Galactic Wind Signatures around High Redshift Galaxies
arXiv:0704.0652v1  [astro-ph]  4 Apr 2007
Draft version November 28, 2021
Preprint typeset using LATEX style emulateapj v. 11/26/04
GALACTIC WIND SIGNATURES AROUND HIGH REDSHIFT GALAXIES
Daisuke Kawata
and Michael Rauch
Draft version November 28, 2021
ABSTRACT
We carry out cosmological chemodynamical simulations with different strengths of supernova (SN)
feedback and study how galactic winds from star-forming galaxies affect the features of hydrogen (HI)
and metal (CIV and OVI) absorption systems in the intergalactic medium at high redshift. We find
that the outflows tend to escape to low density regions, and hardly affect the dense filaments visible
in HI absorption. As a result, the strength of HI absorption near galaxies is not reduced by galactic
winds, but even slightly increases. We also find that a lack of HI absorption for lines of sight (LOS)
close to galaxies, as found by Adelberger et al., can be created by hot gas around the galaxies induced
by accretion shock heating. In contrast to HI, metal absorption systems are sensitive to the presence
of winds. The models without feedback can produce the strong CIV and OVI absorption lines in
LOS within 50 kpc from galaxies, while strong SN feedback is capable of creating strong CIV and
OVI lines out to about twice that distance. We also analyze the mean transmissivity of HI, CIV, and
OVI within 1 h−1 Mpc from star-forming galaxies. The probability distribution of the transmissivity
of HI is independent of the strength of SN feedback, but strong feedback produces LOS with lower
transmissivity of metal lines. Additionally, strong feedback can produce strong OVI lines even in cases
where HI absorption is weak. We conclude that OVI is probably the best tracer for galactic winds at
high redshift.
Subject headings: galaxies: kinematics and dynamics —galaxies: formation —galaxies: stellar content
1. INTRODUCTION
Supernova (SN) explosions are thought to be capa-
ble of ejecting part of the interstellar medium (ISM)
from galaxies. Such outflows are often called “galactic
winds” (Johnson & Axford 1971; Mathews & Baker 1971;
Veilleux et al. 2005). Galactic winds are believed to be
an important mechanism for enriching the intergalactic
medium (IGM) (Ikeuchi 1977; Aguirre et al. 2001; Madau
et al. 2001; Scannapieco et al. 2002; Cen et al. 2005) and
are thought to play a crucial role in shaping the mass-
metallicity relation of galaxies (e.g. Larson 1974; Dekel &
Silk 1986; Arimoto & Yoshii 1987; Gibson 1997; Kawata
& Gibson 2003), and in heating the IGM (Ikeuchi & Os-
triker 1986). Galactic winds have been observed in lo-
cal star-forming galaxies (e.g. Lynds & Sandage 1963;
Martin 1998; Ohyama et al. 2002), where their outflow
morphology and kinematics have been extensively stud-
ied (e.g. Martin 2005; Rupke et al. 2005). Some local
galaxies show outflows to about 20 kpc (Veilleux et al.
2003). Such outflows are expected to be more common at
high redshift, where star formation is more active (z > 1)
(Madau et al. 1996). Galactic winds are also believed to
terminate star formation in ellipticals (Mathews & Baker
1971; Kawata & Gibson 2003) and have been invoked
at high redshift to explain the high age of their stel-
lar populations (Kodama & Arimoto 1997; Labbé et al.
2005; Kriek et al. 2006). Theoretical studies suggest that
for progenitors of disk galaxies the gas outflow due to
SN heating at high redshift can effectively suppress star
formation, leading to a less dense stellar halo (Brook
et al. 2004) and a larger disk (Sommer-Larsen et al.
1 The Observatories of the Carnegie Institution of Washington,
813 Santa Barbara Street, Pasadena, CA 91101
2 Swinburne University of Technology, Hawthorn VIC 3122,
Australia
2003; Robertson et al. 2004; Governato et al. 2006) at
z = 0. Therefore, observing galactic wind signatures at
high redshift may elucidate an important ingredient in
the formation of galaxies.
The observational studies at high redshift have uncov-
ered evidence that outflows from star-forming galaxies
at high redshift may be common, as has been shown for
Lyman break selected galaxies (e.g. Pettini et al. 1998;
Ohyama et al. 2003; Shapley et al. 2003). These results
have contributed to the debate about how the outflows
from star-forming galaxies may affect the intergalactic
medium (IGM). Rauch et al. (2001a) uncovered evidence
for repeated injection of kinetic energy into higher den-
sity, CIV absorbing gas, possibly driven by recent galac-
tic winds. In a study of the lower density, general Lyman
alpha forest, Rauch et al. (2001b) found that most that
of the HI absorption systems lack signs for being dis-
turbed by winds, and derived upper limits on the filling
factor of wind bubbles. Simcoe et al. (2002) surveyed
the properties of strong OVI absorption systems at high
redshift and proposed that the apparent temperatures
and the kinematics of the OVI gas as well as their rate
of incidence could be explained if massive Lyman break
galaxies are driving winds out to 50 proper kpc.
Adelberger et al. (2003, 2005a) studied the absorption
line features in the spectra of background QSOs whose
line of sight (LOS) passes close to Lyman-break selected
star-forming galaxies. They found a deficit of neutral hy-
drogen near these galaxies out to 0.5 h−1 comoving Mpc,
accompanied by a surplus of HI beyond that radius, and
suggested that most Lyman break galaxies may reside
bubbles where superwinds have depleted the HI in the
interior and piled up more neutral gas beyond the hot
bubble. The more recent, more statistically significant
of these studies Adelberger et al. (2005a) however, does
not support this claim, but still appears to show a sig-
http://arxiv.org/abs/0704.0652v1
2 Kawata and Rauch
nificant fraction (about 7 out of 24) of the LOS exhibit-
ing weak or no absorption within 1 h−1 comoving Mpc.
Numerical simulations of the IGM without SN feedback
predict a much lower fraction of such weak absorption
systems near the galaxies (Kollmeier et al. 2003), a fact
that may conceivably be explained if at least some galax-
ies have outflows destroying HI in the IGM in their vicin-
ity. Adelberger et al. (2003, 2005a) also show that there
is a correlation between the spatial distribution of CIV
absorption lines and the star-forming galaxies, with the
strongest CIV absorption lines being observed at the LOS
closest to the galaxies (∼ 80 proper kpc). This result may
indicate that outflows related to recent star formation ac-
tivity have enriched the IGM locally. Other evidence for
the association of CIV with galaxies has been reported
(Pieri et al. 2006; Simcoe et al. 2006). Scannapieco et al.
(2006b) compared the observed LOS correlation func-
tions of CIV and SiIV with analytic outflow models and
concluded that the observed correlation can be formally
explained if there are outflows with a scale of about 2
comoving Mpc from large galaxies whose stellar mass is
about 1012 M⊙. However, the large range of the indi-
vidual wind bubbles required to explain the CIV- galaxy
correlation is puzzling. Theoretical arguments suggest
that the clustering of CIV with Lyman break galaxies is
not necessarily proof that those same galaxies produced
the metal enrichment (Porciani & Madau 2005; Scanna-
pieco 2005), and the relation between metals in the IGM
and galaxies clearly needs further study.
Cosmological numerical simulations have proven a use-
ful tool for understanding metal absorption lines (Rauch
et al. 1997). Comparisons between the observations and
the statistics of metal absorption lines derived from nu-
merical simulations have mostly been concerned with
measurements of the metallicity and ionization state of
the IGM (e.g. Davé et al. 1998; Aguirre et al. 2002;
Schaye et al. 2003; Aguirre et al. 2004). The relation
between star-forming galaxies and absorption line fea-
tures have recently been studied by a number of authors
(Croft et al. 2002; Kollmeier et al. 2003; Bruscoli et al.
2003; Kollmeier et al. 2006; Tasker & Bryan 2006), with
several papers (e.g. Croft et al. 2002; Kollmeier et al.
2006; Theuns et al. 2002) considering observable effects
of galactic outflows on the Lyman alpha forest absorption
lines. These studies generally have concluded that such
outflows hardly affect the strength of the HI absorption
lines, because the winds tend to escape into less dense
regions and do not impact the IGM where the density is
high enough to produce HI absorption lines.
So far, most studies based on full-blown cosmologi-
cal numerical simulations have focused on HI absorp-
tion lines. Although there have been a number of pa-
pers discussing the origin and properties of metal ab-
sorption lines (e.g. Rauch et al. 1997; Theuns et al. 2002;
Aguirre et al. 2005; Tasker & Bryan 2006; Oppenheimer
& Davé 2006), the relation between metal absorption
line systems and outflows from the coeval galaxy pop-
ulation has remained largely unclear, in particular as
the effect of winds on the physical properties of metal
lines it is not well known. Nevertheless, the fact that
the metallicity is expected to be increased by galactic
outflows, and the availability of multiple transitions and
several heavy elements with a different enrichment his-
tory should make metal lines a potentially much more
Fig. 1.— Number of SN and chemical yields as a function of the
age of a star particle with a mass of 1 M⊙. The upper panel shows
the total number of SN. The middle and lower panels present the
total ejected carbon and oxygen masses, respectively. The thin
solid, thick solid, and gray thick dotted lines indicate the history
of a star particle with the metallicity of Z = 10−4, 0.02, and 0.04,
respectively.
useful tracer of galactic winds than neutral hydrogen.
The current paper studies the properties of both HI
and metal absorption lines, and looks for observational
signatures of galactic winds. To this end, we run cosmo-
logical simulations with the original version of the Galac-
tic Chemodynamics Code, GCD+ (Kawata & Gibson 2003)
which is capable of tracing the chemical evolution of the
IGM and galaxies self-consistently. We carry out simula-
tions with different strengths of SN feedback, and com-
pare the features of the HI and metal absorption lines
between the simulations, attempting to identify features
sensitive to the presence of galactic winds.
The following section will explain our method including
the description of the numerical simulations with GCD+
and analysis of absorption lines. Unlike previous studies,
we follow different heavy elements separately, and take
into account the abundance evolution for the different
elements when creating fake QSO spectra from the sim-
ulation. This is important because the different elements
come from different types of SNe or are due to mass loss
from inter-mediate mass stars with different life-times.
Section 3 shows our results. First, Section 3.1 focuses
on one galaxy in the simulation volume, and compares
the absorption line features around the galaxy among
models with different SN feedback strengths. Then, in
Section 3.2 we discuss the results more quantitatively us-
ing artificial QSO spectra in 1000 random LOS. Section
4 summarizes our conclusions.
Galactic Wind Signatures 3
Fig. 2.— Projected gas surface density (upper) and temperature (middle) map and star particle distribution (lower) for models NF (left),
SF (middle), and ESF (right).
Fig. 3.— The history of the star formation rate down to z = 2.43 for models NF (left), SF (middle), and ESF (right). The time of first
star formation is different among the three models, although there should not be any difference before feedback from stars happens. This
difference is because different models are carried out on different computers, and the star formation model in GCD+ uses the random number
generator (see Kawata & Gibson 2003, for details) whose sequences are different for different simulations.
2. METHOD
2.1. Numerical Simulations
The simulations were carried out using the Galactic
Chemodynamics Code GCD+ (Kawata & Gibson 2003).
GCD+ is a three-dimensional tree N -body/smoothed par-
ticle hydrodynamics (SPH) code that incorporates self-
gravity, hydrodynamics, radiative cooling, star forma-
tion, SN feedback, and metal enrichment. GCD+ takes
into account chemical enrichment by both Type II
(SNe II) (Woosley & Weaver 1995) and Type Ia (SNe Ia)
(Iwamoto et al. 1999; Kobayashi et al. 2000) SNe and
mass loss from intermediate-mass stars (van den Hoek
& Groenewegen 1997), and follows the chemical enrich-
ment history of both the stellar and gas components of
the system. Figure 1 shows the total number of SN
4 Kawata and Rauch
Fig. 4.— Overdensity, temperature, metallicity, and [C/O] map (from left to right) for model NF at Y= +50 (upper), 0 (middle) −50
(lower) proper kpc, where the biggest galaxy is set to be at the center.
TABLE 1
Properties of the cental galaxy at z = 2.43
Model ESN Mvir
a rvir
b Mgas,vir MDM,vir Mstar,vir M200
c Tvir
Name (erg) (M⊙) (kpc) (M⊙) (M⊙) (M⊙) (M⊙) (K)
NF 0 2.1× 1011 57 1.7× 1010 1.7× 1011 2.5× 1010 2.0× 1011 3.8× 105
SF 3× 1051 2.0× 1011 57 1.5× 1010 1.7× 1011 1.6× 1010 1.9× 1011 3.7× 105
ESF 5× 1051 1.7× 1011 56 1.2× 1010 1.6× 1011 1.7× 109 1.6× 1011 3.2× 105
aVirial Mass in the definition of Kitayama & Suto (1996)
bVirial radius in the definition of Kitayama & Suto (1996)
cMass within a radius which is the radius of a sphere containing a mean density of 200 times the critical
density at z = 2.43
bVirial temperature in the definition of Kitayama & Suto (1996)
(both SN II and SN Ia) and the total amount of car-
bon and oxygen ejected from a star particle with the
mass of 1 Msun as a function of its age. Initially, SNe
II go off, and they continue until the 8 Msun star dies
(∼0.04 Gyr in the case of Z = 0.02). There is no SN
until SNe Ia start to occur around 0.7 Gyr. A star par-
ticle with Z = 10−4 does not lead to SN Ia, because the
adopted SNIa model restricts the metallicity range for
progenitors of SN Ia to logZ/Z⊙ > −1.1 (see Kobayashi
et al. 2000, for details). Oxygen is produced mainly
by SN II. After SN II ceases, the continuous ejection
of oxygen and carbon is mainly due to the contribution
from intermediate-mass stars. Although oxygen yield is
mainly from the pre-enriched ejecta, carbon is newly pro-
cessed in intermediate-mass stars, which explains the sig-
nificant yield in the low metallicity case.
The adopted version of the code also includes non-
equilibrium chemical reactions of hydrogen and helium
species (H, H+, He, He+, He++, H2, H
2 , H
−) and their
cooling processes, following the method of Abel et al.
Galactic Wind Signatures 5
Fig. 5.— Same as Fig. 4, but for model SF.
(1997); Anninos et al. (1997); Galli & Palla (1998). The
details of the non-equilibrium chemical reactions are de-
scribed in the Appendix of Kawata et al. (2006). We
have made the following update from the code used in
Kawata et al. (2006). We adopt a density threshold for
star formation, and permit star formation from gas whose
hydrogen number density (nH = fHρg/mp, where fH, ρg
and mp are the hydrogen mass fraction, density, and pro-
ton mass for each gas particle) is higher than 0.01 cm−3
(Schaye 2004). It is also crucial to take into account the
effect of the UV background radiation when studying the
properties of the IGM. We use the UV background spec-
trum suggested by Haardt & Madau (2001). The code
follows non-equilibrium chemical reactions of hydrogen
and helium species subjected to the UV background. In
addition, radiative cooling and heating due to heavy el-
ements are taken into account based on the Raymond-
Smith code (Raymond & Smith 1977), used in Cen et al.
(1995). The simulation starts at z = 29.7, and initial
temperature and the fractions of hydrogen and helium
species are calculated by RECFAST (Seager et al. 1999,
2000). We turn on the UV background radiation at z = 6
(Becker et al. 2001; Fan et al. 2001).
The cosmological simulation adopts a Λ-dominated
cold dark matter (ΛCDM) cosmology (Ω0=0.24,
Λ0=0.76, Ωb=0.042, h = 0.73, σ8 = 0.74, and ns = 0.95)
consistent with the measured parameters from three-year
Wilkinson Microwave Anisotropy Probe data (Spergel
et al. 2006). We use a multi-resolution technique to
achieve high-resolution in the regions of interest, includ-
ing the tidal forces from neighboring large-scale struc-
tures. The initial conditions for the simulations are con-
structed using the public software LINGER and GRAFIC2
(Bertschinger 2001). Gas dynamics and star formation
are included only within the relevant high-resolution re-
gion (∼6 Mpc at z=0); the surrounding low-resolution
region (∼55 Mpc) contributes to the high-resolution re-
gion only through gravity. Consequently, the initial con-
dition consists of a total of 1350380 dark matter parti-
cles and 255232 gas particles. The mass and softening
lengths of individual gas (dark matter) particles in the
high-resolution region are 7.61 × 106 (3.59 × 107) M⊙
and 1.15 (1.93) kpc, respectively. The high-resolution
region is chosen as the region within 8 times the virial
radius of a small group scale halo with the total mass of
Mtot = 3 × 10
12M⊙ and the virial radius of rvir = 380
kpc at z = 0.
We simulate the following three models with these ini-
tial conditions to investigate the effect of SN feedback.
Model NF is a ”no-SN-feedback” model: although the
model follows the chemical evolution due to SNe and
mass loss from stars, we ignore the effect of energy feed-
6 Kawata and Rauch
Fig. 6.— Same as Fig. 4, but for model ESF.
back by SNe. Model SF is a ”strong feedback” model,
where each SN yields the thermal energy of 3× 1051 erg.
This model produces a feedback effect noticeable in a
number of observables. In the final model, ESF, for ”ex-
tremely strong feedback” model, a thermal energy release
of 5 × 1051 erg per SN is assumed. We found that this
model causes too strong effects of feedback, and produces
too few stars. Therefore, the model is obviously unreal-
istic. However, we retain the model for this extreme case
to help put the other models in perspective.
We analyze the properties of the IGM for all the mod-
els at z = 2.43. As mentioned above, we adopt the
multi-resolution technique. We extract a spherical vol-
ume within the radius of rp = 800 kpc (in physical scale
at z = 2.43) from a galaxy in the high-resolution region.
The central galaxy is the biggest galaxy in the simulation
volume, and the radius is chosen to avoid contamination
from the low-resolution particles. In this paper, we use
the coordinate system where the central galaxy resides at
(x,y,z)=(0,0,0). Figure 2 shows the gas density, temper-
ature, and stellar density map of the central 800 × 800
proper kpc2 region of this volume analyzed for all the
models. Figure demonstrates that the stronger feedback
affects the gas density distribution, and suppresses star
formation more dramatically.
2.2. The Artificial QSO Spectra
The aim of this paper is to search for signatures of
galactic winds among the absorption lines in the back-
ground QSO spectrum. To this end, we construct artifi-
cial QSO spectra, with lines of sight through the simula-
tion volume from various orientation and projected po-
sitions, and compare the absorption line features among
models with different strengths of SN feedback. For a
given line of sight we identify the gas particles whose
projected distance is smaller than their SPH smoothing
length. In this paper, we focus on three absorption fea-
tures, HI1216, CIV1548, and OVI1032, and we call them
HI, CIV, and OVI hereafter. The ionization fractions for
HI, CIV, and OVI for each gas particle are derived as fol-
lows. The HI fraction is self-consistently calculated in our
simulations, because GCD+ follows the non-equilibrium
chemical reactions of the hydrogen and helium ions. The
CIV and OVI fractions for each gas particle are analyzed
with version 6.02b of CLOUDY, described by Ferland
et al. (1998), assuming the condition of optical thin and
ionization equilibrium. Here, we put in the density, tem-
perature, and the abundances of different elements for
the gas particles in the simulation, and run CLOUDY
adopting the same UV background radiation as used in
the simulations (Haardt & Madau 2001). Unfortunately,
we realized that even our (unrealistically) strong feed-
Galactic Wind Signatures 7
Fig. 7.— Density of HI, CIV, and OVI and OVI weighted temperature map (from left to right) for model NF at Y= +50 (upper), 0
(middle) −50 (lower) proper kpc, where the biggest galaxy is set to be at the center. The arrows represent the velocity field weighted by
HI, CIV, OVI, and OVI in the panels from left to right, respectively. The size of arrow corresponds to the amount of velocity, as indicated
in the upper right panel.
back model cannot enrich the lower density regions in
the IGM as much as what is observed (Cowie & Songaila
1998; Schaye et al. 2003; Aguirre et al. 2004). For ex-
ample, some filaments are not enriched at all even in
models ESF as will be seen in Figure 6, although quanti-
tative comparisons with the observational data (see also
Oppenheimer & Davé 2006) will be pursued in a future
paper. This is likely because the limited resolution of our
simulations is unable to resolve the formation of smaller
galaxies which form at a higher redshift and could enrich
the IGM. Thus, to mimic pre-enrichment for the low den-
sity regions we add metals at the level of [C/H]= −3 and
[O/H]= −2.5 to all gas particles. Once the ionization
fractions are obtained, the column densities at the LOS
for each species are analyzed for each gas particle, using
the two-dimensional version of the SPH kernel. The op-
tical depth τ(v) profiles along the LOS are calculated by
the sum of the Voigt-absorption profiles for each particle,
taking into account their temperature and LOS velocity,
vi,LOS, which is the sum of Hubble expansion and pecu-
liar velocity. The final spectra are constructed, assum-
ing an overall signal-to-noise ratio of 50 per 0.04 Å pixel,
the read out signal-to-noise ratio 500, and FWHM=6.7
km s−1. We stress that we take into account the dif-
ference between carbon and oxygen abundances in our
chemodynamical simulations, when obtaining CIV and
OVI fraction.
Note that the simulation volume analyzed is only 1.6
proper Mpc scale at z = 2.43. The Hubble expansion at
z = 2.43 is 236 km s−1 Mpc−1 in the adopted cosmol-
ogy, and 1.6 proper Mpc corresponds to 378 km s−1. In
addition, since the volume is overdensity region, the ex-
pansion velocity is smaller than the Hubble expansion. In
Section 3.2, we analyze the mean transmissivity within 1
h−1 comoving Mpc from the galaxies. The velocity that
corresponds to 1 h−1 comoving Mpc is 94.2 km s−1 at
z = 2.43, which is well within the range of the volume.
However, in the real Universe, if there are absorbers out-
side the volume at the LOS and their peculiar velocity is
large, they can contribute to absorptions in the velocity
range we focus on (see Kollmeier et al. 2003, for more
detailed discussion of such effect). Therefore, we ignore
the contamination from such absorptions, and consider
the ideal absorption systems only by the absorbers which
are spatially close.
3. RESULTS
3.1. Absorption features around a galaxy
To study how the gas outflows from galaxies affect the
absorption line features, we generate two sets of QSO
8 Kawata and Rauch
Fig. 8.— Same as Fig. 7, but for model SF.
spectra for all the models. In this section, we analyze the
spectra whose LOS are chosen as the 5 × 5 grid points
each separated by 50 kpc (in physical scale at z = 2.43) as
projected onto the plane of the sky. The grid is centered
on a galaxy. In the next section, we generate 1000 ran-
dom LOS spectra, and compare them among the three
models. The properties of the central galaxy chosen for
the present section for different feedback models are sum-
marized in Table 1. The total mass of the central galaxy
is slightly smaller than the estimated mass of the BX
galaxies (M200 ∼ 6.3 × 10
11 − 1.6 × 1012 M⊙ in Adel-
berger et al. 2005b) from the observational studies of the
IGM-galaxies connection around z = 1.9 − 2.6 (Adel-
berger et al. 2003, 2005a; Simcoe et al. 2006). However,
the accurate mass range for such rest-frame UV-selected
galaxies are still unknown. In this paper, we simply as-
sume that the central galaxy is a typical UV-selected
star-forming galaxies, and our simulation volume is a
typical environment for such galaxies. This assumption
will be tested in our future studies. Figure 3 shows the
history of the total star formation rate for the star parti-
cles within r = 5 proper kpc at z = 2.43. The extremely
strong feedback in model ESF terminates star formation
in the system. We confirmed that SNe Ia continuously
heat the ISM, which keeps maintaining an outflow in
model ESF.
We chose the X–Y plane in Figure 2 as the projected
plane of the sky, and define the Z-axis as the LOS di-
rection. Figures 4-6 demonstrate the overdensity, tem-
perature, metallicity, and abundance ratio of carbon to
oxygen, [C/O], in the X–Z plane at three different posi-
tion of Y= ±50 and 0 proper kpc for models NF, SF, and
ESF, respectively. In the same planes, Figures 7-9 give
the density distributions of HI, CIV, and OVI, and the
OVI weighted temperature map. In these figures, the
velocity field of the gas component is also shown with
tangential arrows.
In model NF the central galaxy ends up surrounded by
a hot gaseous halo with a radius of about 100 kpc (Fig.
4). This is due to infalling gas being shock-heated to
the virial temperature (Table 1). This situation appears
common for high redshift galaxies (Rauch et al. 1997)
and corresponds to the hot accretion mode of (Kereš
et al. 2005; Dekel & Birnboim 2006)
Because of the inflow, the metallicity of the hot gas
is low. On the other hand, the high density filaments
are cold, and part of the gas keeps accreting through the
filaments onto the galaxy, i.e., the cold accretion mode
described by Kereš et al. (2005) and Dekel & Birnboim
(2006) (see also the velocity field in Fig. 7). As a result,
the HI density is higher along the filament, and gets sig-
nificantly higher in the collapsed region near the central
Galactic Wind Signatures 9
Fig. 9.— Same as Fig. 7, but for model ESF.
Fig. 10.— Composite image of overdensity (blue contour) and
metallicity (red contour) distribution at Y= 0 for model SF. The
edges of contours for overdenisty and metallicity correspond to
log ρg/ < rhog >= −1 and logZ/Z⊙ = −2.5, respectively.
galaxy and the neighboring galaxies. The spatial dis-
tribution of CIV and OVI also traces the filaments, and
their densities are high in the region close to the galaxies.
Model SF produces a more extended hot gas region
than model NF, and the hot gas is now metal-enriched
(Fig. 5), compared to the NF case. Here the hot gas is
dominated by a galactic wind induced by strong SN feed-
back. The temperature map of Figure 5, the arrows in
Figure 8, and Figure 10 demonstrate that the enriched
gas tends to escape toward the lower density regions.
The filament remains unaffected by the wind. As a re-
sult, the cold accretion is still maintained through the
filaments, which continues funneling gas to the galaxy.
This is similar to what was observed in previous numeri-
cal simulations with strong feedback effect (e.g., Theuns
et al. 2002; Kollmeier et al. 2006). Consequently, the dis-
tribution of HI density is not significantly different from
model NF. On the other hand, higher density CIV and
OVI gas extends over a much larger region. This is be-
cause the enriched gas is blown out from the galaxy, and
helps to raise the abundance of carbon and oxygen in the
Interestingly, we find that the gas whose OVI density is
high in model SF has a low temperature which is consis-
tent with photoionization equilibrium for OVI (typical
logarithmic temperatures are around Log(T (K))= 4.2)
as opposed to collisional ionization temperature (typical
Log(T (K))= 5.5). Comparison between the 3rd and 4th
columns in Figure 8 demonstrates that the region where
the density of OVI is high has temperatures around
Log(T (K))< 4.5. This is gas that has been blown out
of the galaxy, and cooled down by radiative cooling after
colliding with the ambient IGM. CIV in low density halo
10 Kawata and Rauch
Fig. 11.— Mosaic of 5× 5 lines of sight spectra of HI, CIV, and OVI lines around a galaxy at z = 2.43 for model NF. The central panel
corresponds to the position of the galaxy, and each line of sight is separated by 50 kpc in physical scale. The numbers at the upper left
corner of each panel present the (x,y) coordinate (in kpc) for each LOS. In the case of CIV and OVI only the stronger line of each doublet
is shown. The LOS velocity is adjusted so that the LOS velocity of the galaxy equals zero.
Fig. 12.— Same as Fig. 11, but for model SF.
Galactic Wind Signatures 11
Fig. 13.— Same as Fig. 11, but for model ESF.
Fig. 14.— The mean transmissivity of HI (upper), CIV (middle), and OVI (lower) for the pixels within 1 h−1 comoving Mpc of the
galaxies as a function of the impact parameter for models NF (left), SF (middle), and ESF (right). The gray lines indicate median and
25th and 75th percentile. Note that the scale of the y-axis, i.e., f1Mpc is different in each panel.
12 Kawata and Rauch
Fig. 15.— The probability of the flux decrement 1 − f1Mpc of HI (upper), CIV (middle), and OVI (lower), where f1Mpc is the
mean transmissivity for the pixels within 1 h−1 comoving Mpc of the galaxies. The left/middle/right panel shows the results of model
NF/SF/ESF. The dotted histogram in the upper panels present the observational results of Adelberger et al. (2005a).
Fig. 16.— The mean transmissivity of the galaxies of HI, f1Mpc,HI for the pixels within 1 h
−1 comoving Mpc as a function of the mean
transmissivity of OVI, f1Mpc,OVI, for models NF (left), SF (middle), and ESF (right). The lines indicate median (black line) and 25th and
75th percentile (grey lines).
and void regions is in a similar thermal state.
Figure 6 shows that the extremely strong SN feedback
in model ESF develops a much stronger galactic wind
and a larger hot gas bubble. Feedback is now strong
enough to affect the filaments. Gas accretion through
the filaments is suppressed, so that star formation in the
central galaxy ceases (Fig. 3). However, even here the
denser regions of the filaments survive, and the density
distribution of HI is similar to the one seen in models NF
and SF. Figure 9 reveals that in model ESF there is more
collisionally ionized OVI especially in the region close to
the galaxy (r ≤ 40 kpc) where the OVI weighted tem-
perature is around Log(T (K))=5.5 and the OVI density
peaks.
Figures 11-13 represent the spectra whose LOS are cho-
sen as the 5× 5 grid points each separated by 50 proper
kpc projected on the sky i.e., in the X–Y plane. In the
figures, the middle panels correspond to the LOS through
the center of the galaxy. First, we compare HI absorp-
tion lines between the three models. At the LOS through
the central galaxy, the HI lines are heavily saturated,
i.e., they produce a damped Lyman alpha line, except in
model ESF. In model ESF, not enough cold gas can sur-
vive in the galaxy due to the extremely strong feedback.
In all the models, the HI absorption becomes weaker with
the projected distance from the galaxy. General features
of HI absorption lines are similar among the three mod-
els. Thus, the signature of a galactic wind seems to be
difficult to see in HI absorption lines, confirming the con-
clusions of previous studies (Theuns et al. 2002; Croft
et al. 2002; Bruscoli et al. 2003; Kollmeier et al. 2006).
However, a more detailed comparison of HI absorption
line features between models NF and SF (Figs. 11 and
12, respectively) shows that the HI absorption lines tend
Galactic Wind Signatures 13
Fig. 17.— The mean transmissivity of CIV (uppera), and OVI (lower) for the pixels within 1 h−1 comoving Mpc of the galaxies as a
function of the impact parameter. The left panel shows the resul of model SF. The middle panel presents the result of model SF, when the
IGM is assumed to be homogeneously enriched at the level of [C/H]=-2.5 and [O/H]=-3. The right panel demonstrates the result of model
SF, when the QSO only UV background radiation suggested by Haardt & Madau (2001) is adopted. The gray lines indicate median and
25th and 75th percentile.
to be stronger in model SF than those in model NF,
even close to the galaxy. This is a counter-intuitive re-
sult, because it seems natural that a strong wind should
predominantly be destroying HI clouds, as suggested by
Adelberger et al. (2003) and Bertone & White (2006).
However, as shown in Figures 2 and 5, strong feedback
redistributes the high density gas in the galaxy to the
surrounding region, and the filaments become broader.
This leads to the stronger HI lines in model SF, espe-
cially in the outer regions of the galaxy. A similar effect
is seen in model ESF (Fig.13). Therefore, we suggest
that stronger SN feedback actually increases the HI ab-
sorption. We will test this more quantitatively in the
next section.
Conversely, we do see a lack of neutral hydrogen along
some LOS close to the galaxy. However, they do not
seem to have anything to do with winds. In model NF,
the LOS at (X,Y) = (−100,−50) and (−100,−100) have
very weak HI absorption lines, although their projected
distance is smaller than ∼120 proper kpc (∼ 300 h−1
comoving kpc). Adelberger et al. (2005a) studied HI ab-
sorption lines around star-forming galaxies at 2 ≤ z ≤ 3,
and found that some LOS which are within 1 h−1 comov-
ing Mpc from the galaxies show weak or absent HI ab-
sorption. They argue that such a lack of absorption may
be caused by a galactic superwind destroying the neutral
hydrogen. However, our model NF does not include any
SN feedback. The LOS at (X,Y) = (−100,−50) corre-
sponds to the line at X= −100 in the lower panels of
Figures 4 and 7. This LOS passes through hot accretion-
shocked gas which cannot accommodate HI, and misses
the more HI-rich filaments. This example demonstrates
that it is possible to have LOS close to galaxies which do
not show any strong HI absorption, without the need for
a galactic wind.
It is also worth mentioning that there are some LOS
which show double HI absorption lines, especially in
model NF, e.g., the LOS at (X,Y) = (50, 0). We find
that this is due to symmetric gas infall from the fila-
ments. The LOS at (X,Y) = (50, 0) in model NF can be
seen at X= 50 kpc at the middle panels of Figure 7. This
LOS passes through two filaments at Z∼ −120 and ∼ 75
kpc, and the velocity map shows the filament at Z∼ −100
(∼ 75) has positive (negative) LOS velocity. As a result,
these two filaments appear as double absorption compo-
nents. Such double HI absorption line features become
less obvious in the cases of stronger feedback, because
outflow from the galaxy makes the velocity field more
chaotic, and fills the gap between the components in ve-
locity space.
We also compare CIV and OVI lines among the models.
In model NF, there are almost no CIV or OVI lines at
R≥ 50 proper kpc. The LOS at (X,Y) = (50,−50) shows
strong CIV lines. However, this is due to the next closest
galaxy, as seen in Figure 7. In contrast, the stronger
feedback in models SF and ESF creates CIV and OVI
lines further away from the galaxy. We also investigated
the absorption lines with a much finer grid of LOS, i.e.,
smaller separations, and found that strong CIV or OVI
lines are rare at projected radii R ≥ 100 kpc even in
models SF and ESF, unless there is another galaxy close
to the LOS. This can also be seen in Figures 8 and 9,
where dense CIV and OVI regions extent to about 100
kpc. OVI lines are very rare in model NF, and only exist
where the HI absorption is saturated. On the other hand,
strong feedback models produce more OVI lines, that
are sometimes stronger than CIV lines. These figures
demonstrate that the OVI lines are the most sensitive
signature of a galactic wind in absorption. We study
this possibility more quantitatively in the next section.
3.2. The mean transmissivity of HI, CIV, and OVI
In this section, we analyze the artificial QSO spectra
in 1000 random LOS. Since our simulation volume is a
spherical volume (Sec. 2.2), we cannot use the LOS at
too large an impact parameter. Therefore, we generate
1000 spectra for the random LOS at the projected radius
of R < 400 proper kpc. We also change the angle of
projection randomly for each LOS.
Within the three dimensional radius of r = 400 proper
kpc, the high-resolution volume of the simulations con-
tains two galaxies whose virial mass is more than 1011
14 Kawata and Rauch
M⊙. Since the virial mass of the observed UV-selected
galaxies is not well known, as mentioned above (see also
Erb et al. 2006), we assume these two galaxies are such
galaxies, and apply a similar analysis to the one done
for the observed UV-selected galaxies (Adelberger et al.
2003, 2005a). Adelberger et al. (2005a) measured the
mean transmissivity of all HI pixel in their QSO spectra
that lie within 1 h−1 comoving Mpc from the galaxies as
determined from the projected distance and the LOS ve-
locity. In this paper we indicate the mean transmissivity
as f1Mpc. We analyze f1Mpc not only for HI (f1Mpc,HI)
but also for CIV (f1Mpc,CIV) and OVI (f1Mpc,OVI) for all
the LOS spectra for our two galaxies. In reality, it is diffi-
cult to measure f1Mpc for metal lines. For example, OVI
lines are often contaminated by interloper Lyman alpha
lines. However, we carry out this theoretical exercize
to understand the effect of galactic winds on absorption
lines quantitatively.
Figure 14 shows the mean transmissivity of HI, CIV,
and OVI as a function of the projected distance from the
galaxy, i.e., the impact parameter, b. Figure 15 shows the
histogram of the probability of the decrement which is de-
fined as 1−f1Mpc from Figure 14. The top panels of Fig-
ures 14 and 15 correspond to the lower panel of Figures
13 and 15 of Adelberger et al. (2005a), respectively. The
left-top panel of Figure 15 shows results similar to those
from previous numerical simulation studies with without
feedback (Kollmeier et al. 2003; Tasker & Bryan 2006).
Although the observational data (dotted histogram) of
Adelberger et al. (2005a) agree as far as strong absorp-
tion is concerned, their data show a much higher proba-
bility for the very weak absorption (1 − f1Mpc < 0.2).
Adelberger et al. (2005a) claim that this may be be-
cause in the real universe galactic winds turn moder-
ate absorption into weak absorption, but do not affect
the strong absorption systems. However, as seen in the
top panels of Figures 14 and 15, our simulations predict
that the existence of the strong galactic wind does not
change the mean flux transmissivity of the HI lines. In-
terestingly, if we compare the transmissivity at the LOS
close to the galaxy (b ≤ 0.4h−1 comoving Mpc) in Figure
14, the stronger feedback leads to slightly stronger mean
absorption, i.e., smaller median f1Mpc,HI, although the
difference is subtle. Therefore, again, we conclude that
HI absorption lines generally are unaffected by galactic
winds. This leaves an inconsistency between the obser-
vations and numerical simulations. Unfortunately, the
current number of the observational sample is not satis-
factory (31 systems in Adelberger et al. 2005a) to reach
firm conclusions. On the theory side, the implementation
of feedback has been one of the more uncertain ingredi-
ents in current numerical simulations (e.g. Okamoto et al.
2005; Kobayashi et al. 2006; Scannapieco et al. 2006a).
We adopt the simplest implementation, but different im-
plementations may lead to different conclusions. Clearly,
further observational studies and numerical simulations
are required to address this problem.
Figures 14 and 15 also show the results for CIV and
OVI lines, which appear to be more sensitive to the ef-
fect of SN feedback. Model NF shows low f1Mpc,CIV to be
almost independently of the impact parameter. On the
other hand, models SF and ESF show that significantly
more LOS have a lower f1Mpc,CIV at b <∼ 0.2h
−1 co-
moving Mpc. The mean transmissivity of OVI lines also
shows a similar trend. However, for OVI the absorption
becomes noticeably stronger i.e., f1Mpc,OVI decreases, as
the impact parameter decreases below b <∼ 0.4h−1 co-
moving Mpc. Figure 15 reveals that model NF barely
shows a decrement 1 − f1Mpc higher than 0.2 for both
CIV and OVI, but models SF and ESF can produce such
strong absorption lines. However, note that y-axis of the
figure is the logarithm of probability, and it represents
only ∼1 % of f1Mpc that show such strong absorption in
models SF and ESF.
In the previous section, we suggested that OVI is a
good tracer for a galactic wind, and in model NF OVI
lines are only observable where the HI lines are satu-
rated. On the other hand, models SF and ESF produce
OVI lines even where the HI is not saturated. Figure 16
plot the f1Mpc,HI against f1Mpc,OVI. Figure clearly shows
that models SF and ESF show the significant fraction of
the spectra with relatively weak HI (f1Mpc,HI > 0) and
stronger OVI (f1Mpc,OVI < 0.95). We have calculated
the probability of low f1Mpc,OVI for the spectra with
f1Mpc,HI > 0.2. Model NF has only 14 % of the spec-
tra with f1Mpc,OVI < 0.95, while models SF and ESF
have 24 and 28 %. Although the difference is small, the
stronger feedback seems to produce more such HI weak
OVI strong lines.
Finally, we briefly mention how our results are sensi-
tive to the distribution of metals and the assumed UV
background radiation. Since re-simulations changing the
metal yields and the UV background radiation are com-
putationally too expensive, we analyzed the results of
model SF, assuming different metal distributions and UV
background radiations. Figure 17 shows the results of
the same analysis as Figure 14 in the cases when the the
IGM is assumed to be homogeneously enriched at the
level of [C/H]= −2.5 and [O/H]= −3 (model SFhZ) and
when the QSO-only UV background radiation suggested
by Haardt & Madau (2001) is adopted (model SFQ),
while model SF includes radiation from both QSO and
galaxies. Model SFhZ demonstrates that if the heavy el-
ements are homogeneously distributed with the assumed
metallicity, there is very little correlation between the
strength of metal lines and the impact parameter. There-
fore, it is important to follow the metal distribution in the
IGM self-consistently. Model SFQ shows that f1Mpc,CIV
and f1Mpc,OVI differ little between the QSO and galaxy
radiation and QSO-only UV background radiation cases,
except for very subtle decrease in CIV absorptions and
increase in OVI absorptions (see also Aguirre et al. 2005).
4. CONCLUSIONS
We have analyzed the QSO absorption features ob-
tained from cosmological numerical simulations with dif-
ferent strengths of SN feedback. Our simulations self-
consistently follow the metal exchange histories among
the IGM, ISM, and stellar components. We investigate
not only the neutral hydrogen absorption lines but also
the ionization lines for heavy elements, keeping track of
the abundance history of the elements.
We have paid particular attention to the properties of
the IGM around high-redshift (z = 2.43) galaxies with
Mvir ∼ 10
11 M⊙. We found that a model without me-
chanical feedback creates hot gas halos around galaxies
due to shock heating, with radii up to 100 proper kpc.
We found that such hot gas can lead to a lack of HI
Galactic Wind Signatures 15
absorption (Fig. 11) even for LOS close to galaxies, as
found by Adelberger et al. (2005a), without having to
invoke a galactic wind.
In our strong feedback models, outflows induced by
SN feedback produce larger hot bubbles around galax-
ies (Figs. 5 and 6). However, such outflows tends to
escape to lower density regions, and hardly affect the
dense filaments producing HI absorption systems so that
the transmissivity of HI Lyman alpha is virtually inde-
pendent of the strength of SN feedback. If anything the
absorption by neutral hydrogen slightly increases in the
presence of a wind. We conclude that the presence or
absence of HI absorption lines is not a good indicator of
the presence or absence of a galactic wind.
On the other hand, we found that the metal lines, es-
pecially OVI, are sensitive to the existence of outflows.
Without feedback, it is difficult to enrich the IGM enough
to produce strong OVI lines further away from galaxies
(Fig. 4), unless there are nearby satellite galaxies inter-
sected by the LOS by chance. We also found that, in
the no-feedback model, strong OVI lines are almost al-
ways associated with saturated HI lines. On the other
hand, the strong feedback model can produce strong OVI
lines even where HI lines are unsaturated, because strong
feedback can re-distribute the enriched gas to relatively
low density regions. We have confirmed this by looking
for the spectra whose OVI flux is less than 0.8 over more
than 5 pixels and whose mean HI flux within ±50 km
s−1 from the HI velocity corresponding to the OVI lines
are higher than 0.2 from 1000 spectra with random LOS.
The no-feedback model has no such spectra, while strong
feedback (model SF) has 12 of such spectra. We point
out that Figure 9 of Simcoe et al. (2004) shows an OVI
line where the HI is not saturated. Our results suggest
that this is likely to be a region where the effect of a
galactic wind is significant. Analyzing the transmissivity
of OVI lines we found that strong feedback creates more
LOS with lower transmissivity, i.e. stronger OVI ab-
sorption, near the star-forming galaxies. The statistical
analysis of transmissivity also shows that there are more
LOS where stronger OVI is associated with weaker HI,
in the presence of galactic winds. We expect that the
pixel-optical depth analysis of OVI against HI (Schaye
et al. 2000) would be sensitive to the presence of a galac-
tic wind, and we will test this idea in a future paper. In
conclusion, OVI appears a theoretically good tracer of
galactic winds that merits further attention.
DK thanks the financial support of the JSPS, through
Postdoctoral Fellowship for research abroad. We ac-
knowledge the Center for Computational Astrophysics of
the National Astronomical Observatory, Japan (project
ID: imn33a), the Institute of Space and Astronautical
Science of Japan Aerospace Exploration Agency, and
the Australian and Victorian Partnerships for Advanced
Computing, where the numerical computations for this
paper were performed. MR is grateful to the NSF for
support under grant AST-05-06845.
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