002; Seo et al., 2019). The different
regions of RCW 49 (dark blue data points in Fig. 9) lie close to the above
analytical relation taking equipartition of thermal, turbulent, and magnetic
pressures into account. This is similar to the sources studied by Pabst et al.
(2021). Thus, the physical conditions in different regions of RCW 49 are in
agreement with basic Strömgren relation for an H II region.

###  7.4 Impact of stellar feedback on different regions

In this section, we compare different regions of RCW 49 based on the derived
physical conditions. The derived physical conditions in the seven regions of
RCW 49 are indicative of different effects of stellar feedback on these
regions. As justified earlier, we use the analysis reported in Sect. 6.1 for
the comparison. We compare the regions p1, p3 and p7 separately from the
northern cloud, ridge, shell and pillar because analysis of the former
includes H2 emission along a line of sight, while analysis of the latter
excludes H2 emission and includes emission only from specific velocity
components corresponding to the given structure.

For $G_{\rm 0}$ values, the different regions can be placed in an increasing
order: northern cloud with the lowest $G_{\rm 0}$, followed by the pillar,
shell and the ridge with the highest $G_{\rm 0}$. Similarly, for the density,
we can place the different regions in an increasing order: pillar with the
lowest $n$, followed by the shell and ridge and the northern cloud with the
highest $n$.

The northern cloud is located farthest from the Wd2 cluster, likely the reason
that it has the lowest $G_{\rm 0}$ of the regions studied. The ridge, which is
$\sim$ 40 pc away from the northern cloud (Furukawa et al., 2009) is closest
to the Wd2 cluster, corresponding to the highest $G_{\rm 0}$ value. Moreover,
Furukawa et al. (2009) suggested that the northern cloud and the ridge
collided $\sim$ 4 Myrs ago forming the Wd2 cluster (age $\sim$ 2–3 Myrs,
Furukawa et al. 2009; Tiwari et al. 2021). The compression associated with
this collision may be the cause for the higher densities derived for these
regions. The high density of the ridge is also consistent with the ongoing
star formation reported in it (Whitney et al., 2004; Tiwari et al., 2021).
Also, the derived densities of the northern cloud and the ridge in this work
are similar to the ones (3–8 $\times$ 103 cm-3) derived by Ohama et al. (2010)
using the LVG analysis of molecular emission lines. The mechanical feedback of
the Wd2 cluster drives the expanding (at $\sim$ 13 km s-1) shell of RCW 49,
which is $\sim$ 6 pc from the cluster. A secondary (younger and lower in mass)
generation of star formation is taking place in the shell of RCW 49 (as
discussed in Whitney et al. 2004; Tiwari et al. 2021). The pillar has similar
$n$ and $G_{\rm 0}$ values (1.5 and 1.6 times lower) as the shell. From the
pillar’s morphology, it seems to be created by the Wd2 cluster. However, it
also has the Wolf-Rayet star, WR20b, relatively closer to it and that could
also be responsible for illuminating the pillar. Currently, we are unable to
quantify the effects of radiative feedback from WR20b in defining the physical
conditions of the pillar. But comparing the masses and bolometric luminosities
of WR20b with that of Wd2 cluster, the stellar (radiative and mechanical)
feedback from WR20b should be not more than 20% of that from Wd2.

Among p1, p3 and p7, p7 has the highest $G_{\rm 0}$ value, while the -15 to 5
km s-1 velocity component of p3 has the lowest $G_{\rm 0}$ value. And p1 has
the highest $n$ value and again the -15 to 5 km s-1 velocity component of p3
has the lowest $n$ value. P7 is the northernmost part of the shell of RCW 49
and in the west of it, the shell gets broken. We also located an early type
O9V star close to this line of sight. Thus, the high $G_{\rm 0}$ value in p7
could be due to its proximity to this star. p1 is closest to the Wd2 cluster
and has dense gas towards it (bright 12CO and 870 $\mu$m emission), suggesting
a possibility of triggered star formation in its early stage. Towards position
p3, we see a superposition of the shell (-15 to 5 km s-1) and the ridge (5 to
20 km s-1) components. The $G_{\rm 0}$ values for the shell component are
lower than that of the ridge. This can be explained from their location with
respect to the Wd2 cluster. As mentioned in the previous paragraph, the ridge
is close to Wd2 while, the shell is about 6 pc away from it. The ridge
component is also denser compared to the shell component in p3, which is
consistent with its ongoing star formation.

##  8 Conclusions

We studied seven different regions of RCW 49 using the [C II] , [O I] , 12CO
(3-2), H2 0-0 S(1) and H2 0-0 S(2) observations. Four of these regions: the
northern cloud, ridge, shell and pillar have the [C II] , [O I] and 12CO (3-2)
data towards them observed through the FEEDBACK program. These regions are
both spatially and spectrally distinct. The other three regions: p1, p3 and p7
have the H2 0-0 S(1) and H2 0-0 S(2) data observed by the Spitzer telescope in
addition to the FEEDBACK data. We presented the [O I] and H2 data (spectra and
emission maps) towards RCW 49 for the first time in this paper. We found that
the H2 0-0 S(1) emission distribution follows that of [C II] while the H2 0-0
S(2) emission distribution follows that of CO, indicating that H2 0-0 S(2)
arises from a denser gas compared to H2 0-0 S(1) in RCW 49.

To determine the physical conditions in different regions of RCW 49, we
compared our observations with the updated PDR models. We justified that the
PDR analysis done using the overlay plots and data fits is a better technique
than using phase-space diagrams to derive the physical conditions in the ISM.
Based on the physical conditions, we studied the effects of stellar feedback
on the evolution of different regions in RCW 49. We found that the ridge
(closest to Wd2) has the highest $G_{\rm 0}$ value, while the northern cloud
(farthest from Wd2) has the lowest $G_{\rm 0}$. However, both the northern
cloud and the ridge have high densities, which is consistent with previous
studies suggesting that these regions are part of large scale clouds, whose
collision led to the formation of the Wd2 cluster. Also, there is evidence of
ongoing star formation in the ridge. Among p1, p3 and p7, p1 has highest
density and based on the bright 12CO and 870 $\mu$m emission towards it, we
suggest early stage star-formation in this region. p7 has highest the FUV flux
and this is attributed to the impinging radiation from an early type O9V star
close to this position.

We also estimated pressures towards these regions and found that the $p_{\rm
turb}$ dominate (when compared to $p_{\rm th}$) the dynamics of the gas in RCW
49. Furthermore, the observed relationship between $p_{\rm th}$ and $G_{\rm
0}$ follow the theoretical relationship derived from the Strömgren relation
for an H II region, which has also been observed toward several other PDRs.

We would like to thank the referee for their constructive comments and helping
to clarify the paper.

This study was based on observations made with the NASA/DLR Stratospheric
Observatory for Infrared Astronomy (SOFIA). SOFIA is jointly operated by the
Universities Space Research Association Inc. (USRA), under NASA contract
NNA17BF53C, and the Deutsches SOFIA Institut (DSI), under DLR contract 50 OK
0901 to the University of Stuttgart. upGREAT is a development by the MPI für
Radioastronomie and the KOSMA/University of Cologne, in cooperation with the
DLR Institut für Optische Sensorsysteme.

Financial support for the SOFIA Legacy Program, FEEDBACK, at the University of
Maryland was provided by NASA through award SOF070077 issued by USRA. M.W.P.
acknowledges support for development of the PDR Toolbox from NASA ADAP
#80NSSC19K0573.

The FEEDBACK project is supported by the BMWI via DLR, Project Number 50 OR
1916 (FEEDBACK). N.S., R.S., and L.B. acknowledge support by the Agence
National de Recherche (ANR/France) and the Deutsche Forschungsgemeinschaft
(DFG/Germany) through the project “GENESIS” (ANR-16-CE92-0035-01/DFG1591/2-1).

This work was partly supported by the Collaborative Research Centre 956,
funded by the DFG.

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##  Appendix A IR spectra

Figure 10 displays the average spectra in $\sim$ 10–20 $\mu$m wavelength
window towards p1, p3 and p7. H2 lines are brightest in p3 by $\sim$ 1.5–3
times compared to p1 and p7. This is similar to our [C II] , [O I] and 12CO
emission, which is brightest towards the ridge. The H2 0-0 S(1) line is
similarly intense in p1 and p7, however the H2 0-0 S(2) line is $\sim$ 2 times
brighter in p1 than in p7. This can be attributed to the fact that p1 is
closer to Wd2 than p7 and is thus exposed to stronger radiation to heat the
gas and excite the H2 0-0 S(2) line, which has an higher upper level energy.

For the ions, [Ne II] and [S III] are brightest in p3 and least bright in p7.
This is similar to the H2 lines. However, the relatively more ionised lines of
[Ne III] and [S IV] are brightest in p1 and least bright in p7. This can again
be explained because of the location of these regions with respect to Wd2: p1
is closest and p7 is farthest.

Figure 10: Average IR spectra towards p1, p3 and p7 (regions shown in Fig. 1).
The six species presented in this work are marked with red dashed lines.

Figure 11: The $n$ and $G_{\rm 0}$ model phase-space for [C II] , [O I] , 12CO
and FIR intensities. left panel: modeled phase-space with respect to [C II]
/CO(3-2) vs. [C II] /FIR. Right panel: modeled phase-space with respect to [C
II] /[O I] vs. [O I] \+ [C II] /FIR. Data points are shown in colored
(different regions) markers. The shaded regions correspond to the spread in
the phase-space as probed by the entire dataset towards different regions as
shown in Fig. 6. The solid and dashed colored curves depict constant $n$ and
$G_{\rm 0}$ values, respectively.

Figure 12: The $n$ and $G_{\rm 0}$ model phase-space for [C II] , [O I] , 12CO
and FIR intensities. left panel: modeled phase-space with respect to [C II]
/CO(3-2) vs. [C II] /FIR. Right panel: modeled phase-space with respect to [C
II] /[O I] vs. [O I] \+ [C II] /FIR. Data points are shown in colored
(different regions) markers. The shaded regions correspond to the spread in
the phase-space as probed by the entire dataset towards different regions as
shown in Fig. 7. The solid and dashed colored curves depict constant $n$ and
$G_{\rm 0}$ values, respectively.

##  Appendix B Phase-space diagrams for specific data points

Figures 11 and 12 show phase-space diagrams overlaid with specific data points
(colored crosses) selected for the PDR analysis presented in Sect 6.1 and
Figs. 4 and 5. These data points mostly lie close to the boundaries of the
shaded areas (corresponding to the space probed by the entire data sets for
different regions) in the phase-space diagrams. For instance, in the left
panel of Fig. 11, the data points used to make the overlay plot for the ridge
and pillar are probing lower values of $G_{\rm 0}$ compared to its shaded
area, while for the shell, it is probing a higher value of $G_{\rm 0}$.
Moreover, the data point corresponding to the pillar in the right panel of
Fig. 11 is probing lower value of $n$ compared to its shaded area. Another
clear example can be seen in the right panel of Fig. 12, where the data point
corresponding to p7 probes lower density compared to its shaded area.