Issue |
A&A
Volume 496, Number 1, March II 2009
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|
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Page(s) | 51 - 56 | |
Section | Extragalactic astronomy | |
DOI | https://doi.org/10.1051/0004-6361/200810760 | |
Published online | 30 January 2009 |
A giant bar induced by a merger event at z = 0.4?
S. Peirani1 - F. Hammer2 - H. Flores2 - Y. Yang2 - E. Athanassoula3
1 - Institut d'Astrophysique de Paris, 98 bis Bd Arago, 75014 Paris, France
2 - GEPI, Observatoire de Paris, CNRS, University Paris Diderot, 5 place Jules Janssen,
92190 Meudon, France
3 - LAM, UMR6110, CNRS/Université de Provence, Technopôle de Marseille-Etoile, 38 rue Frédéric Joliot Curie, 13388 Marseille Cedex 20, France
Received 6 August 2008 / Accepted 3 December 2008
Abstract
Context. Disk galaxies are the most common galaxy population in the local universe. However, the formation of their disks and structures - in particular their bars - is still a matter of debate.
Aims. We present a physical model of the formation of J033239.72-275154.7, a galaxy observed at z=0.41 that contains a large, young bar of size 6 kpc. The study of this system is particularly interesting to the understanding of the connection between mergers and bars, as well as the properties and fate of this system in relation to disk galaxy formation.
Methods. We compare the morphological and kinematic properties of J033239.72-275154.7, the latter obtained by the GIRAFFE spectrograph, to those derived from the merger of two spiral galaxies described by idealized N-body simulations, including a star formation prescription.
Results. We found that the general morphological shape and most of the dynamical properties of the object can be well reproduced by a model in which the satellite is initially placed in a retrograde orbit and the mass ratio of the system is 1:3. In this scenario, a bar forms in the host galaxy after the first passage of the satellite, where a significant fraction of the available gas is consumed in an induced burst. In its later evolution, however, we find that J033239.72-275154.7, whose major progenitor was an Sab galaxy, will probably become a S0 galaxy. This is mainly due to the violent relaxation and the angular momentum loss experienced by the host galaxy during the merger process, which is caused by the adopted orbital parameters. This result suggests that the building of the Hubble sequence is influenced significantly by the last major collision. In the present case, the merger leads to a severe damage of the disk of the progenitor, leading to an evolution towards a more bulge-dominated galaxy.
Key words: galaxies: evolution - galaxies: kinematics and dynamics - galaxies: interactions - methods: N-body simulations
1 Introduction
The formation of disk galaxies remains an outstanding puzzle in contemporary astrophysics (see Mayer et al. 2008, for a review). According to hierarchical models of structure formation, mergers and interaction of galaxies are an essential ingredient of galaxy formation and evolution. Earlier works and numerical simulations demonstrated that the remnants of mergers of entirely stellar progenitors are more likely to be elliptical galaxies (Toomre 1977; Barnes 1988; Barnes & Hernquist 1992; Hernquist 1992; Lima-Neto & Combes 1995; Balcells & González 1998; Naab et al. 1999) and later studies extended this result to gas-rich progenitors (Springel & Hernquist 2005; Robertson et al. 2006; Hopkins et al. 2008). Whether this is consistent, or inconsistent with the fraction of disk galaxies present in the local universe depends on the typical number of expected mergers per galaxy (see for instance Kazantzidis et al. 2007). These predictions appear to be inconsistent with observations that suggest that disk galaxies represent the majority (70%) of the galaxy population observed in the local universe (see Hammer et al. 2005; Nakamura et al. 2004, and references therein). To help resolve these issues, Hammer et al. (2005) suggested that disks could be rebuilt during encounters of gas-rich spirals. This proposition was inspired by the simultaneous increase, up to z=1, in the merger rate, in the fraction of actively star-forming galaxies (including luminous IR galaxies, LIRGs) and in the fraction of peculiar galaxies (including those with compact morphologies). These studies were followed by simulations of gas-rich mergers (Springel & Hernquist 2005; Robertson et al. 2006; Governato et al. 2007; Hopkins et al. 2008), which demonstrating that under certain conditions a disk may be rebuilt after a merger. This conclusion had been already reached by Barnes (2002), although these simulations have been less convincing because they do not include a prescription for star formation. Lotz et al. (2006) also analysed a large suite of simulated equal-mass, gas-rich mergers and found that most merger remnants appear disc-like and dusty. This scenario is also consistent with the results of other simulations and semi-analytical models which claim that, without merger processes, most of galaxies and their host dark matter halos cannot acquire the required angular momentum to form disks (see Peirani et al. 2004; Puech et al. 2007, and references therein). The formation of bars is also a fundamental issue in the evolution of disk galaxies, particularly since it has been shown that only about one third of them were in place at z = 0.8 (Sheth et al. 2008).
Our aim is to reproduce by
numerical modelling the general morphology (e.g. the presence
of a bar and substructures),
the dynamical properties of the gas component (velocity fields),
the photometric properties of the stars (e.g. colors, star formation
rate) of J033239.72-275154.7, a galaxy located at z=0.41 for which
we have high quality imaging and kinematical data. This work will also provide
useful input to disk formation models, since it provides information
about the
potential progenitors of the present-day galaxy disks, as well as
constraints on their formation
(e.g. initial orbital configuration and mass ratio of the system)
and possible fate. This work is a subsidiary project of a VLT large program entitled IMAGES (``Intermediate-MAss Galaxy Evolution Sequence'', Yang et al. 2008), which
is gathering high quality kinematics for a representative sample of 100 massive galaxies at
z=0.4-0.75 and with
.
Using the GIRAFFE
spectrograph at the VLT, the kinematic properties of 65 of these
galaxies, for instance J033239.72-275154.7, have been derived.
This galaxy lies at z=0.41, is
classified as a merger from an analysis lead by Neichel et al. (2008),
and exhibits a large, young bar. This bar has a size of 6 kpc, i.e. quite large,
since 72% of barred galaxies in a sample of 2000 galaxies from the SDSS have a size smaller than this (Barazza et al. 2008). This bar also has an extremely blue color, that is consistent
with a starburst, i.e. with ages well below a few 100s of Myr. This is contrasting to many bars in the local universe known to contain relatively old and red stars (Sheth et al. 2003). This paper is organized as follows: in Sect. 2, we present the general morphological and kinetic properties of J033239.72-275154.7, and in Sect. 3, a short description of
the numerical modelling and the results obtained. We present our conclusions
in Sect. 4.
2 General properties of J033239.72-275154.7
Table 1: Properties of J033239.72-275154.7.
J033239.72-275154.7 was first presented by Yang et al. (2008).
Table 1 summarizes its overall properties, including
photometry, morphological parameters, and kinematical measurements
(provided by the IMAGES data base that can be retrieved from
Yang et al. 2008; Neichel et al. 2008; Puech et al. 2008).
This object has a stellar mass of
,
a
K-band magnitude of
MK=-20.94,
and displays a peculiar morphology and
kinematics. Its center is dominated
by an elongated structure, most likely a giant, thin bar of semi-major
axis 0.85 arcsec or about 6 kpc. Neichel et al. (2008) found that its color is approximately b-z=0.8, i.e. typical of a starburst (see their Figs. 8 and 12).
This bar is embedded within a diffuse region, which is probably a disk,
due e.g. to features that resemble tidal arms (one blue
at the upper right on the sky and another on the left), although very
irregular, with several blue clumps. Closed to its center, there is a
significant concentration of light that we sometimes refer to as the
core, which is probably a small bulge
Its color is roughly b-z=1.8, as found for galaxies of
late type to Sbc (Neichel et al. 2008).
The bar is asymmetric, extending spatially far more towards the
upper-right side of the core.
On the other side of the core, the bar is redder and part of it disappears
at UV wavelengths.
The discrepancy in color between the two sides is roughly
0.3 mag, as shown by the CDFS-GOODS observations
(b-band observations, thus UV rest-frame at z=0.4).
At the bottom of the galaxy, there are two bright
adjacent knots, which dominate the rest-frame UV light. These knots have
the color of a pure starburst and could be hardly resolved. Due to the
irregularity of the disk and the existence of bright, tidal
arms, it is difficult to measure its
position and inclination angle accurately. From the shape of the outer isophotes,
we find that the photometric major axis is roughly
in the upper-right to lower-left direction and the inclination should be
between
and
.
The star formation rate (SFR) is moderate (
,
)
and the dust does not affect its overall color significantly.
Note that SFRs were estimated using the 2800A and 15 micron luminosities,
respectively. Those were converted to SFR using the method described in Kennicutt (1998). Luminosities have been
estimated using interpolated relations between observed photometric points provided by HST/ACS,
EIS and Spitzer/MIPS, using the method described in Hammer et al. (2001, 2005). The ratio of the
to the
is consistent with Av=0.22 (standard extinction curve), which is low.
Finally, the electron density is modest, including at the blue knots.
![]() |
Figure 1:
B-V-z color map of J033239.72-275154.7 from HST/ACS ( upper left panel) and distribution of the galaxy within the GIRAFFE grid ( upper right panel). The lower panels show the velocity field map ( lower left panel) and the associated |
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The velocity field (VF) is obviously complex (see Fig. 1).
The kinematical major axis is almost parallel to the bar and offset by more than one GIRAFFE pixel from the bulge, towards the prominent blue knots (in the bottom). The pixel scale is 0.52 arcsec and the full width half maximum spectral resolution is 23 km s-1. Details of the GIRAFFE instruments can be found in Flores et al. (2006) and Yang et al. (2008).
The small VF
amplitude is dominated by the two star-bursting clumps and by the blue
arm/giant tidal tail. Even so, the velocity
amplitude is relatively small, of the order of 30 km s-1. The velocity
dispersion map is almost featureless at a value of 30-40 km s-1 over most
of the galaxy. Nevertheless, it has a clear, localised
peak
(50 km s-1), which is offset by 2 GIRAFFE pixels from the bulge and could coincide
(to within half a GIRAFFE pixel) with a small and relatively blue clump at the
upper-right located roughly where the blue, tidal arm joins the disk.
The signal-to-noise ratio in this pixel is quite high (
50).
All evidence supports the conclusion that strong interaction and/or merging is at
work in this system. For example, the dynamical axis is off-center and
passes through the region where the two bright knots lie.
There are strong tidal arms, and
the value of
/
is small (about 1-2 depending on the
exact value of the inclination). The velocity dispersion is
about 30-40 km s-1 over most of the galaxy. Given the VF values and the
measured inclination, the rotational velocity is far
below the value expected from the tully-fisher (TF) relation, which is
km s-1. This value was taken from Hammer et al. (2007),
who carefully estimated the various samples used to derive this
relationship. Given the spatial resolution, it is not easy to derive a rotation curve. Nevertheless, thorough modelling of the velocity fields was completed by Puech et al. (2008), who investigated
the redshift evolution of
the Tully-Fisher relation.
Given its kinematical and morphological properties, it is thus unlikely
that J033239.72-275154.7 can be a rotating disk hosting two giant HII regions because: (1) the two knots have a z-band luminosity corresponding to one third of the total z luminosity
(the relative photometry was completed using apertures of diameter equal to the
full width half maximum of the knot luminosity, using the IRAF/DAOPHOT/ package);
(2) the dynamical axis is strongly off-center; (3) the value of
is high and that of
is low.
Could merging of the two bright knots with J033239.72-275154.7 be responsible for producing a giant bar with a relatively blue color? N-body simulations have shown that a small companion merging with a disk galaxy could, depending on its density and its orbit, either destroy a pre-existing bar (Athanassoula 1999; Berentzen et al. 2003), or trigger its formation (Walker et al. 1996; Berentzen et al. 2004).
![]() |
Figure 2:
Time evolution of the projected star number density. The light blue arrows indicate the specific rotation of each galaxy, while the yellow dashed lines show the motion of the satellite. In panel 6, we superposed the GIRAFFE grid. Each frame is 40 kpc |
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3 Simulations
We briefly describe the numerical methodology used to
model J033239.72-275154.7. We use idealized N-body simulations of
the merger of two spiral galaxies. One of the most difficult part of this work results from the
huge number of free parameters in the initial conditions, for instance
related to the mass ratio and the orbital configuration of the system, or the star formation prescription. We performed more than 100 simulations and present our
fiducial modelling. First, as mentioned in Sect. 2, the
dynamical axis of the system is off-center and
the velocity field strongly indicates that the
satellite is moving toward the main object. By trying
different orbital configurations, we found that retrograde orbits
for the satellite are more consistent with the orientation of the tidal
arms of the main galaxy (in the left and upper right parts)
and the specific position of the accreted satellite. Second,
different mass ratios of the system, different inclinations between the
two orbit planes, and different pericentric distances
were tested at the same time in attempting to reproduce the amplitudes of the
observed VF and map. For instance, since the VF amplitudes are
small, the inclination between the
two orbital planes should also be small.
It is also worth mentioning that we restricted our study
to parabolic orbits (Khochfar & Burkert 2006) to reduce
the number of free parameters.
Finally, different parameters related to the star formation
recipes (see below) have also been tested in order to obtain results consistent
with the observations.
3.1 Initial conditions and numerical method
Our galaxies consist of a spherical dark matter halo (with a
Hernquist profile, Hernquist 1990), a disk, containing both stars and gas,
and a bulge. In both objects, the disk and
the bulge represent respectively 15% and 5% of
the total mass. The baryon fraction used here is slightly higher than the
cosmic baryon fraction (16.7%) derived by Komatsu et al. (2008) in
order to compare to and match the results of Barnes (2002).
We note that the gas fraction of the disk is 25% in the
host galaxy, whereas it is 50% in the accreted satellite.
These gas fractions are consistent with the estimations of Liang et al. (2006), Gavazzi et al. (2008), and Rodrigues et al. (2008). However, the Gavazzi et al. (2008)
measurements were completed for the local universe, while
using the M-Z relationship, Liang et al. (2006) were able to
estimate the mass evolution in the gas phase
of z=0.6 galaxies, and found that,
on average, galaxies at z=0.6 have a factor of two more gas than at the present
day. The host galaxy has a stellar mass of
and
is most likely to be a Sab galaxy. Galaxies are created following
Springel et al. (2005). Dark matter haloes have a concentration parameter of
and
for the host and the satellite, respectively, in good agreement with Dolag et al. (2004).
Their spin
parameter is defined to be
,
where J is the angular momentum, E is the
total energy of the halo, and M is its mass, and
equals
to ensure that
the rotation curves of the galaxies are closed to the baryonic
TF relation (McGaugh 2005). Indeed, their
maximum circular velocity is 140 km s-1 and 98 km s-1, respectively.
In our fiducial model, the satellite is placed on a retrograde, parabolic orbit (Khochfar & Burkert 2006) with a pericentric distance
kpc and initial separation of 20 kpc. The inclination between the two orbital planes is
and
the galaxy spins are opposed each other. We note that the initial separation between the two galaxies is quite small in order to reduce the high computational cost of all
experiments. A larger distance would decrease
the halos overlap at the beginning of the simulation, but this should not
change the main results and conclusions of this paper.
The simulation is performed using GADGET2 (Springel 2005) with added prescriptions for cooling, star
formation, and feedback from type Ia and II supernovae (SN). Approximately 275 000 particles are used for the experiment, and the masses (M), gravitational softening lengths (), and
number of particles (N) of each component involved are summarized in Table 2.
Table 2:
Masses (M), gravitational softening lengths ()
and
number of particles (N) of each component used in the simulations.
The cooling and star formation (SF) recipes follow the
prescriptions of Thomas & Couchman (1992) and Katz et al. (1996),
respectively. Gas particles with T>104 K cool at constant density for the duration of a timestep. Gas particles with
K, number density n > 0.1 cm-3,
overdensity
,
and
form stars according to the standard SFR prescription:
,
where
refers to
the stellar density,
is the dynamical timescale of the gas, and c* is the SF efficiency. Assuming a constant, dynamical time scale across the timestep, the fractional change in
stellar density is
.
For each gas particle, we select a random number (r) between 0 and 1 and convert it into a star if
.
Instead of assuming ``instantaneous'' energy injection, we include the effective
lifetime of SN progenitors using the rate of energy injection
.
For this, we consider stellar lifetimes in the mass
ranges of
and
for type Ia and type II progenitors,
respectively. Using a Salpeter initial mass function for type II
SN, we find that:
![]() |
(1) |
where



![]() |
(2) |
Equations (1) and (2) are used to compute the energy released by SN derived from a star particle i (Ei). A fraction





3.2 Results
Figure 2 illustrates the time evolution of the
projected stellar number density of the system. After 0.36 Gyr
from the beginning of the simulation,
the system has a stellar mass of
and exhibits, from a morphological point of view,
a general shape and pattern similar to the observed
galaxy system. The disc of the more massive galaxy has not yet been destroyed,
reflecting the early stage of the dynamical process,
and is viewed in Figs. 2 and 3 at an
inclination of
,
as inferred from its outer isophote. The bar of the host galaxy has roughly the correct orientation and length (
5 kpc), and the projected position of
the satellite remnant is in approximate agreement with these values. However, many
difficulties remain e.g. in reproducing the blue arm or tidal tail
on the upper right and the fact that the companion should be split
into two distinct, bright knots.
In the present scenario, the bar forms after the first encounter with the satellite, since at that time there is a strong triggering. In good agreement with the observations, it is
particularly clear in the newly formed stars, since those, being on near-circular orbits, are
more prone to the bar instability. Also in good agreement is the fact that
the newly formed stars are also located in the satellite remnant (Fig. 3). Moreover, according to our model, J033239.72-275154.7 is observed when the satellite is about to
experience a the second encounter with the host. At this specific time, the SFR is
(Fig. 3) in good agreement with the observational estimation.
![]() |
Figure 3: The star formation rate obtained from our model ( left panel) and projected distribution of newly formed stars at t=0.36 Gyr ( right panel). The arrows in the left panel correspond to the SFR at the first encounter (1), at t=0.36 Gyr (2) and at the final plunge (3). |
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The VF and
map of the gas component derived from our numerical model are
represented in the GIRAFFE format in Fig. 4. The determination of the radial velocities is presented at the orientation of Figs. 2 and 3.
The observed VF and its amplitude are faithfully reproduced and show that
the initial inclination between the two orbital planes is consistent.
The
map is relatively flat around a mean value of 35-40 km s-1, in good agreement with the observations. It exhibits two maxima of roughly the correct value, but located at positions other than those observed. However, these two maxima are located in regions that are not well resolved in gas particles and thus the resulting pixel values are insignificant.
After 4.26 Gyr, corresponding to the present time z=0, the final object is
characterized by a bulge-to-disk ratio close to 1.04,
as derived from the mass-profile decomposition along the major axis
of the final product (see Fig. 5).
For this decomposition, we used standard models for the disk and bulge,
namely an exponential disk (Sérsic profile with index n=1) and a bulge following
de Vaucouleurs law (Sérsic profile with index n=4),
in which the free parameters are the bulge and disk radius and
the bulge and disk flux. The rotational support parameter
,
where
is the maximum of the major axis rotation speed and
the central velocity dispersion. This last parameter is estimated by considering all star particles within a sphere of
radius 0.5 kpc. These previous derived values are strongly indicative of the formation of a
S0 galaxy being observed.
![]() |
Figure 4:
Maps of the velocity field ( on the left) and |
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![]() |
Figure 5: Mass profile decomposition along the major axis of the final product. The mass profile of the final product is show by the black line. The green dotted-dashed line, the blue dashed line and the red dotted line indicate the disk component, the bulge component and the sum of them, respectively. |
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4 Discussion and conclusions
In this paper, we have presented photometric and kinematic data of
J033239.72-275154.7 and completed a detailed comparison with an N-body
simulation. A first interesting and important result is that among all the simulations
that we have realized with specific
orbital parameters and progenitor mass ratios, the combination of
parameters such as the star formation history, VF, and
map place strong
constraints on the models
and can remove the degeneracy in the good candidates based on morphological criteria only.
We found that our fiducial model can approximately reproduce the general
morphological shape and the total stellar mass of
the object, as well as observed SF rate, VF, and intensity
of the
map.
This specifically includes the reproduction of the giant, young bar,
its location and shape, the relative projected location of the
companion, the
overall morphological shape of the galaxy, its low rotation,
and its off-center dynamical axis. The model predicts a deceleration in
the velocity of the host galaxy due to exchanges of angular
momentum with the companion (see Fig. 2).
However, using these idealized simulations, we have been unable to reproduce
many features, such as the blue arm or tidal tail in the upper right
part of the observed galaxy system, the
morphological patterns of the two bright knots of the satellite remnant,
and the location of the
peaks. These discrepancies with
observations can be easily understood as being due to the huge
parameter space covered by the simulations. These include the
properties of the progenitors (halo: disk: bulge: gas mass ratios and
relative extents, shape of their density profiles, and kinematics),
the geometry of the encounter and of its viewing angle, the gas properties,
and the SF, cooling, and feedback modeling.
On the other hand, the discrepancies with observations may also have
an astrophysical origin and/or be due to a more complex formation
scenario. Several possibilities - such as multiple encounters or
progenitors with specific properties - come to mind, but would
increase substantially the already vast free-parameter space.
In spite of these difficulties, it is encouraging that our
simple numerical modelling is able to build a consistent picture of
the formation of J033239.72-275154.7.
This system is consistent with having been formed
from a merger of two objects with a mass ratio 1:3. The simulation indicates
that a bar is forming in the host galaxy after the first passage of the satellite
where an important fraction of available gas is consumed in the induced
burst, giving a plausible explanation of the observed blue colors of
the bar and the satellite remnant.
Moreover, both the VF and
map derived from the simulations are
able to reproduce the
observational values and thus support this scenario.
In its future evolution, we inferred that J033239.72-275154.7
may become a S0 galaxy, as suggested by the results of our simulations.
This is mainly explained by the fact that the host galaxy
experiences a violent relaxation and
loses
angular momentum during the merger process due to the retrograde orbit
of the satellite. Moreover,
by losing angular momentum, some of the gas may also sink toward the
center of the galaxy where it may be converted into new stars and then
accelerate the growth of the bulge.
To finish, we note that recent numerical studies have demonstrated that gas-rich mergers can produce remnant disks, provided that strong feedback processes are considered and if both the stellar and gas components do not experience a significant angular momentum loss (Springel & Hernquist 2005; Robertson et al. 2006; Governato et al. 2007; Hopkins et al. 2008). In this work, we propose an extreme scenario in which: i) the initial orbital configuration of the merger event does not permit the latter criteria to be satisfied; and ii) the progenitors of the system (in our numerical model) may not be sufficiently gas rich to reform a significant disk according those past theoretical studies. However, the present results are consistent with previous studies in highlighting the fundamental role played by the last major event in building the Hubble sequence (see e.g. Hammer et al. 2005 and 2007). It would thus be interesting to determine the frequency of various orbital configurations during merger events, and to compare our theoretical predictions for simulated galaxies at redshift z=0 with the galaxy population in the local universe to confirm this galaxy formation scenario.
Acknowledgements
S. P. acknowledges the financial support through a ANR grant. E. A. from ANR-06-BLAN-0172. We warmly thank the referee Brant Robertson for an insightful review that considerably improved the quality of the original manuscript. S. P. also warmly thank Y. Kakazu for interesting conversations and for her moral support.
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All Tables
Table 1: Properties of J033239.72-275154.7.
Table 2:
Masses (M), gravitational softening lengths ()
and
number of particles (N) of each component used in the simulations.
All Figures
![]() |
Figure 1:
B-V-z color map of J033239.72-275154.7 from HST/ACS ( upper left panel) and distribution of the galaxy within the GIRAFFE grid ( upper right panel). The lower panels show the velocity field map ( lower left panel) and the associated |
Open with DEXTER | |
In the text |
![]() |
Figure 2:
Time evolution of the projected star number density. The light blue arrows indicate the specific rotation of each galaxy, while the yellow dashed lines show the motion of the satellite. In panel 6, we superposed the GIRAFFE grid. Each frame is 40 kpc |
Open with DEXTER | |
In the text |
![]() |
Figure 3: The star formation rate obtained from our model ( left panel) and projected distribution of newly formed stars at t=0.36 Gyr ( right panel). The arrows in the left panel correspond to the SFR at the first encounter (1), at t=0.36 Gyr (2) and at the final plunge (3). |
Open with DEXTER | |
In the text |
![]() |
Figure 4:
Maps of the velocity field ( on the left) and |
Open with DEXTER | |
In the text |
![]() |
Figure 5: Mass profile decomposition along the major axis of the final product. The mass profile of the final product is show by the black line. The green dotted-dashed line, the blue dashed line and the red dotted line indicate the disk component, the bulge component and the sum of them, respectively. |
Open with DEXTER | |
In the text |
Copyright ESO 2009
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