Star Formation in Ram Pressure Stripped Tails

Abstract

We investigate the impact of star formation and feedback on ram pressure
stripping using high-resolution adaptive mesh simulations, building on a
previous series of papers that systematically investigated stripping using a
realistic model for the interstellar medium, but without star formation. We
find that star formation does not significantly affect the rate at which
stripping occurs, and only has a slight impact on the density and temperature
distribution of the stripped gas, indicating that our previous (gas-only)
results are unaffected. For our chosen (moderate) ram pressure strength,
stripping acts to truncate star formation in the disk over a few hundred
million years, and does not lead to a burst of star formation. Star formation
in the bulge is slightly enhanced, but the resulting change in the
bulge-to-disk ratio is insignificant. We find that stars do form in the tail,
primarily from gas that is ablated from the disk and the cools and condenses in
the turbulent wake. The star formation rate in the tail is low, and any
contribution to the intracluster light is likely to be very small. We argue
that star formation in the tail depends primarily on the pressure in the
intracluster medium, rather than the ram pressure strength. Finally, we compare
to observations of star formation in stripped tails, finding that many of the
discrepancies between our simulation and observed wakes can be accounted for by
different intracluster medium pressures.

Full text

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arXiv:1203.0308v1 [astro-ph.CO] 1 Mar 2012
Mon. Not. R. Astron. Soc. 000, 1–?? (2011) Printed 5 March 2012(MN LATEX style file v2.2)
Star Formation in Ram Pressure Stripped Galactic Tails
Stephanie Tonnesen1⋆and Greg L. Bryan2
1Department of Astrophysics, Princeton University, Peyton Hall, Princeton, NJ 08544, USA
2Department of Astronomy, Columbia University, Pupin Physics Laboratories, New York, NY 10027, USA
5 March 2012
ABSTRACT
We investigate the impact of star formation and feedback on ram pressure stripping us-
ing high-resolution adaptive mesh simulations, building on a previous series of papers
that systematically investigated stripping using a realistic model for the interstellar
medium, but without star formation. We find that star formation does not signifi-
cantly affect the rate at which stripping occurs, and only has a slight impact on the
density and temperature distribution of the stripped gas, indicating that our previous
(gas-only) results are unaffected. For our chosen (moderate) ram pressure strength,
stripping acts to truncate star formation in the disk over a few hundred million years,
and does not lead to a burst of star formation. Star formation in the bulge is slightly
enhanced, but the resulting change in the bulge-to-disk ratio is insignificant. We find
that stars do form in the tail, primarily from gas that is ablated from the disk and
the cools and condenses in the turbulent wake. The star formation rate in the tail is
low, and any contribution to the intracluster light is likely to be very small. We argue
that star formation in the tail depends primarily on the pressure in the intracluster
medium, rather than the ram pressure strength. Finally, we compare to observations
of star formation in stripped tails, finding that many of the discrepancies between our
simulation and observed wakes can be accounted for by different intracluster medium
pressures.
Key words: galaxies: clusters, galaxies: interactions, methods: N-body simulations
1 INTRODUCTION
As galaxies orbit within a cluster, their interstellar medium
(ISM) may interact directly with the intracluster medium
(ICM), the hot halo of gas bound by the cluster gravitational
potential. This is often thought of as a gas-only affair, in
which stars remain unaffected. For example, ram pressure
stripping (and related processes) by the ICM is only able to
remove a galaxy’s gas (Gunn & Gott 1972). Although stars
are not directly affected, there is increasing evidence that
ISM-ICM interactions do impact star formation.
In general, galaxies in clusters have lower star formation
rates than galaxies of the same morphological type in the
field (e.g. Hashimoto et al. 1998; Rines et al. 2005; Balogh et
al. 1998; although Gom´ ez et al. (2003) found that a galaxy’s
star formation rate depended more on local galaxy density
than on cluster membership). Gavazzi et al. (2006) found
that the star formation rate in cluster galaxies was related
to the amount of H I: galaxies with normal H I had twice
the Hα equivalent widths of H I deficient galaxies. Koop-
mann & Kenney (2004) found that Virgo spirals are form-
⋆E-mail:
gbryan@astro.columbia.edu (GLB)
stonnes@astro.princeton.edu(ST));
ing stars primarily in the centres of their disks, which can
be explained by the outer gas disk having been stripped by
the ICM. In a more recent study of ten Virgo spiral galaxies,
Crowl & Kenney (2008) find that the star formation history
and quenching time of five of their galaxies is consistent with
ram pressure stripping in the cluster centre, while the other
five have more complicated histories.
Although there is evidence linking ISM-ICM interac-
tions to star-formation quenching, the question of whether
interactions with the ICM could also induce star formation
in a galactic disk remains open. For example, star formation
could be triggered by the increase in surrounding pressure
when a galaxy enters a high-density ICM (Dressler & Gunn
1983; Evrard 1991; Fujita 1998; Smith et al. 2010). Fujita &
Nagashima (1999) found ram pressure induced star forma-
tion in simulated galaxies.
Observations of post-starburst galaxies in clusters in-
dicate that star formation can be induced by environmen-
tal processes (Dressler & Gunn 1983), although the exact
mechanism is still unknown. Post-starburst galaxies reside
preferentially in clusters at z = 0.3 − 0.6 (e.g. Poggianti
et al. 1999, Poggianti et al. 2004; Tran et al. 2004; but
see Balogh & Bower 2003). Although most observations at
both lower and higher redshift have found k+a fractions in-
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2S. Tonnesen and G. L. Bryan
creasing with decreasing galaxy density (e.g. Zabludoff et al.
1996; Hogg et al. 2006; Goto 2007; Yang et al. 2008), Pog-
gianti et al. (2004) find a large population of low-luminosity
post-starburst galaxies in Coma. Tran et al. (2003) exam-
ined E+A galaxies in three intermediate-redshift clusters
and found that the majority were not associated with merg-
ers. By considering the spatial distribution of post-starburst
galaxies in clusters, some observers have found evidence
that both the star formation quenching and earlier starburst
could be related to interactions with the ICM (Poggianti et
al. 2004; Poggianti et al. 2009; Ma & Ebeling 2008).
By investigating whether and how ram pressure strip-
ping by a dense ICM can affect star formation in a galaxy, we
can shed light on what drives the evolution of cluster galax-
ies. From z∼0.5 to z=0, much of the morphological evolution
in clusters has been from spirals to S0s (Dressler et al. 1997).
In order for a spiral galaxy to evolve into an S0, both spec-
troscopic and morphological changes must take place: the
galaxy must become red, it must lose spiral structure in its
disk, and the bulge-to-disk ratio must increase.
Ram pressure stripping, through gas removal, can cause
a galaxy to become red. Passive spirals, red spirals without
star formation, have been observed to reside preferentially in
clusters (Moran et al. 2007; Poggianti et al. 1999). A galaxy-
ICM interaction is the likely mechanism for forming passive
spirals because these galaxies have had their gas removed
but are morphologically spirals (because their stellar disks
are relatively undisturbed). If passive spirals are precursors
to S0s, then a galaxy-ICM interaction is a step in the evo-
lution of normal spirals into S0s.
Ram pressure stripping can also result in a disk galaxy
losing its spiral arms. Bekki et al. (2002) used simulations
to show that if gas is no longer accreted by a galaxy, it
will lose its spiral arms in about 3.5 Gyr. This is in good
agreement with observational and model estimates for how
long morphological transformation may take (e.g Kodama
& Smail 2001; Poggianti et al. 1999; but see Moran et al.
2007 for a shorter estimate).
Finally, a galaxy’s bulge-to-disk (B/D) ratio can in-
crease either by fading the disk or growing the bulge. Ram
pressure stripping can result in the disk of a galaxy fading.
Fading the disk of a galaxy with a B/D = 0.2 will result in a
galaxy with a B/D = 0.5, so an Sb galaxy, after disk fading,
will have the B/D ratio of an S0 galaxy (Fujita & Nagashima
1999; Solanes et al. 1989). Solanes et al. (1989) find that the
bulge luminosities of Sa galaxies are similar to those of S0s,
and all galaxy types have a range of luminosities, so it is
not universally necessary to increase the bulge luminosity
to transform a spiral to an S0. However, if the slow fading
of the disk was the only mechanism at work, the total lu-
minosity of S0s should be less than that of spirals, which
is not generally the case (Burstein et al. 2005). Christlein
& Zabludoff (2004) compared the bulge and disk luminosi-
ties of cluster galaxies, and concluded that S0s in clusters
form by growing the bulge of spiral galaxies. Therefore it
is important to know whether ram pressure can induce star
formation that may grow a galaxy’s bulge.
There can also be star formation in a ram pressure
stripped tail of gas, even though molecular clouds are con-
sidered too dense to be directly stripped from a galactic disk.
It is still unclear how common star formation in stripped gas
tails is, as many tails observed in H I do not have any as-
sociated star formation. For example, several one-sided H I
tails have been observed in Virgo. Chung et al. (2007) found
7 one-sided tails between 0.6-1 Mpc in projected distance
from M87. Oosterloo & van Gorkom (2005) found a ∼110
kpc H I tail associated with NGC 4388. NGC 4438 was orig-
inally believed to be an interacting galaxy (e.g. Hibbard &
van Gorkom 1990; Kenney et al. 1995; Vollmer et al. 2005),
but more recent work has concluded that it is likely only
ram pressure stripped. NGC 4522 is 1 Mpc from M87, and
Kenney et al. (2004) conclude that it is being ram pressure
stripped by an overdense or moving region of the ICM. Fi-
nally, NGC 4402 also has an H I tail (Crowl et al. 2005).
Of these eleven galaxies with clear stripping signatures in
H I observations, only four have been found to have star
formation either from UV emission or H II regions.
The four H I tails with star formation tend to have short
stellar tails with low star formation rates. Abramson et al.
(2011) find 9 UV emitting regions near NGC 4330, for a total
of 4.55 × 106M⊙ of extragalactic stars. Cortese et al. (2003;
2004) found an H II region 3 kpc above the disk of NGC 4402
in the Virgo cluster. Similarly, NGC 4522, and NGC 4438
have ongoing star formation in their stripped tails close to
the galaxy (Kenney & Koopmann 1999; Boselli et al. 2005).
NGC 4438 has the most distant UV emission out to nearly
30 kpc from the disk (Boselli et al. 2005). Although Gerhard
et al. (2002) find an H II region near NGC 4388 in the Virgo
cluster, this H II region is not near the long H I tail observed
by Oosterloo & van Gorkom (2005), and there have been no
stars found associated with the stripped H I.
There have been an increasing number of recent obser-
vations of star formation in stripped gas tails. Long trails
of star-forming knots were observed in two massive galaxy
clusters by Cortese et al. (2007), extending as far as 80 kpc
from one of the galaxies. Cortese et al. (2007) find that a
combination of tidal and ram pressure stripping are affect-
ing the galaxies. In the Coma cluster, Yoshida et al. (2008)
found a complex of Hα filaments and clouds extending up
to 80 kpc from the E+A galaxy RB 199. They also con-
clude that the most likely gas removal scenario involves a
combination of a merger and ram pressure stripping. Yagi
et al. (2010) find another 13 galaxies with young stars or Hα
clouds in tails. Finally, ES0 137-001, which we discussed in
detail in Tonnesen et al. (2011), has 35 HII regions extend-
ing more than 30 kpc from the disk. These HII regions are
spatially correlated with a ram pressure stripped tail of gas
(Sun et al. 2006, 2007, 2010). Hester et al. (2010) found Hα
emission indicative of star formation in the UV tail of IC
3418 (Martin et al. 2005), a low surface brightness galaxy in
the Virgo Cluster.
In a series of 12 SPH simulations including ram pressure
stripping and star formation, Kapferer et al. (2009) found
that stars can form in their ram pressure-stripped tail out
to hundreds of kiloparsecs from the disk. In fact, using fast,
high-density ICM winds, the authors found more stars form-
ing in the stripped tail than in the remaining disk. The star
formation rate increased due to enhanced external pressure
provided by the ICM. They also found that stars formed in
the tail could fall back into the stellar bulge.
Clearly, the standard lore that galaxy-ICM interactions
do not result in stars outside of the galaxy has been over-
turned. In fact, if star formation in stripped tails is common,
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Star Formation in Ram Pressure Stripped Tails3
ram pressure stripping could contribute a significant fraction
of the Intracluster Light (ICL).
In this paper, we run a set of high resolution simulations
(about 40 pc resolution, which is small enough to marginally
resolve giant molecular clouds) to understand whether ram
pressure can induce star formation in a galactic disk, pro-
ducing starburst galaxies or increasing the mass of the bulge,
and whether stars can form in a stripped gas tail and add
to the ICL.
The paper is structured as follows. After a brief intro-
duction to our methodology, we provide the general charac-
teristics of our simulation (§2.1-2). We introduce the param-
eters of our specific simulations in §2.3. In §2.4 we discuss
how we make projections of observables. We then discuss
our results (§3), first focusing on how star formation affects
the disk and then the stripped tail. In §4 we compare our
results to observations and previous simulations. We discuss
the possible effects of our resolution in §5. Finally, we con-
clude in §6 with a summary of our results and predictions
for observers.
2 METHODOLOGY
We use the adaptive mesh refinement (AMR) code Enzo.
To follow the gas, we employ an adaptive mesh for solving
the fluid equations including gravity (Bryan 1999; Norman
& Bryan 1999; O’Shea et al. 2004). The code begins with a
fixed set of static grids and automatically adds refined grids
as required in order to resolve important features in the flow.
Our simulated region is 311 kpc on a side with a root
grid resolution of 1283cells. We allow an additional 6 levels
of refinement, for a smallest cell size of 38 pc. We refine the
grid based on the local gas mass, such that a cell was flagged
for refinement whenever it contained more than about 4900
M⊙. We found that these parameters quickly refined most
of the galactic disk to 38 pc resolution. The run also refined
much of the wake to a spatial resolution of about 76 pc, and
the dense clouds to 38 pc.
The simulation includes radiative cooling using the
Sarazin & White (1987) cooling curve extended to low tem-
peratures as described in Tasker & Bryan (2006). To mimic
effects that we do not model directly (such as turbulence
on scales below the grid scale, UV heating, magnetic field
support, or cosmic rays), we cut off the cooling curve at a
minimum temperature Tmin so that the cooling rate is zero
below this temperature. In the simulations described here we
use Tmin = 300 K, below the threshold for neutral hydro-
gen formation. In previous work (Tonnesen & Bryan 2009;
2010), we have explored the impact of adopting a Tmin value
of 8000 K, finding the effect to be relatively small.
2.1 Star Formation and Feedback Implementation
Star formation occurred in our finest grid cells (38 pc) when
two criteria were met: 1) the gas density in a cell exceeded
a critical overdensity (in our runs, this was set to a density
of about 3.85 × 10−25g cm−3), and 2) the gas temperature
was below 1.1 × 104K.
The reader may note that this set of criteria is missing
two commonly used requirements for star formation (e.g.
Cen & Ostriker 1992): (1) there is a convergent flow, and
(2) the mass in the cell exceeds the Jean’s mass. We chose
not to require a convergent flow because we intend to look
for star formation in the stripped gas tail and may not be
able to resolve the internal structure of the clouds in the tail
accurately. In a previous paper (Tonnesen & Bryan 2010),
we have shown that dense gas can be accelerated to nearly
1000 km s−1. In addition, we have dropped the requirement
that the mass in the cell must exceed the Jeans mass because
with this condition, our minimum temperature floor could
prevent star formation except in the densest cells. Using our
minimum allowed gas density and maximum temperature
for star formation, the Jeans mass is 300 times the mass in
a cell. We fulfill the Truelove criterion (1997) using those
parameters, but will discuss the limits of our resolution in
more detail in Section 4.4.1.
The implementation of star formation and feedback is
explained in detail in Tasker & Bryan (2006), and our sum-
mary here directly reflects that paper (but is included for
completeness). When the gas in a cell meets our require-
ments for star formation, some of the gas is turned into a
star particle. The mass of the star particle is:
m∗ = ǫ∆t
tdynρgas∆x3
(1)
in which ǫ is the star formation efficiency, ∆t is the timestep,
tdynis the local dynamical time, ρgas is the gas density and
∆x is the cell size. The efficiency parameter was chosen to
match the Kennicutt-Schmidt relation (as in Tasker & Bryan
2006). The star formation and feedback parameters we use
are given in Table 1.
If the above requirements are met and the resulting star
particle will have a mass above a minimum mass, m∗min, it is
formed. This mass is chosen so that a large number of small
star particles will not slow down the simulation. However, if
the star particle would have a smaller mass, the probability
that it will form is equal to the ratio of the mass of the
projected star particle to m∗min. If the star particle is then
formed, its mass is the minimum of m∗min and 80% of the
mass in the gas cell. Thus, the probability of forming stars in
any individual cell is low, but this algorithm still produces
stars at the specified rate. This is done by keeping track
of the amount of mass in cells that fulfilled all of the star
formation criteria except the minimum mass requirement.
When the mass of the unformed stars reaches the minimum
mass, star particles will form even if they are below the
specified minimum mass. In order to make sure that we were
not missing any stars formed in the larger volume of the
stripped tail, we allowed for a smaller minimum mass in our
simulation with ram pressure stripping. As we will show,
this lower m∗min does not change the star formation in the
disk, but we did find that the lower minimum was necessary
to form the correct amount of stars in the tail.
We also include stellar feedback from Type II supernova
explosions. Not only may this be important for regulating
star formation in the galactic disk (e.g. Tasker & Bryan 2006;
Robertson et al. 2004), but feedback from star formation in
a stripped gas tail could enhance the rate at which stripped
gas mixes with the ICM. Star formation in a molecular cloud
is likely to be spread out over a dynamical time, and so in
order to calculate the timeline of feedback, stars in a particle
are considered to form according to the relation:
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4 S. Tonnesen and G. L. Bryan
Variable Value
ρmin
Tmax
3.85 × 10−25g cm−3
1.1 × 104K
0.5%
104M⊙ (SFNW) 102M⊙ (SFW)
10−5
ǫ
m∗min
ǫSN
Table 1. Star Formation and Feedback Parameters
mstar(t) = m∗
?t
tSF
(t − tSF)
τ2
exp−(t − tSF)
τ
dt,(2)
where tSF is the time at which the star particle was formed
and τ = max(tdyn, 10 Myr). As in Tasker & Bryan (2006),
over a time period of a few τ, 10−5of the rest-mass energy of
the stars is added to the thermal energy of the gas in the cell
in which the star particle has been created. This corresponds
to approximately 56 solar masses of stars formed for each
1051erg SN.
2.2 The Galaxy
Our galaxy is placed at a position corresponding to
(155.5,155.5,68.42) kpc from the corner of our cubical 311
kpc computational volume, so that we can follow the
stripped gas for more than 200 kpc. The galaxy remains
stationary throughout the runs. The ICM wind flows along
the z-axis in the positive direction, with the lower z bound-
ary set for inflow and upper z boundary set as outflow. The
x and y boundaries are set to outflow in all three cases.
We chose to model a massive spiral galaxy with a flat
rotation curve of 200 km s−1. It consists of a gas disk that is
followed using the adaptive mesh refinement algorithm (in-
cluding self-gravity of the gas and any newly formed stars),
as well as the static potentials of the (pre-existing) stellar
disk, stellar bulge, and dark matter halo. We directly follow
Roediger & Br¨ uggen (2006) in our modeling of the stellar
and dark matter potential and gas disk. In particular, we
model the stellar disk using a Plummer-Kuzmin disk (see
Miyamoto & Nagai 1975), the stellar bulge using a spherical
Hernquist profile (Hernquist 1993), and the dark matter halo
using the spherical model of Burkert (1995). This dark mat-
ter halo model is compatible with observed rotation curves
(Burkert 1995; Trachternach et al. 2008). The equation for
the analytic potential is in Mori & Burkert (2000). We de-
scribe our disk in detail in Tonnesen & Bryan (2009, 2010).
Briefly, our stellar disk has a radial scale length of 4 kpc,
a vertical scale length of 0.25 kpc and a total mass of 1011
M⊙; the stellar bulge has a scale length of 0.4 kpc and a
total mass of 1010M⊙; and the dark matter halo has a scale
radius of 23 kpc and a central density of 3.8×10−25g cm−3.
The gas disk has about 10% of the mass in the stellar disk,
and radial and vertical scales of 7 kpc and 0.4 kpc, respec-
tively.
To identify gas that has been stripped from the galaxy
we also follow a passive tracer that is initially set to 1.0 in-
side the galaxy and 10−10outside. In the following analysis,
we will use a minimum tracer fraction of 25% to find gas
stripped from the galaxy (our conclusions do not change if
we use 10% instead).
2.3 The Simulations
All three of the galaxies we discuss in this paper initially
evolve in a static, high-pressure medium with ρ = 9.152 ×
10−29g cm−3and T = 4.15 × 106K, to allow cool, dense gas
to form in the galaxy. This naturally generates a multiphase
ISM (see Tasker & Bryan (2006) and Tonnesen & Bryan
(2009) for more discussion of the ISM properties).
In our simulation without star formation, after 155
Myrs we reset the boundary conditions to generate a con-
stant ICM inflow along the inner z-axis, which is always
face-on to the galaxy. We chose this time so that the galaxy
would have formed high density gas clouds (ρ > 10−23g
cm−3) by the time the wind hits the galaxy (190 Myr after
the start of the simulation). In our comparison case with star
formation, we delay the onset of the wind by about 18 Myr
in order to allow the galaxy to evolve for more than 200 Myr
before the wind hits the disk. We do this for two reasons:
first, because Tasker & Bryan (2006) found that the star for-
mation rate in the disk settles to a relatively constant value
after about 200 Myr of star formation, and second, because
at that point in our simulations the azimuthally-averaged
relation between gas surface density and SFR surface den-
sity agrees with that found in Kennicutt (1989, 1998). While
we could have chosen to wait longer, this change would have
no qualitative effect on our conclusions.
In this paper we discuss three simulations. All of these
runs have the same initial conditions (same galaxy density
profiles evolving in a static ICM with ρ = 9.152 × 10−29g
cm−3and T = 4.15 × 106K and a cooling curve following
Sarazin & White (1987) extended to Tmin = 300 K). Two of
these simulations include star formation, SFNW and SFW.
SFNW evolves in a static ICM, while SFW initially evolves
in a static ICM and is later stripped by a higher-density
ICM wind. The final simulation we discuss in this paper is
NSFW, which is the same simulation as SFW without star
formation. NSFW is the Tmin = 300 K case discussed in our
earlier paper, Tonnesen & Bryan (2010). In both wind cases,
Pram = ρv2
km s−1. The ICM wind has a T = 4 × 107K and ρ = 3.2
× 10−28g cm−3.
ICM= 6.4× 10−12dynes cm−2, and vICM = 1413
2.4Projections
Enzo outputs the density and temperature of the gas in
each cell. To transform these values into H I column den-
sity and Hα intensity, we used Cloudy, version 08.00 of the
code, last described by Ferland et al. (1998). Using a grid
of temperatures and densities, we calculated the hydrogen
neutral fraction and Hα emissivity. In the Cloudy calcula-
tion, we included cosmic microwave background radiation,
the cosmic ray background, bremsstrahlung radiation from
the ICM and the 2005 version of the Haardt & Madau
(2001) z = 0 metagalactic continuum, as implemented by
Cloudy. We found that including a local interstellar radia-
tion field emission resulted in somewhat lower amounts of
neutral gas. Since much of our gas is very distant from the
galaxy, we decided not to include this radiation. We also
found that removing bremsstrahlung radiation did not sig-
nificantly change any of the values we considered.
We chose to calculate the neutral fraction and Hα emis-
sivity (from collisionally heated diffuse gas; emission from
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Star Formation in Ram Pressure Stripped Tails5
HII regions will be dealt with separately) for a thin plane-
parallel gas cloud of width 100 pc. We selected this width
because it corresponds roughly to the cell size of most of
the gas in the highly resolved tails, and accounts approxi-
mately for radiative transfer effects. If we assumed the ra-
diative thin limit, it would slightly decrease the amount of
H I we predict, and slightly increase the Hα emission for
dense, low-temperature gas. Ideally, we would include the
radiation field with radiative transfer directly in the sim-
ulation, but this is not yet feasible (and only has a slight
impact on the dynamics); instead we post-process these re-
sults to get reasonable predictions for the ionization fraction
and Hα emissivity (see Furlanetto et al. 2005 for a discussion
of various approaches in the context of Lyα emission). For
a more detailed discussion, we refer the reader to Tonnesen
& Bryan (2010).
3 RESULTS
3.1Star Formation in the Galactic Disk
We will first compare the star formation in the galactic disks
of the SFNW and SFW cases to determine if and how ram
pressure stripping affects disk star formation. We define the
disk to extend 2 kpc from the the central disk plane, which
includes all of the star formation in the SFNW run. First,
in the top panel of Figure 1 we plot the star formation
rate (SFR) as a function of time for the SFNW and SFW
runs. For the first 220 Myr the SFRs are nearly identical.
The higher m∗min in SFNW only makes the SFR slightly
less smooth over time, and does not affect the agreement of
SFNW and SFW. Shortly after the ICM wind hits the SFW
galaxy disk, its SFR drops precipitously. This figure clearly
shows that ram pressure stripping quickly lowers the SFR
and does not induce even a short-lived burst of star forma-
tion (at least at the level of ram pressure modeled here).
While our exact star formation recipe results in a SFR in
our models that is high for an isolated Milky-Way sized spi-
ral galaxy, we emphasize that it is the comparison between
SFNW and SFW that is important in this work.
In the bottom panel of Figure 1 we plot the radius in-
cluding 95% of the new stars formed in the disk against
time. As in the panel above, we find that the two cases are
nearly identical for the first 200 Myr. However, it takes ∼
70 Myr longer than for the SFR (until ∼290 Myr) for the
star-formation radii to diverge. This is likely because dense
clouds have formed up to 17 kpc from the disk centre that
cannot be instantly stripped by the ICM wind. It is only
after these clouds have formed stars or been ablated — had
their own edges stripped by the wind until they are destroyed
or are of low enough density to be removed from the disk by
the ICM wind — that the star-formation radius drastically
drops.
In the two previous figures we focused on the total
star formation rate. Now we will look at how star forma-
tion rate relates to the gas surface density. Even though
the star formation rate does not spike when the wind hits
the galaxy, it could be high relative to the surface den-
sity of the gas remaining in the disk. In Figure 2 we plot
the Schmidt-Kennicutt relationship of each galaxy for each
timestep (Schmidt 1959; Kennicutt 1989). We focus on the
Figure 1. The top panel plots star formation rate (SFR) as a
function of time in the simulated galaxy with no wind (SFNW,
black line) and in the galaxy that is hit by an ICM wind after
about 210 Myr (SFW, red line). About 10 Myr after the wind
hits the disk, the SFR of the galaxy begins quickly decreasing,
and continues to decrease throughout our run, as the disk gas is
stripped. The bottom panel plots the disk radius (in kpc) that
includes 95% of the new stars formed as a function of time for
the same runs as in Figure 1. Once the wind hits the SFW galaxy
(red), this outer star forming radius decreases, but not as quickly
as the SFR (Figure 1), probably because some dense clouds in
the disk cannot be stripped by the wind, so instead form stars.
outputs at times later than 250 Myr (shown as diamonds
in this figure), as this is both when the SFR becomes con-
stant in the SFNW case (see Fig.
SFW galaxy begins being ram pressure stripped. Limiting
ourselves just to those points after 250 Myr, we can see that
the Schmidt-Kennicutt relationship closely follows the ob-
served relation in both cases, plotted as the solid line (2.5
× 10−4(Σgas/ 1 M⊙ pc−2)1.4). At the latest times in the
SFNW galaxy, the galaxy evolves only very slowly in gas sur-
face density or SFR surface density. The SFW galaxy has a
slightly increasing gas surface density with time because the
outer, lower-density regions are being stripped.
In Figure 3 we plot a “local” Schmidt Law–the SFR
surface density against the gas surface density in rings with
1 kpc width. This relationship is plotted for a single time
snapshot–460 Myr into the simulation, or 250 Myr after the
wind hits the SFW galaxy. We choose this late time because
it maximizes the difference between the SFR and the star-
formation radius of the SFW and SFNW galaxies (as seen in
Figure 1). However, clearly the local relationships between
SFR surface density and gas surface density are very similar
in both runs. They even seem to flatten at about the same
1), and also after the
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6 S. Tonnesen and G. L. Bryan
Figure 2. The Kennicutt-Schmidt relation of the two simulated
galaxies. For each timestep we calculate the gas and SFR surface
density within the radius containing 95% of the new star forma-
tion for that timestep. Plus symbols are used for the first 250
Myr, before the disk has completely settled, and then diamonds
are used thereafter. At early times, the star formation relation lies
below the line denoting the observed Kennicutt-Schmidt relation
(Kennicutt 1989), while at late times the galaxies lie very close
to the Kennicutt-Schmidt Law.
Figure 3. The star formation-surface density relation computed
in rings with widths of 1 kpc. The SFNW run is shown with black
diamonds and SFW is the red triangles, both shown at a time 250
Myr after wind hits disk (460 Myr into the simulation). The solid
line is the empirical Kennicutt relation (2.5 × 10−4(Σgas/ 1 M⊙
pc−2)1.4) (Kennicutt 1989), and the dash-dotted line denotes the
gas surface density at which Leroy et al. (2008) found the local
star formation efficiency flattens.
gas surface density, ∼20 M⊙ pc−2, which is in remarkably
good agreement with the gas surface density at which the
local SF efficiency flattens as observationally found by Leroy
et al. (2008) (14 ± 6 M⊙ pc−2; see their Figure 5). The most
notable difference between our two runs is that the SFW
galaxy has fewer points than the SFNW galaxy. This could
be because either the gas density is zero in a 1 kpc ring
in the galaxy, or the SFR surface density is zero. In fact,
both the gas density and SFR surface density are very low
outside of a radius of about 12 kpc, continuing to hold to a
correlation between gas and SFR surface density.
Thus far we have found that including a ram pressure
stripping wind decreases the total SFR and focuses the SF
towards the centre of the galaxy (Figure 1), but only slightly
increases the gas surface density in the disk and does not
change the relationship between gas surface density and SFR
surface density (Figures 2 and 3). We finally consider the
total amount of newly formed stellar mass in the disk and
in the bulge of the galaxy. As shown in the top panel of
Figure 4, the total stellar mass formed in the disk, since the
beginning of the simulation, is nearly identical in SFNW
and SFW for the first 200 Myr, but once the wind hits the
disk, the two lines begin to diverge. By the end of the SFW
simulation, the SFW galaxy has about 7 × 108M⊙ less
stellar mass in its disk.
If we focus only on the (newly formed) bulge stars, as
defined by all of the stars formed since the simulation began
in a sphere with a 3.4 kpc radius from the galaxy centre
(this includes 80% of the mass in the spherical Hernquist
bulge we initially used to determine our galaxy potential),
we find that including the wind leads to more stars in the
galactic bulge, as shown in the bottom panel of Figure 4. As
we have discussed in Tonnesen & Bryan (2009), we find that
gas clouds that are not stripped are able to spiral towards
the centre of the disk (initially seen by Schulz & Struck
2001). It is this inflow of gas within the disk that adds most
of the stars to the bulge (rather than stellar fallback).
Our bulge-to-total ratio of new stars is 0.1 in the SFNW
galaxy and 0.2 in the SFW galaxy. However, this galaxy ini-
tially had a stellar disk mass of 1011M⊙ and a stellar bulge
mass of 1010M⊙, so the the new stars change the bulge-to-
total ratio of stellar mass by at most a few percent. In sum-
mary, we find that ram pressure stripping lowers the SFR of
a galaxy without an initial burst, despite the fact that our
stripping ICM has a higher pressure than the static ICM.
The relationship between gas density and SFR is similar in
both simulations, with a slight increase in the total gas sur-
face density for the SFW run, as would be expected from
both the ram pressure and the higher-pressure surrounding
ICM. Finally, although ram pressure does cause SFW to
form more stars in the bulge and less stars in the disk than
SFNW, it is not enough to overcome the initial mass profile
of the galaxy. The galaxy would have to have about 2 or-
ders of magnitude less stellar mass for the difference in star
formation to have a significant effect on the bulge-to-total
mass ratio.
3.2Gas in the Galactic Disk
We will now examine if including star formation and feed-
back affects the remaining gas disk of a ram pressure-
stripped galaxy. First we plot the amount of gas in the disk
as a function of time for all three simulations in Figure 5.
Gas is counted as being in the disk if it has a tracer frac-
tion of more than 0.6 and is within 2 kpc of the disk central
plane. Focusing on SFNW (the black solid line), we see that
including star formation results in disk gas being used to
form new stars throughout the simulation. Comparing the
two ram pressure stripped galaxies (SFW: red dash-dotted
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Star Formation in Ram Pressure Stripped Tails7
Figure 4. The top panel shows the total amount of newly formed
stellar mass in the disk (i.e. within 2 kpc of the disk plane), as
a function of time. Shortly after the wind hits the galaxy on
the right, the star formation rate decreases, and by the end of
the SFW run, it has about 7 × 108M⊙ less stellar mass in the
disk than the SFNW run. The bottom panel plots mass of newly
formed bulge stars, as a function of time. If we focus only on
the bulge – all stars within 3.4 kpc of the galaxy centre – we see
that ram pressure does lead to more stars in the bulge. However,
the difference does not significantly change the B/T ratio, which
begins at 0.1 with MBulge= 1010M⊙and MDisk= 1011M⊙.
line and NSFW: blue dashed line), we see that after 250
Myr of stripping (the end of each line), the galaxies lose
very similar amounts of gas. The SFW run has about 109
M⊙ less gas left in the disk than the NSFW galaxy, which
is less than the amount of gas that formed stars (1.4 × 109
M⊙), so 4 × 108M⊙ less gas was actually removed from
SFW than from NSFW by the ICM wind. We conjecture
that this is because dense gas clouds near the outer edges
of the disk formed stars in SFW rather than being ablated
and eventually stripped by the ICM wind.
This picture of outer gas clouds being either stripped
(NSFW) or forming stars (SFW) also agrees with Figure
6, where we plot the radius of dense gas in the disk in all
three simulations (see also Tonnesen & Bryan 2008). This
is the radius of gas with a density above 10−24g cm−3. For
each timestep we consider twelve wedge-shaped sections of
the disk and measure the largest radius at which there is gas
with a density above 10−24g cm−3. Each of these individual
measurements are shown as dash-dotted lines. We also plot
the mean radius of all wedges against time as the thick solid
line. In all three cases, we see the collapse of the disk gas
into dense clouds as the early increase in the radius of this
Figure 5. The amount of gas in the disk in all three simulated
galaxies. The SFW galaxy has less gas than the NSFW galaxy
after being stripped for 250 Myr, but only by 109M⊙. This is
less than the 1.4 × 109M⊙ of stars that form throughout the
SFW simulation.
dense gas. The ram pressure stripping wind affects both the
disk with and the disk without star formation in a similar
fashion: the gas disks have similar ranges of measured radii
and a similar mean radius of about 18-19 kpc.
We next consider how star formation affects the density
and temperature distribution of gas in the galactic disk. In
Figure 7, we show contours of gas mass in the disk, as a
function of gas density and temperature. These are all 250
Myr after the wind has hit the galaxy (or simply 460 Myr
into the SFNW simulation). Including star formation and
feedback spreads the distribution of high density (ρ > 10−23
g cm−3) gas in the disk to include lower densities and higher
temperatures. This may lower the surface density of gas in
the galaxies with star formation, making it easier to strip.
Finally we consider the velocity structure of the gas in
the disk. In Figure 8, we plot contours of gas mass in the disk
(defined as before), as a function of density and gas velocity
in the wind direction, for the same time as in Figure 7. First,
we note that the velocity spread in the SFNW run indicates
the density and velocity of the disk gas that is accelerated
due to the inclusion of thermal feedback from supernovae
(i.e. only gas with ρ ≤ 10−24g cm−3is affected).
Gas accelerated by this process likely explains most of
the differences between the SFW and NSFW cases. The
most obvious difference is in the gas with negative velocities
– while in the NSFW runs, most of the gas with negative
velocities is likely the fallback of stripped gas into the cen-
tral region of the disk, in the SFW simulation a portion of
the gas is also accelerated due to feedback. A smaller effect
is seen in the positive velocity gas – the same density gas
can have a higher positive velocity in the SFW case than in
the NSFW case.
However, this figure tells us that the gas being stripped
(with large positive velocities) has densities less than 10−23
g cm−3, so the differences that star formation and feed-
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8S. Tonnesen and G. L. Bryan
Figure 6. The maximum radius of gas with a density above 10−24g cm−3as a function of time. Dash-dotted lines show this radius for
each of twelve wedges of the disk, while the solid line is the mean of the wedges. From left to right, the panels show SFNW, SFW, and
NSFW. Being stripped by the ICM wind lowers the galaxy radius, but including star formation does not have much affect on the radius
of the dense gas that remains in the disk (at least for 250 Myr of stripping).
Figure 7. Contours of gas mass in the disk (defined as gas with a tracer fraction of more than 0.6 and a height above the disk of less
than 2 kpc), as a function of gas density and temperature. These are all shown 250 Myr after the wind has hit the galaxy (or simply 460
Myr into the simulation with no ICM wind). Including star formation and feedback allows gas in the disk to have lower densities and
higher temperatures which may make gas in the disk with star formation easier to strip.
back cause in the density and temperature distribution of
disk gas (Figure 7) do not strongly affect gas that will be
stripped. Although we do not carry out a series of runs with
higher ram pressures that also include star formation (due to
prohibitively long run times), we have analyzed simulations
with higher ram pressures (but without star formation) in
Tonnesen et al. (2011). In that case, with a ram pressure
of 4 × 1011dynes cm−2, gas with densities as high as 6 ×
10−23g cm−3can be stripped. We predict that star forma-
tion would result in that galaxy being stripped more quickly
than in our simulations with no star formation.
Including star formation only has a slight effect (∼10%)
on the amount of gas left in the disk (Figure 5) and has very
little effect on the residual size of the unstripped gas disk
(Figure 6). We do find noticeable differences in the density-
temperature and z-velocity distributions of the gas in the
disk, likely due to feedback from star formation.
3.3Gas in the Stripped Tail
We turn now to the gas in the stripped tail. First, as with
the disk gas, we plot contours of gas mass as a function of
density and temperature in Figure 9. This figure is a snap-
shot 250 Myr after the wind has hit the disk, including all
of the gas more than 10 kpc above the disk with a tracer
fraction above 0.25 (i.e only gas that originated in the disk).
It is immediately clear that the gas density and tempera-
ture distributions in the tail are very similar whether or not
the simulation includes star formation. In the SFW panel
we have denoted the range of temperatures and densities at
which gas may form stars. The NSFW contours reach some-
what lower densities at low temperatures. In Figure 10 we
plot contours of gas mass as a function of velocity in the
wind direction and height above the disk. Once again, the
distributions are very similar, although on closer inspection
there is a slight difference in the placement of the highest
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Star Formation in Ram Pressure Stripped Tails9
Figure 8. Contours of gas mass in the disk as a function of velocity in the wind direction and gas density, 250 Myr after the wind has
hit the galaxy. Including star formation and feedback causes gas in the disk to have negative and positive velocities. In addition, at any
density, the run that includes star formation has slightly more gas at higher velocities.
Figure 9. Contours of gas mass in the tail (tracer fraction of more than 0.25 and height greater than 10 kpc), as a function of gas
density and temperature, both at a time 250 Myr after the wind has hit the galaxy. Including SF and feedback has a minimal effect on
the ρ-T distribution of the tail. The dotted red lines denote the minimum mass and maximum temperature necessary for stars to form
from a gas cell.
contour (see also Figure 11), and the SFW tail has a nar-
rower velocity distribution at large distances from the disk.
We can see how the similarity in density, temperature,
and velocity plays out in observables, specifically H I col-
umn density and Hα emission. In Figure 11 we display the
projections of the SFW run on the left and the NSFW run
on the right. For these plots, we restrict ourselves to diffuse
Hα emission, neglecting the contribution from HII regions
(see the next section for the stellar contribution). Unlike the
previous figures, here the tails look somewhat different, with
the SFW tail having much of its dense gas farther from the
disk than the NSFW tail. The small difference in the highest
contour of Figure 10 has resulted in a significantly different
distribution of bright H I and Hα emission.
These differences in the distribution of the stripped gas
are due to feedback near the disk. Thermal feedback results
in gas outflows above and below the disk. This means that
some gas that will be removed from the galaxy is already
moving away from the galaxy in the direction of the ICM
wind. This is why dense gas is farther from the disk in the
SFW simulation than in the NSFW simulation. There is not
much star formation in the stripped tail, as we will show be-
low, so it does not have such a large effect on the morphology
of the tail.
As we would expect from the similarity between SFW
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10 S. Tonnesen and G. L. Bryan
Figure 10. Contours of gas mass in the tail as a function of gas velocity in the wind direction and height above the disk. These are
both 250 Myr after the wind has hit the galaxy. Including star formation results in the bulk of the stripped gas being farther above the
disk, although the range of z-velocities is very similar both cases.
and NSFW in Figure 9, the range of H I column densities
and Hα intensities are the same in the two tails. The total
emission from the tails is also very similar. Including star
formation results in slightly less H I gas (20% less), possibly
because some of the most dense gas turns into stars. The
(diffuse) Hα emission is only 2% higher when including star
formation and thermal feedback. Although we do not show
a projection here because the X-ray brightness is too low
to be observed, including star formation reduces the X-ray
luminosity by only a small amount, 13%.
3.4 Star Formation in the Stripped Tail
Finally, we consider star formation in the tail. In Figure 12
we plot the z-velocity of the tail stars against the height
above the disk for a single output 250 Myr after the wind
has hit the disk. We plot these points over the contours of
gas mass.
The escape velocity from the galaxy as a function of
height above the disk is shown by a red dash-dotted line.
We see that most of the stars are moving more slowly than
the bulk of the stripped gas at that height. This is because
the stars are moving at the velocity of the gas when they
are formed, but then are no longer accelerated by the ICM
wind and begin to slow down due to the galaxy’s poten-
tial. Most of the stars with negative velocities are near the
disk, but some stars out to ∼60 kpc have negative velocities,
indicating that they are falling back onto the disk.
If we ran the simulation for long enough (inside a large
box) we would expect all of the stars below the red line
to begin to eventually fall back towards the disk. For most
of these stars to be tidally stripped by the cluster potential
(which we have not considered so far), the tidal radius would
Figure 12. The z-velocity and height above the plane of the
star particles are shown as diamonds (showing only those with
heights above 1 kpc) are overplotted on the SFW gas contours
from Figure 10. The edges of the gas contours are overplotted in
green for clarity. The stars begin with the velocity of the gas in
the tail from which they are formed, and then slow down as they
are no longer accelerated by the wind.
need to be about 60 kpc. If our galaxy were in the Virgo
cluster, which has a velocity dispersion of about 700 km
s−1, it would therefore need to be about 200 kpc from the
cluster centre. In Coma, which has a velocity dispersion of
about 1000 km s−1, the galaxy would need to be about 300
kpc from the cluster centre for tidal stripping to unbind a
significant number of stars.
In Figure 13 we show projections of the stellar surface
density and the surface brightness of Hα from HII regions.
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Star Formation in Ram Pressure Stripped Tails11
Figure 11. Projections of HI column density (top row) and Hα intensity (bottom row). The galaxy with star formation and feedback
(SFW) is on the left, and without (NSFW) is on the right. Including star formation results in slightly longer tails.
First we focus on the left panel, the stellar mass surface den-
sity. We see that there are a few clumps of ∼ 3 × 104M⊙
kpc−2, which are aligned with where the recent star forma-
tion has taken place (compare to the right panel). There is
also a more diffuse component with surface densities about
an order of magnitude less. We can estimate if we should see
these tails in deep images of clusters. Each star particle is
the size of a small cluster of stars that is formed at the same
time using a Salpeter mass function ranging from 0.1-100
M⊙. Mengel et al. (2002) find the Lv/M for young star clus-
ters to range from 0.5-2 for ages ranging from 106-108yr. If
we assume the highest Lv/M, the surface brightness of the
bright knots of ∼ 3 × 104M⊙ kpc−2is ∼29.5 mag/arcsec2.
This is well within the range of V-band surface brightness
of the ICL, and dimmer than the ICL observed by Mihos
et al. (2005). The dimmer, more diffuse stellar component
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12 S. Tonnesen and G. L. Bryan
that is clearly connected to the disk would be very difficult
to distinguish from the general ICL.
In the right panel of Figure 13 we show a projection
of the surface brightness of Hα from HII regions. Our sim-
ulation outputs the mass and formation time of each star
particle, from which it is easy to compute a SFR. The sim-
ulation does not directly calculate an Hα luminosity, so in
order to compare with observations of HII regions (which
are observationally distinguished from diffuse Hα emission
because they are small bright regions), we must assume
that newly-formed (within 10 Myr) star particles produce
Hα emission from HII regions. We use the observation-
ally determined relation between the (recent) star forma-
tion rate and Hα luminosity from Kennicutt (1998): SFR
(M⊙yr−1) = 7.9×10−42L(Hα) (erg s−1). In this calculation
we are simply assuming that Hα emission only measures the
SFR within the last ∼ 10 Myr. We compute the equivalent
star formation rate by selecting only star particles which
are younger than 10 Myr (and hence will have associated
gas and bright Hα emission).
As in previous work (Tonnesen & Bryan 2009; 2010), we
find that ram pressure stripping cannot remove the densest
clouds in the disc without breaking them apart. Stars formed
in the tail must therefore be created from the less dense gas
that has cooled and condensed in the tail. Disk gas (that
is not in dense clouds) with a range of densities is stripped
continuously, so there is a range of gas densities throughout
the tail. Accordingly, the time it takes for radiative cooling
and compression by the ICM, and consequent star forma-
tion to occur, varies throughout the tail as well. Figures 7
and 8 show that gas with ρ < 10−24can be stripped from
the disk and that some gas at these lower densities has T
< 105. Because of these low temperatures, this gas will cool
and condense into clouds rather than mix into the ICM (as
discussed in Tonnesen & Bryan 2010, 2011). The exact tem-
perature and density of the gas will determine how long this
takes, which is why there is a large spread in the height of
stars above the disk. This is illustrated in the two lines in the
upper panel of Figure 14, which shows the SFR in the tail
above either 2 kpc (the black solid line), or 20 kpc (the red
dashed line). Star formation in the tail occurs throughout
the simulation from close to the disk (∼ 2 kpc) to far from
the disk (well beyond 20 kpc, see also Figure 13). Hester et
al. (2010) also find that this scenario of star-forming clouds
condensing within the stripped tail agrees well with their
observation of a tail from a galaxy in the Virgo cluster.
There are clearly stars in the tail, and if unbound,
they can contribute to the ICL. We evaluate this possibility
with Figure 14, which shows the cumulative amount of stars
formed in the tail. We find that 250 Myr after the SFW
galaxy has begun to be stripped by the ICM wind, there is
about 4.2 × 106M⊙ of stellar mass more than 20 kpc above
the galaxy. From Figure 12 we know that this is an overes-
timate of the number of stars that will escape this galaxy.
Therefore we find it unlikely that ram pressure stripping is
a large contributor to the ICL, as 4.2 × 106M⊙ is less than
1% of the stellar mass formed in the disk.
Figure 14. The top panel shows the SFR in the tail gas either
above 2 kpc (solid black line) or above 20 kpc (dashed red line).
There is star formation throughout the tail. At early times the
SFR rate is very low, in rough agreement with the observations of
Gerhard et al. (2002) and Cortese et al. (2003; 2004). The bottom
panel plots the amount of stellar mass more than 20 kpc above
the disk vs time. Although there certainly are stars in the tail,
this amount of stellar mass will not be a large fraction of the ICL
even if it all escapes into the ICM.
4 DISCUSSION
4.1Comparison with previous work
We compare our results with the simulations in K09, as they
both make predictions for star formation in the tails of ram-
pressure stripped galaxies. As we discuss in the introduction,
K09 ran 12 simulations of ram pressure stripped galaxies,
varying the ICM density and velocity. Our simulation is most
similar to their run 2, which has an ICM density slightly
above ours (5 × 10−28g cm−3rather than 3.2 × 10−28g
cm−3) and a wind velocity slightly below ours (1000 km
s−1vs 1413 km s−1). The amount of gas stripped is very
similar – in 250 Myr about 65% of the original gas mass is
stripped from the K09 simulation (their figure 15), while our
SFW run has about 58% of the gas mass stripped (Figure
5). Given that the SFW galaxy is more massive than the
K09 galaxy, this is reasonably good agreement. K09 find
that tail gas with T > 106K has a mean density of about
10−24g cm−3, while our mean density is lower, at ∼ 2 ×
10−25g cm−3. This is both because our surrounding ICM
density is lower and because we are including all gas with
a tracer fraction above 0.6 in this measurement, which we
would expect to skew our results to lower densities. The
results regarding the stripped gas are similar between SFW
and K09.
However, our star formation results are very different
from the K09 results. In direct opposition to our results,
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Star Formation in Ram Pressure Stripped Tails13
Figure 13. Projections of stellar surface density on the left and the surface brightness of Hα in HII regions on the right (computed from
stars formed in the last 10 Myr, see text). We smooth both projections using a 1 kpc gaussian. Both old (stars formed within the stripped
gas but older than 10 Myr) and new stars are well-distributed throughout the tail, reflecting the range of densities and dynamical times
of stripped gas.
they find that adding a ram pressure stripping wind results
in more stars being formed in the simulation. In the run most
similar to ours, after 250 Myr there are nearly as many stars
formed in the wake as in the disk, while in our SFW run only
about 1% of the new stars formed are in the tail.
We checked whether a difference in our threshold den-
sity for star formation could have a large effect on the star
formation rate in our simulations. To do this, we ran a com-
parison simulation identical to SFW in which we allowed
stars to form if gas had a density above 3.85 × 10−26g
cm−3, a factor of 10 below our standard prescription. The
results were quite similar to our standard run. The SFR in
the disk had differences of less than 10% at every output. Af-
ter 460 Myr, the total stellar mass in the disk with the lower
density star-formation threshold was only 4% larger than in
the SFW disk, and the amount in the bulge was only 4%
lower than in the SFW bulge. The tail results differed a bit
more: about 50% more stars were formed in the tail over the
length of the simulation. This makes sense – decreasing the
star formation threshold substantially increases the amount
of gas that can form stars (see Figure 9), but the gas has
a longer dynamical time so the net change in star forma-
tion is not large. This increase in stellar mass formed in the
stripped tail is not nearly large enough to account for the
difference between our results and those of K09. Increasing
our star formation efficiency would probably increase the
stellar mass in the tail, but this could affect our agreement
with the empirical Schmidt-Kennicutt Law.
Our simulations are very different, so there are many
possible reasons for different results, and we will list a few of
the most salient differences now. First, K09 use GADGET-
2, an SPH code, and include a subgrid model that increases
the gas pressure at high density. Their galaxy is less mas-
sive than ours, with a circular velocity of 160 km s−1. It is
also more gas-rich, with a gas fraction of 25% of the total
disk mass. They allow radiative cooling to 104K, and the
maximum temperature at which gas can form stars is 106K.
Their star formation prescription, like ours, is proportional
to the dynamical time of the gas. The reason for the very
different predicted star formation rates in the tail is hard
to pin down, but we suggest two key differences. First, the
subgrid model in the Springel & Hernquist (2003) prescrip-
tion has a very stiff equation of state in dense gas, while our
dense gas typically is cold and has a low thermal pressure.
Second, inspection of the K09 results suggests that stripped
gas does not mix with the ICM, resulting in a high frac-
tion of the stripped gas ending in large, cold clumps. We
note that SPH has difficulties in resolving instabilities at in-
terfaces (Agertz et al. 2007), resulting in undermixing and
reduced stripping. This, combined with the stiff equation of
state and the high ICM pressure, lead to large star formation
rates in the stripped gas.
4.2 Comparison With Observations
We will first compare our results to observations of ram
pressure stripped galaxies in the Virgo cluster, which have
largely been identified due to their H I tails (but have rela-
tively little star formation). We will then compare our results
to observations of H II regions and/or stellar tails associated
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14 S. Tonnesen and G. L. Bryan
with galaxies that are likely undergoing ram pressure strip-
ping in more massive clusters. Finally, we discuss the key
physics that control the star formation rate in our simulated
tails.
4.2.1 Virgo Tails
As we have discussed in the Introduction, eleven galaxies
in Virgo have clear stripping signatures in H I observations,
but only four have been found to have star formation either
from UV emission or H II regions. This begs the question
of whether we should expect star formation in our stripped
tail. To answer this question we will compare our results to
some of these observations.
Let us first consider the galaxies that have star for-
mation in their tails. Cortese et al. (2003) found an H II
region 3 kpc from the disk in the stripped tail of NGC 4402.
Their measured Hα luminosity results in a SFR (using the
Kennicutt (1998) equation) of 2.3 × 10−3M⊙ yr−1. Our
simulated ram pressures are similar to that likely experi-
enced by Virgo galaxies, and we have SFRs of tail gas (more
than 2 kpc above the disk) ranging up to 6 × 10−2M⊙
yr−1(over the 250 Myr that the galaxy is being stripped).
If we choose an early output, our SFR agrees with the ob-
servations of Cortese et al. (2003; 2004). However, Crowl &
Kenney (2008) use stellar populations in the disk to predict
that this galaxy has been stripped for nearly 200 Myr. This
means that we overpredict both the SFR and the distance
to which star formation would be observed in this galaxy.
Of course, because we are not actually modeling this galaxy,
our galaxy velocity is slightly higher than that expected by
Crowl et al. (2005), and our galaxy is being stripped face-on
rather than at an angle. Our ICM density is very similar to
that near NGC 4402 (Schindler et al. 1999). Similarly, NGC
4330, NGC 4522, and NGC 4438 have ongoing star forma-
tion in their stripped tails (Abramson et al. 2011; Kenney &
Koopmann 1999; Boselli et al. 2005). As in NGC 4402, the
observed stars in the tails are closer to the galaxy than the
more extended tails we simulate.
Finally, we discuss IC 3418, which has not been ob-
served in H I, but has a UV and Hα tail (Martin et al.
2005; Hester et al. 2010). IC 3418 is likely to be in a higher-
density ICM than we model by about a factor of 5. This UV
tail extends 17 kpc from the disk, and the authors calculate
a lower limit for the SFR of 6 × 10−3M⊙ yr−1(because
they do not correct for any dust extinction). They use the
star formation truncation time to estimate that this galaxy
has been stripped for 100 Myr. After 100 Myr of stripping,
our stellar tail is about 20 kpc, in good agreement with this
observation. However, our simulated SFR is still a factor of
∼3 above that in IC 3418.
There are, in addition, seven galaxies that do not have
any stellar light associated with their H I tails. Six of these
galaxies are at or beyond 700 kpc from M87, and so may be
in lower-pressure ICM regions than we simulate. The excep-
tion is NGC 4388, which, based on the models of Vollmer &
Huchtmeier (2003), may be in a similar, or higher, density
region of the ICM than we use in our simulations.
Although we are not attempting to directly model any
single galaxy, in general we have a higher SFR in our tail
and a longer stellar tail than in most of these galaxies. There
are a number of reasons we find higher star formation in our
stripped tail. (i) We may be overestimating the amount of
star formation in the tail due to our star formation method.
We could lower the SFR in our tail by changing our star
formation criteria–while we saw in Section 4.1 that lowering
the star formation density threshold by an order of magni-
tude only changed the amount of star formation by 50%, we
could raise the threshold to the point where there was very
little star formation in the tail. (ii) We may also be overesti-
mating the survival of star-forming clouds to large distances
above the disk, a point we will discuss in more detail in Sec-
tion 4.4.2. (iii) We may also have more star formation in
our tail because we have a face-on wind that can strip more
gas, or (iv) because we are modeling a higher-pressure ICM
than surrounds most of the observed galaxies with H I tails
in Virgo. This point will be discussed in more detail below.
4.2.2 Stellar Tails in Massive Clusters
Turning to more massive clusters, Yoshida et al. (2008) ob-
served “fireballs” around a merger galaxy in the Coma clus-
ter. They find that these blue or Hα emitting filaments are
found on one side of the galaxy, extending up to 80 kpc
from the disk. The morphology of this tail of “fireballs” is
roughly in agreement with our simulation. They find that
the Hα knots and filaments are farther from the disk than
the blue knots and filaments, which also tends to be the case
in our tail (Figure 13), but is not as clear-cut as in the case
of RB 199. However, they only find 13 knots and filaments,
while it is clear that we have many more star particles. Fur-
ther, they estimate the mass of the stars in their tail to add
up to 108M⊙–at least a factor of 25 larger than the stellar
mass in our stripped tail (but see below).
Recently Yagi et al. (2010) observed 14 galaxies with
stellar Hα clouds in tails in the Coma cluster. Their images
show a range of possible tails, some of which look more like
our simulated tail than others. A number of their tails show
very linear trails of either young stars or H II regions, which
is in qualitative agreement with the right panel of Figure
13. Finally, Sun et al. (2007) focus on the H II regions in
the stripped tail of ESO 137-001. They find 29 H II regions
with Hα luminosities ranging from 1038to 1040erg s−1.
They calculate that the total mass of all the H II regions
should be about 107M⊙, and using the Kennicutt (1998)
relationship between Hα luminosity and SFR estimate an
instantaneous SFR of about 0.7 M⊙ yr−1. This includes a
bright H II region slightly less than 2 kpc above the disk,
and without this H II region the SFR would be about 0.53
M⊙ yr−1. Both are larger than what we find, by about an
order of magnitude.
Why do we predict significantly less star formation in
the stripped tail than in these observations? Yoshida et al.
(2008) determine an ICM density much like the one in our
simulation, so there should not be an increased SFR in tail
gas due to higher ICM density. A possible explanation is
that RB 199 has the disturbed morphology of a merger rem-
nant, which may have moved large star-forming clouds far-
ther from the centre of the galaxy to regions where they
could be more easily stripped by ram pressure.
We may have less stellar mass in our tails than in ESO
137-001 and most of the galaxies observed by Yagi et al.
(2010) because we simulate a lower density ICM than in
those regions of the Coma and Norma clusters. Kapferer et
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Page 15

Star Formation in Ram Pressure Stripped Tails15
al. (2009) found that increasing the ICM density increased
the star formation in the tail, which might cause stripped
clouds to be compressed more quickly. This agrees with ob-
servations of the lower Hα luminosities of H II regions in
ram pressure stripped galaxies in the lower-density ICM of
the Virgo cluster (Kenney & Koopmann 1999; Cortese et al.
2004), as highlighted in the previous section.
In order to test this idea using our simulations, we can
make a rough estimate of the star formation rate of gas
in the tail in our two simulations modeled after ESO 137-
001, and examined in detail in Tonnesen et al. (2011) (run-
ning these simulations including star formation at the same
resolution as SFW is too computationally costly). From
equation (1) we can calculate the star formation rate sim-
ply by knowing the gas mass and density in a cell, using
tdyn = (3π/32Gρ)1/2and recalling that we have an effi-
ciency of 0.5%. We first test this method by comparing the
star formation rate estimated in this way to the measured
rate for SFNW and SFW. First, we consider just the disk
gas, and find that, averaged over time, the estimated rate
is within 1% of the measured rate, with a variation ranging
from 10% too large, to 2% too small. More importantly, in
the tail of the SFW run, the estimated rate is within about
10% of the measured rate, with a similar level of scatter.
Now that we have confirmed that the errors in our esti-
mation scheme are low in comparison to the measured rates,
we can predict the star formation rate in our two simulations
using similar ICM conditions to those around ESO 137-001
(Sun et al. 2006; 2010; see Tonnesen et al. 2011 for details of
these runs). Using the gas between 2 and 40 kpc above the
disk in order to match the observations of Sun et al. (2006),
we predict a SFR of 0.094 M⊙ yr−1(for the T3vl run, which
had somewhat lower ram and thermal pressures) and 0.32
M⊙ yr−1(for the T3vh run, which had higher pressures).
We see that the T3vh estimate is within a factor of ∼2.5
of the observed Sun et al. (2007) value of 0.7 M⊙ yr−1(or
within a factor of 2 of the corrected value of 0.53 M⊙ yr−1,
beyond 2 kpc above the disk), while the T3vl estimate is in
poorer agreement. As discussed in Tonnesen et al. (2011),
the T3vl case also does not agree with the non-detections of
H I in the tail, and these star formation estimates lend more
credibility to our claim that T3vh is a stronger match to
the observations of ESO 137-001. We also predict that there
will be star formation in the stripped tail between 40-80 kpc
with a SFR of about 0.075 (T3vh, 0.51 for T3vl) M⊙ yr−1,
which, if correct, will make it very difficult to observe.
4.2.3 What drives the rate of star formation in the tail?
Finally, in this section we try to determine the key physical
effect that determines the star formation rate in the tail. To
do this, we compare the T3vl and SFW runs – the velocities
in the two runs are similar, and although the different ram
pressures (T3vl is about a factor of 10 larger than the SFW
run) lead to different mass loss rates, we can compare the
two runs when there are similar amounts of gas in the tail.
When we do that, the T3vl run still produces a much larger
(estimated) star formation rate, indicating it is not simply
the amount of stripped gas. Instead, we argue that, the en-
hanced rate is largely due to the increased thermal pressure
in the ICM. Since the cold tail gas is largely in pressure equi-
librium with the ICM, increasing the ICM pressure moves
that gas to higher density (see Figure 9), increasing the star
formation rate. Of course, the amount of stripped gas is also
important, which is controlled by the ram pressure strength,
but once the gas is stripped, it is largely the ICM pressure
which controls the star formation rate in the wake.
4.3 Star Formation Recipe
As noted in Section 2.1, we allow stars to form in gas with
densities greater than 3.85 × 10−25g cm−3and tempera-
tures below 1.1 × 104K. We do not; however, require that
the mass of the cell exceeds the Jeans mass, or that the
cooling time be less than the dynamical time, as used in,
for example, Tasker & Bryan (2006). If we used more strict
star formation criteria, we would expect our star formation
rate to decrease, although based on inspection of Figure 7,
the density criteria would have to change by more than an
order of magnitude before the mass of gas able to form stars
would be significantly affected. Nevertheless, we note that
our more generous star formation criteria make our con-
clusions regarding intracluster light and the change of the
bulge-to-disk ratio conservative. In addition, we note that
our results are based on a comparison between the SFR in
SFNW and SFW and so should not be affected by our exact
star formation recipe.
Our star formation recipe may have a greater impact on
the SFR we measure in the stripped tail of gas. We discuss
this in Section 4.1, but it is worth reiterating here. If we
used more strict star formation criteria, we may indeed see
less star formation in the tail, possibly in better agreement
with observations of Virgo galaxies. While it is possible to
change our calculated values of the SFR by changing our star
formation recipe, our conclusion that higher ICM pressure
results in more star formation is robust.
4.4 Resolution
4.4.1 Star Formation
This work has included star formation in order to determine
how ram pressure stripping affects star formation rates in
the disk and tail of a stripped galaxy. Therefore it is im-
portant to discuss how resolution may affect our results.
As we discussed above (Section 2.1), star formation can oc-
cur in gas with densities greater than 3.85 × 10−25g cm−3
and temperatures below 1.1 × 104K. This corresponds to
a Jeans length of about 1.9 kpc, which is resolved by 50
cells at the finest level of resolution. These cases meet the
Truelove criterion (Truelove et al. 1997), which requires a
minimum of four cells per Jeans length. However, there is
some gas with much higher densities and lower temperatures
in both the disk and tail (Figures 7 & 9). Some of this high
density gas has a Jeans length of less than 10 pc, so our
simulations may include artificial fragmentation that could
increase our star formation rate. Our densest gas is found in
the galaxy disk, and we have found that our galaxies do lie on
the Kennicutt-Schmidt relation. Further, our newly formed
stellar mass closely matches that predicted by Equation 2.1.
However, this does not prove that we do not have artificial
fragmentation on small scales in high-density clouds, and so
our exact measures for the SFR in the galaxy disks could be
incorrect. The comparisons between the SFNW and SFW
c ? 2011 RAS, MNRAS 000, 1–??