II]) (K km s-1). The resulting map is presented in Fig. 20. Comparing this map
with Fig. 2 in Suzuki et al. (2021), it is found that the structure and
observed intensity of the region covered so far by the FEEDBACK observations
is extremely similar. This is an important validation for the data presented
in this paper.

Figure 20: The FEEDBACK [C II] integrated intensity map smoothed to a spatial
resolution of 90′′. It looks exactly like the 90′′ map presented in Fig. 2 of
Suzuki et al. (2021).

##  Appendix B [C II] self-absorption demonstrated with [13C II]

To verify whether [C II] self-absorption affects the observed spectrum, the
three [13C II] hyperfine structure lines were used as they are expected to be
optically thin (Ossenkopf et al., 2013). As the [13C II] hyperfine structure
lines are at a relative velocity of -65.2 km s-1 (F=1-0), 11.7 km s-1 (F=2-1)
and 62.4 km s-1 (F=1-1) compared to the [C II] line, these three hyperfine
transitions are also covered by the upGREAT receiver. The brightest transition
(F=2-1) is closest to the [C II] and contaminated by the [C II] outflow wing
for RCW 36. The F=2-1 transition is presented in Fig. 5 and the two other
hyperfine transitions are presented in Fig. 21. These figures show, as far as
possible for F=1-0 and F=1-1 because of the noise, that there is one main
velocity component at 7 km s-1 which leads to strong absorption that is seen
with [C II], [O I], 12CO(3-2) and even a bit with 13CO(2-1). This result is
also consistent with Fig. 2 in Fissel et al. (2019) which shows that the
optically thin N2H+(1-0) dense gas tracer is centered on 7 km s-1.
To visualize the [O I] absorption, the spectra were integrated between 3 and 7
km s-1 (the velocities where the [O I] emission goes under the baseline in
Fig. 2). The resulting map is shown in Fig. 22 which shows that the strongest
absorption is located at the densest clumps of the ring.

Figure 21: The [C II] line compared with the [13C II] F=1-0 and F=1-1
hyperfine transitions for the ring of RCW 36. Both [13C II] transitions are
detected and indicate that the [13C II] emission is centered on the dip at 7
km s-1 of the [C II] line, implying self-absorption of the [C II] line. Figure
22: The [O I] integrated intensity map from 3 to 7 km s-1 with the white
contours indicating the column density. Towards the dense ring, and in
particular the highest column density region, there is absorption in this
velocity interval.

##  Appendix C Estimation of the FUV field and [C II] temperature

In order to estimate a temperature for the gas traced by [C II] we will work
with the PDR Toolbox (Kaufman et al., 2006; Pound & Wolfire, 2008), but first
we must estimate the strength of the FUV field. Two methods are used to
estimate the field strength independently. First we use the method presented
in Kramer et al. (2008); Schneider et al. (2016) which makes use of the
Herschel intensity at 70 $\mu$m and 160 $\mu$m to estimate FUV field in Habing
(Habing, 1968) using

$\text{F}_{\text{FUV}}[\text{G}_{0}]=4000\pi\text{I}_{\text{FIR}}/1.6$ (C1)

where I${}_{\text{FIR}}$ is the far-infrared intensity from adding the
intensity at 70 $\mu$m and 160 $\mu$m. The resulting map of the estimated FUV
field strength is presented in Fig. 23 which shows that the FUV field strength
is generally between G0 = 300-2000 for RCW 36 within 1-2 pc from the cluster
center. This method assumes that the [C II] emission originates from roughly
the same region as this warm dust emission, which likely is a relatively good
assumption.
The second method to calculate the FUV field at a specific location, presented
in Bonne et al. (2020b), directly works with the ionizing radiation intensity
from the O stars in the central cluster. The FUV irradiation from these O
stars is calculated from the Kurucz (1979) models with solar abundances
(Grevesse & Sauval, 1998). In order to select the Kurucz models to calculate
the FUV field, the stellar parameters for the spectral types of the O stars
were taken from Pecaut & Mamajek (2013). This results in an FUV field of G0 =
1.9$\times 10^{3}$ at a distance of 1 pc and G0 = 4.7$\times 10^{2}$ at a
distance of 2 pc. These calculations of the FUV field take no extinction on
the path of the FUV photons in account, but fit well with the values obtained
from the previous method which suggests that the FUV field at the expanding
shells in the cavities is of the order of G0 = 500-2000.
From this information it is possible to estimate a temperature in the PDR
using the PDR Toolbox. The surface temperature of the PDR as a function of the
density is presented in Fig. 24 which indicates that a surface temperature of
100-500 K is expected for a FUV field strength between G0 = 500-2000 and
densities n = 102-105 cm-3 (Minier et al., 2013). This surface temperature of
the PDR corresponds to the temperatures traced by [C II]. To verify this
method, it is also found with the PDR Toolbox that the observed [C II]
intensity of $\sim$ 100 K km s-1 towards the shells of RCW 36 fits with
G${}_{0}\sim$ 103 and n $\sim$ 103-104 cm-3 which is expected for the
expanding shells.

Figure 23: The FUV field map of RCW 36 on a log scale calculated from the
Herschel 70 $\mu$m and 160 $\mu$m maps. The blue crosses indicate the
locations of the O stars. Figure 24: The surface temperature as a function of
the density and FUV field in the PDR Toolbox (Kaufman et al., 2006; Pound &
Wolfire, 2008). Working with G0 = 500-2000 and possible densities in the
shells and cavities (102 and 105 cm-3) points to temperatures between 100 and
500 K for [C II].

##  Appendix D Fitting the X-ray spectra

The fitting of the X-ray spectra, presented in Fig. 25, uses the XSPEC
package777https://heasarc.gsfc.nasa.gov/docs/xanadu/xspec/manual/manual.html
(Arnaud, 1996), and takes into account diffuse emission components from hot
plasmas, unresolved stellar clusters, background emission and a constant
absorption from foreground material over a considered field of view. When
fitting a hot plasma the chemical abundance values from Wilms et al. (2000)
are assumed. The fitting results are presented in Tab. 8 and Fig. 25.
The spectrum of the central region around the OB cluster is best fitted by a
hot plasma and an unresolved stellar cluster. The unresolved cluster has an
X-ray luminosity of 0.17 L⊙ between 0.5 and 7 keV. Converting this to a
bolometric luminosity using log(Lx/Lbol) = -3.6 (Bhatt et al., 2013), gives a
a bolometric luminosity of 6.8$\times$102 L⊙. The fitted unresolved cluster
can thus correspond to low-mass stars from the 1.3$\times$105 L⊙ OB cluster
(Verma et al., 1994) that were not resolved with ACIS. For the fitted hot
plasma, the ionization timescale $\tau$ can range between 108 and
5$\times$1013 s cm-3. The lower limit indicates a plasma with non-equilibrium
ionization (NEI) and the upper limit indicates a plasma with collisional
ionization equilibrium (CIE). For the central region, $\tau$ thus indicates a
low to intermediate timescale.
The other four regions contain no unresolved cluster and are instead fitted by
two hot plasma components because a single hot plasma component did not give a
satisfactory fit to the spectra. During the fit of the two plasma components
in the bipolar cavities, the first plasma component went to small ionization
timescales and the second one to large ionization timescales, while
temperatures of both hot plasma components converged to the same temperature.
Therefore, in a second run, the temperatures for both plasmas were imposed to
be the same. The presence of NEI and CIE with the same temperature in the fit
indicates a recombining plasma that is in transition between from NEI to CIE.
From the larger surface brightness in both cavities for the CIE components, it
is found that the plasmas are dominated by CIE. In the region outside the H II
region, i.e. north and southeast, the temparatures of both plasma components
did not converge. The first component has a temperature similar to the ones in
the cavities, and is likely the hot plasma that is leaking from the bipolar H
II region. The second plasma component is slightly hotter and is probably
fore- or background emission as RCW 36 is still relatively close to the
galactic plane. The surface brightness indicates that the leaking plasma is
dominating the emission.
The fitting also shows that the observed diffuse X-ray emission is
significantly impacted by absorption from foreground layers with N${}_{H}\sim$
5$\times$1021 cm-2. This is particularly seen in column 23 of Tab. 8, where it
is found that the intrinsic surface brightness of both cavities, corrected for
absorption, is higher than for the regions outside the H II region.
Interestingly, the fitted absorbing column density value for the eastern
cavity is extremely similar to the typical column density found in the
Herschel maps. For the western cavity the fitting had more difficulties to
constrain the typical column density, so we decided to supply the foreground
column density based on the Herschel data.

Figure 25: Chandra/ACIS spectra of diffuse X-ray emission in RCW 36, in the
regions defined by Fig. 14 and repeated in the upper left panel here. This
panel also demonstrates that these regions are exclusive, with no X-ray event
appearing in more than one spectrum: Center and North events are color-coded
in blue, Western Cavity and Southeast events are in green, and Eastern Cavity
events are in red. The spectra show the total model in black, the first
diffuse thermal plasma component in red, the second diffuse thermal plasma
component in green, and background emission in magenta. For the Center region
(upper right panel), cyan curves model the emission from unresolved pre-main
sequence stars in the young cluster. Lower panels in all spectra show the fit
residuals. Spectral fit parameters and model details are given in Tab. 8.
Table 8: Spectral Fits for Diffuse X-ray Regions

Extraction Region | Quality | Diffuse Components | Unresolved Stars | Background |  | Total Diffuse |  |
---|---|---|---|---|---|---|---|---
|  |  | pshock1 | pshock2 |  | apec4+apec5 | apec6 |  | Emission |  |
Name | SNR | Area |  | $\chi^{2}\\!/\\!\rm{DOF}$ |  |  |  |  |  |  | $N_{H}^{1}$ | $kT^{1}$ | ${\tau}^{1}$ | $SEM^{1}$ | $SB^{1}_{tc}$ | $N_{H}^{2}$ | $kT^{2}$ | ${\tau}^{2}$ | $SEM^{2}$ | $SB^{2}_{tc}$ |  |  |  |  |  | $N_{H}^{stars}$ | $SEM^{5}$ | $SB^{stars}_{tc}$ | $L^{stars}_{tc}$ | $N_{H}^{6}$ |  | $SEM^{6}$ | $SB^{6}_{tc}$ |  |  |  |  | $SB^{1+2}_{tc}$ | $L^{1+2}_{tc}$ |  |
| (cts) | (pc2) |  |  |  | ( $N_{H}$: $10^{22}$cm-2 $kT$: keV $\tau$: log cm-3 s $SEM$: log cm-3 pc-2 $SB$: log erg s-1 pc-2 $L$: log erg s-1 ) |  |
(1) | (2) | (3) |  | (4) |  |  |  |  |  |  | (5) | (6) | (7) | (8) | (9) | (10) | (11) | (12) | (13) | (14) |  |  |  |  |  | (15) | (16) | (17) | (18) | (19) |  | (20) | (21) |  |  |  |  | (22) | (23) |  |
Center | 1650/61 | 0.9 |  | 54/ 53 |  |  |  |  |  |  | $0.55\phantom{.}_{-0.3}^{+0.3}$ | $0.23\phantom{.}_{-0.09}^{+0.37}$ | ${\phantom{0}}8.98\phantom{.}_{-0.4}^{+0.3}$ | $55.9\phantom{.}_{\cdots}^{+1.5}$ | 32.26 | $\cdots$ | $\cdots$ | $\cdots$ | $\cdots$ |  $\cdots$ |  |  |  |  |  | $2.8\phantom{.}_{-0.8}^{+0.9}$ | $54.7\phantom{.}_{-0.09}^{+0.08}$ | 31.88 | 31.82 | $4.8\phantom{.}\\!*$ |  | $53.9\phantom{.}\\!*$ | 31.10 |  |  |  |  | 32.26 | 32.20 |  |
Eastern Cavity | 4240/126 | 4.7 |  | 77/ 62 |  |  |  |  |  |  | $0.51\phantom{.}_{-0.07}^{+0.11}$ | $0.15\phantom{.}_{-0.02}^{\cdots}$ | ${\phantom{0}}8.00\phantom{.}\\!*$ | $56.5\phantom{.}_{\cdots}^{+0.4}$ | 31.69 | $0.51\phantom{.}\\!*$ | $0.15\phantom{.}\\!*$ | $13.7\phantom{.}\\!*$ | $55.3\phantom{.}_{-0.3}^{+0.3}$ | 32.00 |  |  |  |  |  | $\cdots$ | $\cdots$ |  $\cdots$ |  $\cdots$ | $2.4\phantom{.}\\!*$ |  | $53.9\phantom{.}\\!*$ | 31.10 |  |  |  |  | 32.17 | 32.85 |  |
Western Cavity | 3493/132 | 6.5 |  | 60/ 58 |  |  |  |  |  |  | $0.11\phantom{.}_{\cdots}^{+0.3}$ | $0.17\phantom{.}_{-0.02}^{+0.03}$ | ${\phantom{0}}8.66\phantom{.}_{-0.5}^{+0.3}$ | $55.3\phantom{.}_{\cdots}^{+0.4}$ | 30.92 | $0.80\phantom{.}\\!*$ | $0.17\phantom{.}\\!*$ | $13.7\phantom{.}\\!*$ | $55.1\phantom{.}_{-0.4}^{+0.3}$ | 31.91 |  |  |  |  |  | $\cdots$ | $\cdots$ |  $\cdots$ |  $\cdots$ | $2.0\phantom{.}\\!*$ |  | $53.9\phantom{.}\\!*$ | 31.10 |  |  |  |  | 31.95 | 32.76 |  |
Southeast | 5269/116 | 3.6 |  | 64/ 59 |  |  |  |  |  |  | $0.29\phantom{.}_{-0.1}^{+0.1}$ | $0.26\phantom{.}_{-0.08}^{+0.13}$ | $11.4\phantom{.}_{-0.4}^{+0.6}$ | $54.5\phantom{.}_{\cdots}^{+0.7}$ | 31.83 | $0.14\phantom{.}_{\cdots}^{+0.1}$ | $0.68\phantom{.}_{\cdots}^{\cdots}$ | $12.9\phantom{.}_{\cdots}^{\cdots}$ | $53.7\phantom{.}_{-0.1}^{+0.1}$ | 31.03 |  |  |  |  |  | $\cdots$ | $\cdots$ |  $\cdots$ |  $\cdots$ | $4.8\phantom{.}\\!*$ |  | $53.9\phantom{.}\\!*$ | 31.10 |  |  |  |  | 31.90 | 32.45 |  |
North | 2207/97 | 3.1 |  | 39/ 40 |  |  |  |  |  |  | $0.36\phantom{.}_{\cdots}^{+0.1}$ | $0.16\phantom{.}_{-0.02}^{+0.03}$ | $13.7\phantom{.}\\!*$ | $55.0\phantom{.}_{\cdots}^{+0.2}$ | 31.78 | $0.36\phantom{.}\\!*$ | $0.76\phantom{.}_{-0.2}^{\cdots}$ | $11.3\phantom{.}_{\cdots}^{\cdots}$ | $53.2\phantom{.}_{-0.3}^{+0.2}$ | 30.81 |  |  |  |  |  | $\cdots$ | $\cdots$ |  $\cdots$ |  $\cdots$ | $4.1\phantom{.}\\!*$ |  | $53.9\phantom{.}\\!*$ | 31.10 |  |  |  |  | 31.83 | 32.32 |  |

Properties of the diffuse regions:
Col. (1): Name of the diffuse extraction region shown in Figure 14.
Col. (2): Signal-to-noise ratio in the total energy band (0.5–7 keV),
expressed as $NetCounts/\sigma_{NetCounts}$.
Col. (3): Geometric area of the region in square parsecs, irrespective of
point source masking, assuming a distance of 900 pc.

Details of the spectral fits:
Col. (4): “Reduced $\chi^{2}$” = $\chi^{2}$ divided by “degrees of freedom”
for the best-fit model.
All fits used the source model “tbabs1*pshock1 \+ tbabs2*pshock2 \+
tbabsstars(apec4 \+ apec5) + tbabs6*apec6” in XSPEC.
$N_{H}$ columns report the best-fit value for the extinction column density
(tbabs components).
$kT$ columns report the best-fit value for the plasma temperature (pshock and
apec components).
$\tau$ columns report the ionization timescale for the plasma (pshock
components).
$SEM$ columns report the surface emission measure for the plasma (pshock and
apec components), which is independent of distance.
Unresolved pre-main sequence stars are modeled with a pair of soft (0.86 keV)
and hard (2.6 keV) thermal plasmas (apec4 \+ apec5) with a fixed normalization
ratio ($norm4=0.5\times norm5$). Background emission (apec6) is frozen at 10
keV; its absorption (tbabs6) is always larger than the largest tbabs needed
for other model components.

Quantities marked with an asterisk (*) were frozen in the fit, because the fit
parameter ran to one of its limits or because it was poorly constrained.
Uncertainties represent 90% confidence intervals. More significant digits are
used for uncertainties $<$0.1 in order to avoid large rounding errors; for
consistency, the same number of significant digits is used for both lower and
upper uncertainties. Uncertainties are missing when XSPEC was unable to
compute them or when their values were so large that the parameter is
effectively unconstrained. $SB_{tc}$ and $L_{tc}$ columns report the
absorption-corrected surface brightness and luminosity of model components in
the total band (0.5–7 keV). $SB$ is independent of distance; $L=SB\times
Area$.
$SB^{stars}$ and $L^{stars}$ are the surface brightness and luminosity,
respectively, of unresolved pre-main sequence stars, calculated as the sum
from components apec4 and apec5.
Col. (22): total surface brightness from both pshock components combined.
Col. (23): total luminosity from both pshock components combined.

##  Appendix E Calculating the pressure terms in RCW 36.

The hot plasma thermal pressure is obtained from fitting the X-ray spectra and
the thermal pressure information for the ionized gas in the center is obtained
from Minier et al. (2013). The thermal pressure for the ionized gas in the
cavities was derived from the SHASSA Hα data (Gaustad et al., 2001). From the
SHASSA data it is possible to calculate the emission measure (EM) using EM =
4.2$\times$1017 Iα after correction for the dust extinction, with Iα the
intensity in the SHASSA region. The correction for the dust exctinction uses
the curve from Calzetti et al. (2000) with k($\lambda$) = 2.38, AV = 3, based
on the Herschel column density towards the cavities, and RV = 3.1 (Draine,
2003). Assuming a physical size of 1 pc for the cavities then gives an
electron density of 58-62 cm-3. Combining this with an electron temperature of
8000 K then gives the thermal pressure. This method based on the SHASSA data
was not applied for the central region in the ring because the associated AV
cannot be accurately constrained. To calculate the ionizing photon flux from
the emission measure, we then use $\dot{N}_{ion}$ =
EM$\times$A$\times\alpha_{B}$. Here, A is the surface area and $\alpha_{B}$
the recombination coefficient (= 2.7$\times$10-13 cm3 s-1). The radiation
pressure was directly estimated using I/c, where I is the radiation intensity
and c the speed of light. The radiation intensity was determined from the
luminosity of the central cluster (1.3$\times$105 L⊙) and a distance from the
cluster of 0.9 pc for the molecular ring and a distance of 1.5 pc for the
cavity shells.
To calculate the ram pressure on the expanding shells, it was assumed that
they are expanding in the ambient cloud of Vela C with a density n${}_{H_{2}}$
= 102 \- 5$\times$102 cm-3 and a temperature of 50 K which is a typical
temperature for the cold neutral medium. This density was estimated using a
column density N${}_{H_{2}}$ = 3$\times$1021 cm-2 in the ambient cloud, based
on the Herschel data, and an estimated size in the line-of-sight of 2-10 pc
which is poorly constrained. Inside the compressed shells of the cavities, we
worked with n${}_{H_{2}}$ = 103-3$\times$103 cm-3 and a temperature of 250 K
that was determined from the PDR modeling. The density is estimated using a
column density N${}_{H_{2}}$ = 3$\times$1021 cm-2 and a depth of 0.3 to 1 pc.
To calculate turbulent pressure, the [C II] linewidth, associated with the
expanding shell (i.e. the blueshifted velocity component varying between 1 and
6 km s-1), was fitted. The details of this, both for the cavity shells and
ring, are presented in the next paragraph. To calculate the ram pressure on
the molecular ring, it is assumed that the ring is expanding into an ambient
cloud/sheet with n${}_{H_{2}}$ = 103-104 cm-3 and a temperature of 20 K which
is typical for these slightly denser molecular regions in the cloud. This
large range of densities was assumed as it is difficult to constrain the
density around the ridge observationally. For this estimate we assumed a
typical column density around the ridge of 1022 cm-2 and a potential size
range of 0.3 to 3 pc. To calculate the upper limit on the turbulent pressure
in the ambient cloud, the velocity dispersion of 2.0 km s-1 is taken (Fissel
et al., 2019). The assumed internal density of the ring is n${}_{H_{2}}$ =
2.4$\times$104 cm-3 based on a typical N${}_{H_{2}}$ = 1.5$\times$1022 cm-2
over the full ring and a diameter of 0.2 pc. The magnetic field strength in
the ambient cloud was calculated using the assumed density range and the
Crutcher (2012) relation. The magnetic field in the shell of the cavities was
estimated using equation 3 in van Marle et al. (2015), which assumes
(initially) super Alfvénic expansion speed, which is uncertain leading to an
upper limit. This assumption is likely not valid for the low expansion
velocity of the ring, such that it is not calculated there. Future work using
HAWC+ data will further address the role of the magnetic field in the ring
(Bij et al. in prep.).

To obtain the linewidth of the expanding shell in the cavities, the [C II]
velocity component between 0 and 6 km s-1 was fitted with a single Gaussian.
This fitting is complicated by contamination from the bright [C II] emission
at more redshifted velocities (the main velocity component associated with the
ring). Therefore, a fitted Gaussian profile was only accepted if the the peak
brightness is between 3.5 and 40 K. This peak brightness interval covers the
full observed brightness range for the velocity component below 6 km s-1,
associated with the expanding shell, based on a visual inspection. As a result
this constraint prevents that the Gaussian fitting tries to include additional
emission that is associated with the bright component at 7-8 km s-1 from the
ring. The resulting [C II] linewidth distribution for the expanding shells is
presented in Fig. 26.
Fitting the linewidth was also done for 13CO(3-2). The resulting linewidth
distribution for the detected emission over the full map is compared with the
linewidth in the region that is identified as the dense ring in Fig. 27. This
demonstrates that there is an increase in linewidth towards the dense ring. To
determine the associated turbulent pressure and turbulent energy in the ring,
it has to be taken into account that 13CO(3-2) is optically thick there which
leads to opacity broadening given by (Phillips et al., 1979)

$\frac{\Delta v}{\Delta
v_{int}}=\frac{1}{\sqrt{\text{ln}2}}\large[\text{ln}\large(\frac{\tau_{0}}{\text{ln}(\frac{2}{\text{exp}(-\tau_{0})+1})}\large)\large]^{0.5}$
(E1)

with $\Delta$vint the intrinsic velocity dispersion we are interested in,
$\Delta$v the observed velocity dispersion and $\tau_{0}$ the opacity. As
there is no optically thin tracer it is not straightforward to estimate the
opacity. Therefore, RADEX (van der Tak et al., 2007) was employed to estimate
the optical depth. This was done using the observed linewidth, a typical
column density N${}_{H_{2}}$ = 1.5$\times$1022 for the ring, [12CO]/[H2] =
10-4, [12CO]/[13CO] = 60, Tkin 15 K and n${}_{H_{2}}$ = 2.4$\times$104 cm-3.
This gives an estimated opacity of $\sim$ 2 which results in a correction
factor 1/1.35 for for the linewidth.

Figure 26: Histogram of the [C II] linewidths in the expanding shells in the
bipolar cavities. The dashed and dotted vertical lines indicate the mean and
median linewidth, respectively. Figure 27: Histogram of the 13CO(3-2)
linewidth over the full map of RCW 36 (blue). The green histogram indicates
the linewidth over the swept-up ring, pointing to an increased linewidth with
respect to the rest of the mapped region of Vela C.

##  Appendix F Spitzer and Weaver solutions

In order to estimate the expansion timescale of the different components in
RCW 36, we use the Spitzer (Spitzer, 1978) and Weaver (Weaver et al., 1977)
solution. First we will discuss the Weaver solution which is derived from Eqs.
51 and 52 in Weaver et al. (1977). This results in

$\text{t}_{exp}=\frac{16}{27}\frac{\text{R}}{\text{v}}$ (F1)

with R the radius of the expanding structure and v the observed velocity. For
the dense ring, we use R = 0.9 pc and v = 1.0-1.9 km s-1. For the shells in
the cavities, we use R = 1.0 pc and v = 5.2 km s-1.
The Spitzer solution for the expansion timescale is obtained from

$\text{R}=\text{R}_{St}\left(1+\frac{7}{4}\sqrt{\frac{4}{3}}\frac{\text{c}_{i}\text{t}_{exp}}{\text{R}_{St}}\right)^{4/7}$
(F2)

with ci the sound speed of the ionized gas (at a kinetic temperature Tk = 8000
K) and RSt the Strömgren radius which is given by

$\text{R}_{St}=\left(\frac{3\text{N}_{LyC}\text{m}_{p}^{2}}{4\pi\alpha_{B}\rho_{0}^{2}}\right)$
(F3)

with NLyC the number of photons emitted that will ionize the Strömgren radius,
mp the proton mass (= 1.67$\times$10-27 kg), $\alpha_{B}$ the recombination
coefficient (= 2.7$\times$10-13 cm3 s-1) and $\rho_{0}$ the density of the
neutral medium. For the number of ionizing photons we use 1048 s-1 inside the
ring and 8.5$\times$1046 s-1 in the cavities. This was derived based on the
radiocontinuum for the ring and the SHASSA H$\alpha$ flux for the cavities
(see above). Based on the current mass in the dense ring, we can assume a
density n${}_{H_{2}}$ = 2.4$\times$104 cm-3 for the neutral medium which gives
texp = 1.2 Myr. However, in Sec. 6.5 we show that potentially half of the mass
has already been removed from the central region, which would give a neutral
medium density of n${}_{H_{2}}\sim$5$\times$104 cm-3 and a resulting texp =
1.7 Myr. Considering that a significant amount of mass has been removed from
the region might thus increase the Spitzer expansion timescale. Even though we
have argued that the hot plasma from stellar winds might drive the expansion
in the cavities, we also calculated the resulting Spitzer timescale for this
regions. Using the above mentioned ionizing photon flux and a neutral medium
density between n${}_{H_{2}}$ = 102 cm-3 and 5$\times$102 cm-3 gives an
expansion timescale between 0.2 and 0.4 Myr. These timescales are also
consistent with the ones derived assuming a constant expansion velocity and
confirm that, even if expansion in the cavities is driven by ionizing
radiation, the expansion timescales of the cavities is short compared to the
lifetime of the H II region. All the obtained values for the expansion
timescales using different methods are summarized in Tab. 9.

Table 9: Obtained expansion timescales for the dense molecular ring and cavities assuming (1) a constant expansion velocity, (2) the Weaver et al. (1977) solution and (3) the Spitzer (1978) solution. region | constant velocity | Weaver et al. (1977) | Spitzer (1978)
---|---|---|---
dense ring | 0.5-0.9 Myr | 0.3-0.6 Myr | 1.2-1.7 Myr
cavities | 0.2 Myr | 0.1 Myr | 0.2-0.4 Myr