# Three supernova shells around an M33 star cluster

A. Camps-Fariña1,2, John E. Beckman1,2,3, J. Font1,2 and A. Borlaff1,2,J.
Zaragoza-Cardiel1,2,4 and P. Amram5
1Instituto de Astrofísica de Canarias, Tenerife, Spain, E38205
2Departamento de Astrofísica, Universidad de la Laguna, Tenerife, Spain,
E38205
3CSIC, Madrid, Spain, 28006
4Instituto de Astronomía, Universidad Nacional Autónoma de México, D.F. México
5Laboratoire d’Astrophysique de Marseille, Aix Marseille Université, CNRS,
F-13388, Marseille, France
<EMAIL_ADDRESS>
###### Abstract

Using a specialized technique sensitive to the presence of expanding ionized
gas we have detected a set of three concentric expanding shells in an HII
region in the nearby spiral galaxy M33. After mapping the kinematics in
H$\alpha$ with Fabry-Perot spectroscopy we used slit spectra to measure the
intensities of the [SII] doublet at $\lambda\lambda$ 671.9, 673.1 nm and the
[NII] doublet at $\lambda\lambda$ 645.8, 658.3 nm to corroborate the
kinematics and apply diagnostic tests using line ratios. These showed that the
expanding shells are shock dominated as would be the case if they had
originated with supernova explosions. Estimating their kinetic energies we
find fairly low values, indicating a fairly advanced stage of evolution. We
obtain density, mass and parent star mass estimates, which, along with the
kinetic energies, are inconsistent with the simplest models of shock-
interstellar medium interaction. We propose that the presence and properties
of an inhomogeneous medium offer a scenario which can account for these
observations, and discuss the implications. Comparing our results with data
from the literature supports the combined presence of an HII region and
supernova remnant material at the observed position.

###### keywords:

HII regions – ISM: bubbles – ISM: supernova remnants – galaxies: star
clusters: general

††pagerange: Three supernova shells around an M33 star
cluster–References††pubyear: 2002

##  1 Introduction

It is recognized that feedback from stellar winds and supernovae has had a
decisive effect on galaxy formation and evolution (Dekel & Silk, 1986). Galaxy
formation models without strong stellar feedback convert their gas into stars
too rapidly by two orders of magnitude (e.g. Hopkins, Quataert, & Murray
(2011)) compared to observations (Kennicutt, 1998), form too many stars by
factors up to $10^{3}$, (e.g. Faucher-Giguère, Kereš, & Ma (2011)) and cannot
explain the distribution of heavy elements in the intergalactic medium (e.g.
Wiersma et al. (2010)). Approximations to the effects of supernovae on the
interstellar medium are now a feature of simulation models (e.g. Marinacci,
Pakmor, & Springel (2014))which, as well as making ad hoc approximations to
ensure that the feedback is sufficiently strong to reproduce the properties of
observed galaxies, have traditionally assumed as simplifying limitations that
the surrounding ISM is homogeneous, and supernovae occur singly (Chevalier,
1974; Thornton et al., 1998). These limitations have been gradually relaxed,
as computing power has increased.

It is known that the real ISM is highly inhomogeneous (Osterbrock & Flather,
1959; McKee & Ostriker, 1977; Giammanco et al., 2004) and this has been used
in a recent set of models by Martizzi, Faucher-Giguère, & Quataert (2015). The
effect of multiple supernovae occurring within a stellar cluster was modelled
by Kim, Kim, & Ostriker (2011) and more recently by Hennebelle & Iffrig
(2014). Martizzi, Faucher-Giguère, & Quataert (2015) produced models with both
of these features. The aim of all the models is to reproduce the balance
between star formation enhancement and quenching due to the effects of the
supernovae on their surroundings, and thereby to predict the macroscopic
properties of galaxies. As is frequent in all issues related to star
formation, particularly massive star formation, a key difficulty for modellers
is the lack of observational constraints; in this article we report
observations which take a step forward in improving them.

##  2 Observations

We observed a field on the southern arm of M33 centered on the HII region
studied in this article during on 31st October 2014 using the Galaxy H$\alpha$
Fabry-Perot System (GH$\alpha$FaS, see Hernandez et al. (2008); Fathi et al.
(2008)), and followed this up using the Intermediate dispersion Spectrograph
and Imaging System (ISIS) on 25th December 2014. The region in question is
listed as BCLMP17-d, the central part of the BCLMP17 complex, in the 1974
Boulesteix HII region catalogue (Boulesteix et al., 1974; Hodge, Skelton, &
Ashizawa, 2002).

GH$\alpha$FaS is an integral field spectrometer mounted at the Nasmyth focus
of the 4.2m William Herschel Telescope (WHT) at the Observatorio del Roque de
los Muchachos (ORM, La Palma), producing a 3.4x3.4 arcmin data cube with
seeing-limited angular resolution (spaxel size $\sim 0.2$ arcsec) and a
spectrum at each point on the object, with 48 channels covering a spectral
range close to $\sim 400$ km/s, yielding a nominal velocity resolution of
$\sim$8 km/s for H$\alpha$. The data reduction procedures for GH$\alpha$FaS
have been well described in the literature. Here we give a very brief summary.
GH$\alpha$FaS does not use an optical derotator so a correction is applied by
software, as described in detail in Blasco-Herrera et al. (2010). Velocity
calibration, phase correction, sky removal and adaptive binning are performed
as described in Daigle et al. (2006) producing the final cube used here.

ISIS is a high-efficiency, double-armed, medium-resolution (8-120 Å/mm)
spectrograph,at the Cassegrain focus of the WHT, capable of providing up to 4’
slit length and 22” slit width. We observed the region with a 10” wide slit
using the red arm R1200R grating centred at 6800 Å. This allowed us to observe
simultaneously the H$\alpha$, SII and NII lines with a spectral resolution of
$\sim$12 km/s. We used a slit width of 10” to encompass in a single
observation the central part of the region where we detected the bubbles with
GH$\alpha$FaS.

We reduced the spectrum using IRAF. As we use only line flux ratios no flux
calibration was necessary.

##  3 Expansion maps

To search for expansive components we analysed the GH$\alpha$FaS data cube
with bubbly (Camps-Fariña et al., 2015), a program designed to find and fit
multi-component spectral line data cubes. The analysis is automatic, using a
mathematical property of Gaussians to detect multiple components, which are
then fitted (if they are above a threshold signal-to-noise) and checked to
look for pairs of components equidistant in velocity (within uncertainty
limits) from the central peak representing the bulk of the HII region. These
detections are then mapped in expansion velocity giving ”expansion maps”,
which we use to find the bubbles.

We used the default configuration in the BUBBLY parameter file with a minimum
S:N ratio of 5 to detect the expansion components. Because of the relatively
low signal-to-noise ratio threshold we added visual inspection of the
expansion maps to ensure that detection was spatially coherent. In Fig. 1 we
show the line profile of H$\alpha$ emission which indicates the presence of
the three shells, with a central main peak and three pairs of secondary
components. This profile comes from only a single resolved element of the
velocity map.

Figure 2 was produced by mapping all the pixels where a given pair of peaks
was detected, which displays the projected extent of the shell; the process
was repeated for each of the three shells. We can see that the smallest shell
has the most rapid expansion velocity, and the largest shell has the slowest,
as expected for shells originating in the centre and doing work on the
surrounding interstellar medium during their expansion.

Figure 1: Example of a spaxel from the GH$\alpha$FaS data cube which shows
three pairs of symmetrically placed secondary peaks, indicating expansion. The
detection and fitting was performed automatically using bubbly. Figure 2:
Expansion maps of the three detected bubbles, which show the detected
expansion velocity in each pixel, all in the same velocity scale. Overlaid are
contours of the region’s H$\alpha$ emission, it can be seen that the bubbles
are roughly concentric with each other and the region

##  4 Properties of the bubbles

We can measure the radius and expansion velocity directly on the expansion
maps and, using the H$\alpha$ flux of the shells, can calculate their
density(Relaño & Beckman, 2005):

$L_{H{\alpha}}(shell)=4\pi R^{2}\Delta R\,n_{e}^{2}\alpha(H,T)h\nu$ (1)

We cannot independently calculate the shell width $\Delta R$, so we took a
lower limit arising from the evolution of a strong shock disregarding
instabilities(Lozinskaya, 1992). A lower limit in shell thickness means a
lower limit in the total mass of the shell. The fact that we measure a lower
limit in mass will be important later in the discussion.

The flux is measured calculating the fraction of flux the pair of secondary
components emits and using a calibrated H$\alpha$ image of M33 (Hoopes,
Walterbos, & Bothun, 2001). We chose this method because we are more confident
in the calibration of the photometry of this image rather than calibrating
GH$\alpha$FaS.

We can also calculate mass and energy and using these three properties:

$\displaystyle{M=4\pi R^{2}\Delta R\>n_{e}m_{p}}$ ;
$\displaystyle{E_{k}=\frac{1}{2}M*v_{exp}^{2}}$ (2)

The age is estimated using the relation of $t=2/7*(R/v)$, from the analytic
expression for the evolution of a supernova remnant in the snowplough phase
(McKee & Ostriker, 1977). Using this expression is justified in the following
sections, where we also discuss the validity of the values calculated.

Lastly, we calculate the mean pre-shock density for each bubble assuming the
shells swept out all of the gas they encountered, dividing the total mass of
the shell by the total spherical bubble volume. We disregard the masses of the
SN ejecta (a few solar masses), as the shell masses have much larger values.

The ISIS spectrum was used to compute emission line ratios, to probe the
origin of the shells. To obtain results for each shell separately we applied a
modified version of bubbly to each spectrum over the region in the spatial
direction for each of these lines: H$\alpha$, [SII]6716, [SII]6731 and
[NII]6584. To compute the doublet [NII] we used the relation
$[\mathrm{NII}]_{6584}=3*[\mathrm{NII}]_{6548}$ which is independent of the
physical parameters of the gas.

To improve the statistical significance of the ratios we repeated the
detection using binning of 2, 3, 5, 6, 15 and 30 elements in the spatial
direction, and used the mean value with the standard deviation as the error
after eliminating clearly discrepant values.

We could in principle also use the [SII]6716/[SII]6731 ratio as an indicator
of density, but unfortunately the values measured are not within the range in
which the ratio is sufficiently sensitive to density. Table 1 lists all
parameters measured in the bubbles, both from GH$\alpha$FaS and ISIS data.

Table 1: Physical properties of the bubbles

| Bubble 1 | Bubble 2 | Bubble 3
---|---|---|---
Radius | 22 $\pm$ 2 | 16 $\pm$ 2 | 12.5 $\pm$ 1
(pc) |  |  |
$v_{exp}$ | 55 $\pm$ 5 | 115 $\pm$ 12 | 165 $\pm$ 17
(km/s) |  |  |
$L_{H\alpha}$ | 4.4 $\pm$ 0.4 | 1.4 $\pm$ 0.2 | 0.61 $\pm$ 0.06
($10^{36}$ erg/s) |  |  |
ne | 10.2 $\pm$ 1.1 | 9.3 $\pm$ 1.1 | 8.8 $\pm$ 1.0
(cm-3) |  |  |
Mass | 337 $\pm$ 34 | 118 $\pm$ 12 | 54 $\pm$ 5
(M⊙) |  |  |
Ek | 1.0 $\pm$ 0.2 | 1.6 $\pm$ 0.3 | 1.4 $\pm$ 0.3
($10^{49}$ erg) |  |  |
Age* | 114 $\pm$ 16 | 40 $\pm$ 6 | 21 $\pm$ 3
(kyr) |  |  |
n0* | 0.31 $\pm$ 0.06 | 0.28 $\pm$ 0.06 | 0.26 $\pm$ 0.05
(cm-3) |  |  |
$[\mathrm{SII}]/\mathrm{H}\alpha$ | 0.4 $\pm$ 0.15 | 0.45 $\pm$ 0.16 | 0.43 $\pm$ 0.09
$[\mathrm{NII}]/\mathrm{H}\alpha$ | 0.44 $\pm$ 0.2 | 0.35 $\pm$ 0.09 | 0.27 $\pm$ 0.13
$[\mathrm{SII}]_{1}/[\mathrm{SII}]_{2}$ | 1.38 $\pm$ 0.25 | 1.41 $\pm$ 0.18 | 1.42 $\pm$ 0.26

Notes. See section 4 for details on how the properties were determined. ne is
the electron density inside the shell, while n0 is the average ambient
particle density it encountered. The properties marked with * (Age and ambient
density) are only valid for a shock evolving in a homogeneous medium. The
calculations use a lower limit for shell thickness, this implies that mass and
kinetic energy are lower limits while density is an upper limit. Even then,
the dependence on shell thickness is $\sqrt{\Delta R}$, so a factor 10
increase in the thickness would increase mass and energy a factor $\sim$ 3.

###  4.1 Archival properties of the cluster

We have supporting information from the literature obtained by querying the
position with VizieR; the presence of a supernova remnant in this HII region
detected using radio observations at lower angular resolution is reported by
Gordon et al. (Gordon et al., 1999) (Tables 1,2 #90). Also relevant is a
survey of HII regions and molecular gas by Miura et al. (2012) which used
archival Hubble Legacy Archive images to date the accompanying star cluster by
fitting isochrones to a color-magnitude diagram extracted using aperture
photometry (see section 3.2.2). This gives an age range of 6-12 Myr for the
cluster, which is also reported to have 51 O stars (Table 3 #27). This survey
also shows a $2.2\cdot 10^{6}$ M⊙ molecular cloud adjacent to the cluster
(Table 2 #24).

###  4.2 Identification as supernova remnants

The identification as supernova remnants was made using forbidden line ratios
[SII]/H$\alpha$ and [NII]/H$\alpha$ to place our bubbles within diagnostic
plots which distinguish between supernova remnants, HII regions and planetary
nebulae (Sabbadin, Minello, & Bianchini, 1977; Garcia Lario et al., 1991). The
classic diagnostic of [SII]/H$\alpha$ $>$ 0.4 is included in these plots as
the region for SNR is delimited by this value on the X axis (see Fig. 3). All
three fall well within the defined regions for SNR, though for bubbles 1 and 2
the error reaches the area for photoionized gas. Significantly, the same
analysis for the central peak places it, as expected, within the area for HII
regions.

Though the ratios indicate otherwise, an alternative origin for the bubbles
might be stellar winds, but this is ruled out in practice; as winds are
continuously injecting energy since the stars’ birth they cannot produce the
segregated concentric structure observed, as the individual bubbles would
blend and merge into one structure before growing to their current size, which
already encompasses most of the cluster. The only possibility for winds to
produce bubbles afterwards would be a red supergiant wind accelerated by a
subsequent Wolf-Rayet phase (see Freyer, Hensler, & Yorke (2006)), but the
stars capable of producing this effect have masses over 25 M⊙ and could no
longer be present, barring multiple star formation bursts.

Figure 3: Diagnostic diagram used to distinguish between supernova remnants
(SNR), HII regions (HII) and planetary nebulae (PNe) based on the relation
between the SII to H$\alpha$ and NII to H$\alpha$ line ratios. We plot the
remnants with their identifier numbers as well as the HII region.

###  4.3 Physical properties

The evolution of supernova remnants is divided into three main phases, an
initial free expansion while the mass of the ejecta dominates, an adiabatic
expansion once the swept-up material mass is similar to that of the ejecta and
a radiative phase, also called snowplough phase, in which the losses by
radiation become significant and the gas cools. Given that the initial energy
of a supernova is around $10^{51}$ erg our remnants currently have lost some
99% of their energy, which means that they must have been radiative for some
time. We can check this using the expressions for the adiabatic phase (Hnatyk,
Petruk, & Telezhyns’Kyi, 2007):

$\displaystyle{R_{ST}(t)=\left(\frac{25E_{SN}}{4\pi\rho_{0}}\right)^{1/5}}$ ;
$\displaystyle{t_{tr}=2.84\cdot 10^{4}E_{51}^{4/17}n_{0}^{-9/17}}$ (3)

where the left expression is the evolution of the radius and on the right we
have the transition time to radiative phase. For our measured ambient density
we find a transition time of 54 kyr which already does not fit the data, two
of our remnants should not have transitioned yet. Furthermore, even with the
faster expansion of the adiabatic phase we find predicted sizes roughly double
the measured radii. We will address this apparent inconsistency again shortly.

The mass range for the progenitor stars, calculated from the age of the
cluster, is from $\sim$15 to 19 M⊙, which is not unreasonable if we find three
similar mass stars in a cluster of this size. Given the relatively short
intervals between the supernova explosions, their progenitors must have had
very similar masses.

The swept-up masses of the shells are very interesting. Their values are not
consistent with the simplest models but they give us a clue to solve the
inconsistency in the physical parameters. In a homogeneous medium the
expanding shock should sweep out almost all the surrounding gas into the shell
leaving a very rarefied medium, so that a second supernova soon after the
first should sweep up hardly any material, lengthening the free-expansion
phase. In spite of this, we find that each consecutive remnant has accumulated
tens of solar masses of material. This is so even though our value for the
thickness of each shell is a lower limit, meaning the masses are also lower
limits as they scale with the square root of the thickness. A mechanism is
needed to replenish gas inside a remnant for the next shell to encounter
sufficient material; the most plausible is mass-loading from cold dense
clouds. It has long been suspected that the interior of an HII region is not
homogeneous nor fully ionized, and that small molecular clouds can survive for
a long time inside it as they slowly evaporate (Bok & Reilly, 1947).

Hydrodynamic ablation by shocks or supersonic winds can also disrupt such
clouds, so each supernova shell feeds gas to the next one. The effects of the
interaction on the shell are also important. Simulations show (Chevalier,
1974) that upon encountering a cold, dense cloud, the shell will tend towards
radiative evolution, rapidly slowing down and losing energy, greatly advancing
its stage of evolution. In the cited paper a simple model with a single
density jump at a fixed radius is used, while we expect clusters of cloudlets
instead. Either way, the effect of a density jump on the shell is what we have
observed in our remnants, so adding the condition of inhomogeneity to our
model explains the apparent inconsistencies. It also means that the ages
estimated assuming a homogeneous medium are not correct. An inhomogeneous
medium slows down the remnants, so our values are upper limits, but they
should be within a small factor of the real expansion ages.

##  5 Discussion

We have presented results which seem to indicate the presence of three
currently expanding supernova remnants within a single young HII region. Our
key claim is that the expansion detected must arise from three different
objects. Initially, we considered other geometrical configurations, such as
typical complex structures in HII regions with several arches and cavities,
but we detect very clearly an approaching and receding component for all three
along the same lines of sight, so they are either nested or aligned on the
line of sight. Given the similar size of the bubbles relative to the cluster,
and their concentricity, we believe that nested shells offer the simpler
explanation.

We found that the physical properties appeared inconsistent given their
masses, kinetic energies and ages. This is solved considering an inhomogeneous
medium with embedded dense clumps of molecular gas, which would rob the
remnants of energy as well as refill the ambient gas to be swept up by the
subsequent supernova explosions. Because of this, our interpretation requires
such clumps to survive for a long time inside the HII region.

The survival of clumps inside HII regions has long since been documented, (Bok
& Reilly, 1947), and they have been found even in extreme examples such as 30
Doradus (Indebetouw et al., 2013), where a complex of molecular clouds is
found associated with the brightest star cluster.

The effect of stellar winds and radiation on their parent molecular clouds has
been studied in simulations by Rogers & Pittard (2013) and Dale et al. (2014).
In the former, three massive $\sim$30 M⊙ stars affect a surrounding 4 pc
massive molecular cloud. The main finding is that the winds carve low density
channels through which the gas escapes, making the denser regions surprisingly
resistant to ablation. Another important result is that dense clumps are
largely impervious to supernova shocks. They find that the dense material is
either destroyed or pushed outside the original 4 pc radius in just over 6
Myr, which coincides with the lower limit of the age for our HII region. Dale
et al. (2014) perform cluster-scale simulations and find results compatible to
those of Rogers & Pittard (2013) with a very structured dense medium. They use
molecular clouds with a variety of physical parameters finding that the
denser, more massive clouds are much harder to disrupt. For lower mass clouds
the stars carve irregular cavities on scales of 10 pc. The molecular cloud
associated with our cluster is similar to the densest and most massive of Dale
et al. (2014)’s.

Both studies show that the presence of embedded clumps is expected in clusters
born from massive molecular clouds, and that they can easily survive until the
stars explode as supernovae. We also need to estimate the rate of cloud
evaporation required to refill the gas and reproduce the masses in the
observed shell. From the masses and ages of the remnants we require an
injection rate of about 25 M⊙ per 10 kyr. Dale et al. (2014) find that except
at the onset of star formation, winds hardly affect the disruption of the
molecular cloud; with photoevaporation dominating this process, so here we
need consider only cloud destruction by photoevaporation.

Using the relation from Pittard (2007) for photoevaporation and taking a
sample cloud with radius 1 pc at a distance of 10 pc from the centre of the
cluster, we estimate the flux of ionizing photons from the cluster using its
H$\alpha$ luminosity and obtain an injection rate of 0.8 M⊙ per 10 kyr. We can
therefore easily reproduce the required mass injection without an overly large
number of clouds, even accounting for the fact that our masses are lower
limits.

Feedback from stellar winds and supernovae plays an important role in many
processes relevant to galaxy evolution. As well as those mentioned above there
is the possible dissipation of nuclear dark matter cusps (Pontzen & Governato,
2012), the dissipation of metals both within galaxies (Spitoni et al., 2009)
and from galaxies into the intergalactic medium (Heckman, Armus, & Miley,
1990), and the enhancement of the infall rate of low metallicity intergalactic
gas to the disc, by its interaction with supernova ejectae in the galactic
halo (Marasco, Fraternali, & Binney, 2012). To incorporate any of these
effects into models of the formation and evolution of galaxies, it is clearly
important to have not only model inputs to the dynamics of the interaction of
the winds and supernovae with the surrounding ISM, but measurements which can
help quantify the process. As the stars in massive star clusters form
virtually simultaneously, cases of multiple supernova explosions must be
common enough to be incorporated into the relevant models.But much more
generally to make realistic models mass loading of the expansion via the
dissipation of dense cool clumps must be taken into consideration. The triple
supernova has allowed us to illustrate its importance, and make quantitative
estimates of relevant parameters for much wider application.

## Acknowledgements

We thank M. Beasley, J.E. Dale, and the anonymous referee for useful
discussion and comments. This research has been supported by the Instituto de
Astrofísica de Canarias under project P/308603. JEB acknowledges financial
support to the DAGAL network from the People Programme (Marie Curie Actions)
of the European Union’s Seventh Framework Programme FP7/2007-2013/ under REA
grant agreement number PITN-GA-2011-289313.

Based on observations made with the William Herschel Telescope under
Director’s Discretionary Time of Spain’s Instituto de Astrofísica de Canarias.

The William Herschel Telescope is operated on the island of La Palma by the
Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of
the Instituto de Astrofísica de Canarias.

This research has made use of the VizieR catalogue access tool, CDS,
Strasbourg, France. The original description of the VizieR service was
published in A&AS 143, 23.

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