Smooth and Starburst Tidal Tails in the GEMS and GOODS
Fields
Debra Meloy Elmegreen
Vassar College, Dept. of Physics & Astronomy, Box 745, Poughkeepsie, NY 12604;
elmegreen@vassar.edu
Bruce G. Elmegreen
IBM Research Division, T.J. Watson Research Center, P.O. Box 218, Yorktown Heights,
NY 10598, bge@watson.ibm.com
Thomas Ferguson
Vassar College, Dept. of Physics & Astronomy, Box 745, Poughkeepsie, NY 12604;
thferguson@vassar.edu
Brendan Mullan
Vassar College, Dept. of Physics & Astronomy, Box 745, Poughkeepsie, NY 12604 and
Colgate University, Dept. of Astronomy, Hamilton, NY; bmullan@mail.colgate.edu
ABSTRACT
GEMS and GOODS fields were examined to z ∼1.4 for galaxy interactions
and mergers. The basic morphologies are familiar: antennae with long tidal tails,
tidal dwarfs, and merged cores; M51-type galaxies with disk spirals and tidal
arm companions; early-type galaxies with diffuse plumes; equal-mass grazing-
collisions; and thick J-shaped tails beaded with star formation and double cores.
One type is not common locally and is apparently a loose assemblage of smaller
galaxies. Photometric measurements were made of the tails and clumps, and
physical sizes were determined assuming photometric redshifts. Antennae tails
are a factor of ∼ 3 smaller in GEMS and GOODS systems compared to local
antennae; their disks are a factor of ∼ 2 smaller than locally. Collisions among
early type galaxies generally show no fine structure in their tails, indicating that
stellar debris is usually not unstable. One exception has a 5×109 M⊙ smooth red
clump that could be a pure stellar condensation. Most tidal dwarfs are blue and
probably form by gravitational instabilities in the gas. One tidal dwarf looks like
it existed previously and was incorporated into the arm tip by tidal forces. The
star-forming regions in tidal arms are 10 to 1000 times more massive than star
complexes in local galaxies, although their separations are about the same. If
http://arxiv.org/abs/0704.0911v1
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they all form by gravitational instabilities, then the gaseous velocity dispersions
in interacting galaxies have to be larger than in local galaxies by a factor of ∼ 5
or more; the gas column densities have to be larger by the square of this factor.
Subject headings: galaxies: formation — galaxies: merger — galaxies: high-
redshift
1. Introduction
Galaxy interactions and mergers are observed at all redshifts and play a key role in
galaxy evolution. Two percent of local galaxies are interacting or merging (Athanassoula &
Bosma 1985; Patton et al. 1997), and this fraction is larger at high redshift (e.g., Abraham
et al. 1996b; Neuschaefer et al. 1997 ; Conselice et al. 2003; Lavery et al. 2004; Straughn
et al. 2006; Lotz et al. 2006, and others). Conselice (2006a) estimates that massive galaxies
have undergone about 4 major mergers by redshift 1. Toomre (1977) described a sequence
of merger activity ranging from separated galaxies with tails and a bridge between them,
to double nuclei in a common envelope with tails, to merged nuclei with tails. Ground-
based (Hibbard & van Gorkom 1996) and space-based (Laine et al. 2003; Smith et al. 2007)
observations of this sequence show optical, infrared, and radio activity in the tails and nuclei.
High resolution images and numerical simulations of nearby interactions demonstrate
how star formation and morphology are affected. General reviews of interaction simulations
are given by Barnes & Hernquist (1992) and Struck (1999). The initial galaxy properties,
such as mass, rotational velocity, gas content and dark matter content, and their initial sep-
arations and velocity vectors, all play a role in generating structure. The viewing angle also
affects the morphology. Early-type galaxies with little gas are expected to display smooth
plumes and shells, while spiral interactions and mergers should exhibit clumpy star forma-
tion along tidal tails, and condensations of material at the tail ends. Equal mass companions
may show bridges between them. A prominent example of a tidal interaction is the Antennae
(NGC4038/9), a merging pair of disk galaxies with rampant star formation in the central
regions, including young globular clusters (Whitmore et al. 2005). Its interaction was first
modeled by Toomre & Toomre (1972). The Cartwheel galaxy is a collisional ring system
rimmed with star formation from a head-on collision (Struck et al. 1996). Sometimes polar-
ring or spindle galaxies are the result of perpendicular collisions (Struck 1999). The Mice
(NGC 4676) has a long narrow straight tail and a curved tidal arm (Vorontsov-Velyaminov
1957; Burbidge & Burbidge 1959); numerical simulations reproduce both features well in a
model with a halo:(disk+bulge) mass ratio of 5 (Barnes 2004). The Superantennae (IRAS
19254-7245) is a pair of infrared-luminous merging giant galaxies having Seyfert and star-
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burst nuclei and ∼ 200 kpc tails with a tidal tail dwarf (Mirabel, Lutz, & Maza, 1991). The
Leo Triplet includes NGC 3628 with an 80 kpc stellar tail containing star-forming complexes
with masses up to 106 M⊙ (Chromey et al. 1998). The Tadpole galaxy UGC10214 (Tran et
al. 2003; de Grijs et al. 2003; Jarrett et al. 2006), the IC2163/NGC2207 pair (Elmegreen
et al. 2001, 2006), and Arp 107 (Smith et al. 2005) are all interacting systems observed
with HST and SST and modeled in simulations. Many local mergers have intense nuclear
activity, such as the Seyfert galaxy NGC 5548, which also has an 80 kpc long, low surface
brightness (V=27-28 mag arcsec−2) tidal tail and a 1-arm diffuse spiral (Tyson et al. 1998).
The GEMS (Galaxy Evolution from Morphology and SEDs; Rix et al. 2004), GOODS
(Great Observatories Origins Deep Survey; Giavalisco et al. 2004), and UDF (Ultra Deep
Field; Beckwith et al. 2006) surveys done with the HST ACS (Hubble Space Telescope
Advanced Camera for Surveys) have enabled high resolution studies of the morphology of
intermediate and high redshift galaxies. Light distribution parameters such as the Gini co-
efficient (Lotz et al. 2006) and concentration index, asymmetry, and clumpiness (CAS;
Conselice 2006) have been applied to galaxies in these fields to study possible merger
systems. For GEMS and GOODS, John Caldwell of the GEMS team has posted images
(archive.stsci.edu/prepds/gems/datalist.html) of several galaxies from each field, including
peculiar and interacting systems with tails and bridges. Here we examine the entire GEMS
and GOODS fields systematically for such galaxies and study their tails, bridges, and star-
forming regions. Their properties are useful for understanding interactions and interaction-
triggered star formation, and for probing the relative dark matter content (e.g., Dubinski,
Mihos, & Hernquist 1999).
2. The Sample of Interactions and Mergers
The GOODS and GEMS images from the public archive were used for this study. They
include exposures in 4 filters for GOODS: F435W (B435), F606W (V606), F775W (i775), and
F850LP (z850); and 2 filters (V606 and z850) for GEMS. The public images were drizzled to
produce final archival images with a scale of 0.03 arcsec per px. GEMS, which incorporates
the southern GOODS survey (Chandra Deep Field South, CDF-S) in the central quarter of
its field, covers 28 arcmin x 28 arcmin; there are 63 GEMS and 18 GOODS images that
make up the whole field. The GOODS images have a limiting AB mag of V606= 27.5 for
an extended object, or about two mags fainter than the GEMS images. There are over
25,000 galaxies catalogued in the COMBO-17 survey (Classifying Objects by Medium-Band
Observations, a spectrophotometric 17-filter survey; Wolf et al. 2003), and 8565 that are
cross-correlated with the GEMS survey (Caldwell et al. 2005).
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Interacting galaxies with tails, bridges, diffuse plumes and other features were identified
by eye on the online Skywalker images and examined on high resolution V606 fits images.
The lower limit to the length of detectable tails is about 20 pixels. Snapshots of several
different morphologies for interacting galaxies are shown in Figures 1-6. Out of an initial
list of about 300 galaxies, a total of 100 best cases are included in our sample: 14 diffuse
types, 18 antennae types, 22 M51 types, 19 shrimp types, 15 equal mass interactions, and
12 assemblies, as we describe below.
GEMS and GOODS galaxy redshifts were obtained from the COMBO-17 list (Wolf et
al. 2003). Our sample ranges from redshift z = 0.1 to 1.4 in an area of 2.8×106 square arcsec.
The linear diameters of the central objects were determined from their angular diameters
and redshifts using the appropriate conversion for a ΛCDM cosmology (Carroll et al., 1992;
Spergel et al., 2003). The range is ∼ 3 to 33 kpc. Projected tail lengths were measured in
a straight line from the galaxy center to the 2σ noise limit (25.0 mag arcsec−2) in the outer
tail.
Photometry was done on the whole galaxies, on each prominent star-forming clump, and
on the tails using the IRAF task imexam. A box of variable size was defined around each
feature; the outer limits of the boxes were chosen to be where the clump brightness is about
3 times the surrounding region. Sky subtraction was not done because the background
is negligible. The photometric errors are ∼0.1-0.2 mag for individual clumps. The V606
surface brightnesses of the tidal tails were determined using imagej (Rasband 1997) to trace
freehand contours around the tails, so that they could be better defined than with rectangular
or circular apertures.
Figure 1 shows galaxies with diffuse plumes and either no blue star formation patches
or only a few tiny patches (e.g., galaxies number 5 and 6); we refer to these interactions as
diffuse types. The colors of the plumes match the colors of the outer parts of the central
galaxies, indicating the plumes are tidally shorn stars with little gas. There is structure in
most of the plumes consisting of arcs or sharp edges. This is presumably tidal debris from
early type galaxies with little or no gas (e.g. Larson & Tinsley 1978; Malin & Carter 1980;
Schombert et al. 1990). This type of interaction is relatively rare in the GEMS and GOODS
images, perhaps because the tidal debris is faint. The best cases are shown here and they
all have relatively small redshifts compared to the other interaction types (the average z is
0.23 and the maximum z is 0.69).
The image in the top left panel of Figure 1 (galaxy 1) has a giant diffuse clump in the
upper right corner. This could be a condensation in the tidal arm, or it could be another
galaxy. In either case, it has the same color as the rest of the tidal arm nearby. That is,
V606 − z850 = 0.90 ± 0.5 for the clump and also in six places along the tail; the color is
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essentially the same, 0.94 ± 0.05, in the core of the galaxy. The absolute magnitude of the
clump is MV = −18.41 for redshift z = 0.15. The mass is ∼ 5 × 10
9 M⊙ (Sect. 3.2). If
this clump is a condensation in the tail, then it could be a rare case where a pure stellar arc
has collapsed gravitationally into a gas-free tidal dwarf. The final result could be a dwarf
elliptical. Usually tidal dwarfs form by gaseous condensations in tidal arms (Wetzstein,
Naab, & Burkert 2007).
Figure 2 shows interactions that resemble the local Antennae pair, so we refer to them as
antennae types. These types have long tidal tails and double nuclei or highly distorted centers
that appear to be mergers of disk galaxies. Note that antennae are not the same as “tadpole”
galaxies (Elmegreen et al. 2005a; de Mello et al. 2006; Straughn et al. 2006), which have one
main clump and a sometimes wiggly tail that may contain smaller clumps. Some antennae
have giant clumps near the ends of the tails which could have formed there (galaxies 16 and
17) and are analogous to the clump at the end of the Superantennae (Mirabel et al. 1991).
Galaxy 18 is in a crowded field with at least two long tidal arms; here we consider only the
tail system in the north, which is in the upper part of the figure. These long-tail systems are
relatively rare and all the best cases are shown in the figure; their average redshift, 0.70, is
typical for GEMS and GOODS fields. Galaxy 24 is somewhat like a tadpole galaxy, but its
very narrow tail and protrusion on the anti-tail side of the main clump are unlike structures
seen in tadpoles of the Ultra Deep Field.
For the antennae galaxies in Figure 2, the tails have an average (V-z) color that is
negligibly bluer, 0.10±0.25 mag, than the central disks. In a study of tidal features in local
Arp atlas galaxies, Schombert et al. (1990) also found that the tail colors are uniform and
similar to those of the outer disks. They noted that the most sharply-defined tails are with
spiral systems and the diffuses plumes are with ellipticals. This correlation may be true
here also, but it is difficult to tell from Figure 1 whether the smooth distorted systems are
intrinsically disk-like.
Galaxy 20 in Figure 2 is an interesting case. It has an elliptical clump at the end of its
tail that could be one of the collision partners. There are two central galaxy cores, however,
and their interaction may have formed the tidal arms without this companion. Furthermore,
the clump at the tip is aligned perpendicular to the tail, which is unusual for a tidal dwarf.
Thus it is possible that the clump was a pre-existing galaxy lying in the orbital plane of one
of the larger galaxies now at the center. Presumably this former host is the galaxy currently
connected to the dwarf by the tidal arm. The interaction could have swung it around to
its current position at the tip. A similar case occurs for the local IC 2163/NGC 2207 pair,
which has a spheroidal dwarf galaxy at the tip of its tidal arm (Elmegreen et al. 2001). Such
swing-around dwarfs should have the same dynamical origin as the large pools of gas and
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star formation that are at the tips of superantenna-type galaxies; i.e. the whole outer disk
moves to this position during the interaction (Elmegreen et al. 1993; Duc et al. 1997).
Figure 3 shows examples of interactions that we refer to as M51-type galaxies, where
the tidal arms can be bridges that connect the main disk galaxy to the companion (galaxy
33), or tails on the opposite side of the companion (e.g., galaxies 34 and 35), or both (galaxy
36). In galaxy 44, the tidal arm looks like the debris path of a pre-existing galaxy that lies at
the right; the orbit path apparently curves around on the left. The M51-types usually have
strong spirals in the main disk. In the top row, the tails and bridges are thin and diffuse.
The galaxy on the left in the lower row (galaxy 42) has a thick, fan-shaped tail opposite
the companion. Some bridges have star formation clumps (galaxy 40) and others appear
smooth (galaxy 33). Interactions like this, especially those with small companions, are more
common than the previous two types and only a few best cases are shown in Figure 3 and
discussed in the rest of this paper.
Figure 4 shows examples of galaxies dominated by one highly curved, dominant arm and
large, regularly-spaced clumps of star formation. We call these “shrimp” galaxies because
of their resemblance to the tail of a shrimp. Although their star formation indicates they
contain gas and therefore are disk systems, there are no well-defined spirals (except for the
prominent arm), merging cores, or obvious central nuclei. The clumps resemble the beads-
on-a-string star formation in spiral density waves and probably have the same origin, a
gravitational instability (Elmegreen & Elmegreen 1983; Kim & Ostriker 2006; Bournaud,
Duc, & Masset 2003). The J-shaped morphology is reminiscent of the 90 kpc gas tail of M51
(Rots et al. 1990) and the 48 kpc gas tail observed in NGC 2535 (Kaufman et al. 1997).
Rots el al. point out that the M51 gas tail is much broader (10 kpc) than the narrow tails
seen in merging systems like the Antennae. The broad tail in galaxy 42 (Fig. 3) is similar
to the M51 tail. Sometimes there is a bright tail with no obvious companion (galaxies 56,
57, and 60); one of these, galaxy 56, was in our ring galaxy study (Elmegreen & Elmegreen
2006). Asymmetric, strong arm galaxies like this are not common in GEMS and GOODS;
this figure shows the best cases.
Figure 5 has a selection of irregular galaxies that appear to be interactions. Most of
them suggest an assembly of small pieces, so we refer to them as assembly types. If they were
slightly more round in overall shape, with more obvious interclump emission, then we would
classify them as clump-clusters, as we did in the UDF (Elmegreen et al. 2005a). The galaxy
in the lower left (galaxy 83) is like this. The resemblance of these types to clump-clusters
suggests that some of the clumps are accreted from outside the disk and others form from
gravitational instabilities in a pre-existing gas disk, as suggested previously (Elmegreen &
Elmegreen 2005). The system in the lower right (galaxy 85) could be interacting spirals, or a
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triple system, or a bent chain (as studied in Elmegreen & Elmegreen 2006). There are many
examples of highly irregular galaxies like these in the GEMS and GOODS fields; indeed most
galaxies at z > 1.5 are peculiar in this sense (Conselice 2005). In what follows, we discuss
only these 12 galaxies.
Figure 6 has samples of grazing or close interactions, with spirals at the top of the
page (numbers 86-93), ellipticals lower down (numbers 95-97) and two polar-ring galaxies
(numbers 99 and 100) in the lowest row at the middle and right bottom. We refer to these
paired systems as “equals” because their distinguishing feature is that the two galaxies have
comparable size. The pair number 89 has a bright oval in the smaller galaxy, which is char-
acteristic of recent tidal forces for an in-plane, prograde encounter such as IC2163/NGC2207
(Sundin 1993; Elmegreen et al. 1995 ). There is a spiral-elliptical pair on the right in the
middle row (galaxy 94). Double ellipticals in the UDF were studied previously (Elmegreen
et al. 2005a, Coe et al. 2006). Near neighbors like this have been studied previously in the
GEMS field; 6 double systems out of 379 red sequence galaxies were identified as being dry
merger candidates, as reproduced in simulations (Bell et al. 2006). The models of mergers
of early-type systems by Naab et al. (2006) apparently account for kinematic and isophotal
properties of ellipticals better than the formation of ellipticals through late-type mergers
alone. For the pairs in our figure, both components have the same COMBO17 redshift.
There are many other examples of close galaxy groups and near interactions in the GEMS
and GOODS surveys. In what follows we discuss only the properties of those shown in Figure
The interacting types shown in the figures are meant to be as distinct as possible. These
and other good cases are listed in Table 1 by running number, along with their COMBO-
17 catalog number, redshift, and R magnitude. There is occasionally some ambiguity and
overlap in the interaction types, particularly between M51-types and shrimps when the M51-
types have small or uncertain companions at the ends of their prominent tails. Projection
effects can lead to uncertainties in the classifications as well, particularly for antennae whose
tails may be foreshortened. Nevertheless, these divisions serve as a useful attempt to sort out
the most prominent features among interacting galaxies. There are numerous other galaxies
in GEMS and GOODS that are apparently interacting, but most of them are too highly
distorted to indicate the particular physical properties of interest here, namely, disk-to-halo
mass ratio and star formation scale.
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3. Photometric Results
3.1. Global galaxy properties
The integrated Johnson restframe (U-B) and (B-V) colors from COMBO-17 for the
observed galaxies with measured redshifts are shown in a color-color diagram in Figure 7.
The crosses in the diagram are Johnson colors for standard Hubble types (Fukugita et al.
1995). Our sample of galaxies spans the range of colors from early to late Hubble types,
although the bluest are bluer than standard irregular galaxies (a typical Im has U-B= −0.35,
B-V= 0.27). The reddest galaxies tend to be the diffuse types, thought to originate with
ellipticals involved in interactions. The two reddest galaxies in our sample are the diffuse
types number 1 and 2 in Figure 1. The bluest tend to be the assemblies, consistent with
their having formed recently.
Figure 8 shows a restframe color-magnitude diagram. Early and late type galaxies
usually separate into a “red sequence” and a “blue cloud” on such a diagram (Baldry et
al. 2004; Faber et al. 2005). The solid line indicates the boundary between these two
regions from a study of 22,000 nearby galaxies (Conselice 2006b). The short-dashed lines
are the limits of the Conselice (2006) survey; local galaxies are brighter than the vertical
short-dashed line and their colors lie between the horizontal short-dashed lines. The long-
dashed lines approximately outline the bright limit for the local blue cloud galaxies. Our
galaxies fall in both the red sequence and the blue cloud. The restframe colors in Figure 8
are consistent with their morphological appearances. The red sequence galaxies in the figure
usually appear smooth (the diffuse types) or lack obvious huge star formation clumps (the
equal mass mergers), while the blue cloud galaxies usually have patches of star formation
(the M51-types, shrimps, assemblies, and many antennae). We see now why the redshifts of
the diffuse galaxies (z < 0.3) are much lower than the others: this is a selection effect for the
ACS camera. These tails comprise old stellar populations without star-forming clumps, and
their intrinsic redness makes them difficult to see at high redshifts. Also, they tend to have
intrinsically low surface brightnesses because of a lack of star formation, and cosmological
dimming makes them too faint to see at high redshift. Hibbard & Vacca (1997) note that it
is difficult to detect tidal arms beyond z ∼ 1.5.
3.2. Clump properties
Prominent star-forming clumps are apparent in many of the interacting galaxies. Their
sizes and magnitudes were measured using rectangular apertures. The observed magnitudes
were converted to restframe B magnitudes whenever possible, using linear interpolations
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between the ACS bands. For example, GEMS observations are at two filters, V606 and z850.
GEMS galaxies with redshifts z between 0.39 (= 606/435 − 1) and 0.95 (= 850/435 − 1)
were assumed to have restframe blue luminosities given by LB,rest = LV,obs(0.95− z)/(0.95−
0.39) + Lz,obs(z − 0.39)/(0.95 − 0.39). The restframe B magnitude is then −2.5 logLB,rest.
For GOODS galaxies, the conversions were divided into 3 redshift bins to make use of
the 4 available filters, and a linear interpolation was again applied to get restframe clump
magnitudes. For the GOODS galaxies, the restframe magnitudes determined by interpolation
between the nearest 2 filters among the 4 filters are within ±0.2 mag of the restframe
magnitudes determined from only the V and z filters. Thus, the GEMS interpolations are
accurate to this level. (We do not include corrections for intergalactic absorption in these
colors, because we are comparing them directly with their parent galaxy properties. Below,
when we convert the colors and magnitudes to masses and ages, absorption corrections are
taken into account.)
The apparent restframe B magnitudes of the clumps were converted to absolute rest-
frame B magnitudes using photometric redshifts and the distance modulus for a ΛCDM
cosmology. These absolute clump magnitudes are shown as a function of absolute galaxy
magnitude in Figure 9. The clump absolute B magnitudes scale linearly with the galaxy
magnitudes. The clumps are typically a kpc in size (∼ 3 to 8 pixels across), comparable
to star-forming complexes in local galaxies (Efremov 1995), which also scale with galaxy
magnitude (Elmegreen et al. 1996; Elmegreen & Salzer 1999).
Clump ages and masses were estimated by comparing observed clump colors, magni-
tudes, and redshifts with evolutionary models that account for bandshifting and intergalac-
tic absorption and that assume an exponential star formation rate decay (see Elmegreen &
Elmegreen 2005). Internal dust extinction as a function of redshift is taken from Rowan-
Robinson (2003). The GEMS galaxy clumps only have (V606-z850) colors, so the ages are not
well constrained. For the GOODS galaxies, the additional B and I filters help place better
limits on the ages, although there is still a wide range of possible fits.
Figure 10 shows sample model results for redshift z = 1. The different lines in each
panel correspond to different decay times for the star formation rate, in years: 107, 3× 107,
108, 3×108, and 109, and the sixth line represents a constant rate. Generally the shorter the
decay time, the redder the color and higher the mass for a given duration of star formation.
This correspondence between color and mass gives a degeneracy to plots of mass versus color
at a fixed apparent magnitude (top left) and apparent magnitude versus color at a fixed mass
(top right). Thus the masses of clumps can be derived approximately from their colors and
magnitudes, without needing to know their ages or star formation histories.
Figure 11 shows observations and models in the color-magnitude plane for 6 redshift in-
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tervals spanning our galaxies. Each curve represents a wide range of star formation durations
that vary along the curve as in the top right panel of Fig. 10; each curve in a set of curves is
a different decay time. The different sets of curves, shifted vertically in the plots, correspond
to different clump masses, as indicated by the adjacent numbers, which are in M⊙. Each
different point is a different clump; many galaxies have several points. Only clumps with
both V606 and z850 magnitudes above the 2σ noise limit are plotted in Figure 11. The clump
(V606 − z850) colors range from 0 to 1.5. The magnitudes tend to be about constant for each
redshift because of a selection effect (brighter magnitudes are rare and fainter magnitudes
are not observed).
Figure 11 indicates that the masses of the observable clumps are between 106 and 109
M⊙ for all redshifts, with higher masses selected for the higher redshifts. The masses for all
of the clumps are plotted in Figure 12 versus the galaxy type (types 1 through 6 are in order
of Figs. 1 through 6 above). The masses are obtained from the observed values of V606 and
V606 − z850 using the method indicated in Figure 11. The different mass evaluations for the
six decay times are averaged together in the log to give the log of the mass plotted as a dot
in Figure 12. The rms values of log-mass among these six evaluations are shown in Figure
12 as plus-symbols, using the right-hand axes. These rms deviations are less than 0.2, so
the uncertainties in star formation decay times and clump ages do not lead to significant
uncertainties in the clump mass. (Systematic uncertainties involving extinctions, stellar
evolution models, photometric redshifts, and so on, would be larger.)
The clump ages cannot be determined independently from the star formation decay
times with only the few passbands available at high angular resolution. Figure 13 shows
model results that help estimate the clump ages. As in the other figures, each line is a
different exponential decay time for the star formation rate. If we consider the two extreme
decay times in this figure (continuous star formation for the bottom lines in each panel and
107 years for the top lines), then we can estimate the age range for each decay time from
the observed color range. For V606 − z850 colors in the range from 0 to 0.5 at low z (cf. Fig
11), the clump ages range from 107 to 1010 yr with continuous star formation and from 107
to 3 × 108 yr with a decay time of 107 yrs. For colors in the range from 0 to 1.5 at higher
redshifts, the age ranges are about the same in each case. For intermediate decay times, the
typical clump ages are between ∼ 107 years for the bluest clumps and ∼ 109 years for the
reddest clumps. These are reasonable ages for star formation regions, and consistent with
model tail lifetimes.
The star-forming complexes in the GEMS and GOODS interacting galaxies are 10 to
1000 times more massive than the local analogs seen in non-interacting late-type galaxies
(Elmegreen & Salzer 1999), but the low mass end in the present sample is similar to the
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high mass end of the complexes measured in local interacting galaxies. For example, the
Tadpole galaxy, UGC 10214, contains 106 M⊙ complexes along the tidal arm (Tran et al.
2003; Jarrett et al. 2006). The interacting galaxy NGC 6872 has tidal tails with 109 M⊙ HI
condensations (Horellou & Koribalski 2007), but the star clusters have masses only up to
106 M⊙ (Bastian et al. 2005). The most massive complexes in the tidal tail of NGC 3628
in the Leo Triplet are also ∼ 106 M⊙ (Chromey et al. 1998). The NGC 6872 clusters differ
qualitatively from those in our sample in being spread out along a narrow arm; ours are big
round clumps spaced somewhat evenly along the arm. Small star clusters are also scattered
along the tidal arms the Tadpole and Mice systems; they typically contain less than 106 M⊙
(de Grijs et al. 2003). The NGC 3628 clusters are also faint with surface brightnesses less
than 27 mag arcsec−2; they would not stand out at high redshift.
It is reasonable to consider whether the observed increase of complex mass with increas-
ing redshift is a selection effect. Our clumps are several pixels in size, corresponding to a
scale of ∼ 1 kpc. Individual clusters are not resolved and we only sample the most massive
conglomerates. These kpc sizes are comparable to the complex sizes in local galaxies, but the
high redshift complexes are much brighter and more massive. They would be observed easily
in local galaxies. The massive complexes in our sample are more similar to those measured
generally in UDF galaxies (Elmegreen & Elmegreen 2005).
Clump separations were measured for clumps along the long arms in the shrimp galaxies
of Figure 4. They average 2.20±0.94 kpc for 49 separations. This is about the same separa-
tion as that for the largest complexes in the spiral arms of local spiral galaxies (Elmegreen
& Elmegreen 1983, 1987), and comparable to the spacing between groups of dust-feathers
studied by La Vigne et al. (2006). Yet the clumps in shrimp galaxies and others studied here
are much more massive than the complexes in local spiral arms, which are typically < 106
M⊙ in stars and ∼ 10
7 M⊙ in gas. This elevated mass can be explained by a heightened
turbulent speed for the gas, combined with an elevated gas density. Considering that the
separation is about equal to the two-dimensional Jeans length, λ ∼ 2a2/ (GΣ) for velocity
dispersion a and mass column density Σ, and that the mass is the Jeans mass, λ2Σ, the
mass scales with the square of the velocity dispersion, M = M0 (a/a0)
for fixed length
λ0 = 2a
0/ (GΣ0) and M0 = λ
0Σ0. The mass column density also scales with the square of
the dispersion, Σ = Σ0 (a/a0)
to keep λ constant. Thus the interacting tidal arm clumps
are massive because the velocity dispersions and column densities are high. Another way
to derive this is to note that for regular spiral arm instabilities, 2Gµ/a2 is about unity at
the instability threshold, where µ is the mass/length along the arm (Elmegreen 1994). Thus
cloud mass scales with a2 for constant cloud separation. High velocity dispersions for neutral
hydrogen, ∼ 50 km s−1, are also observed in local interacting galaxies (Elmegreen et al. 1993;
Irwin 1994; Elmegreen et al. 1995; Kaufman et al. 1997; Kaufman et al. 1999; Kaufman
– 12 –
et al. 2002). Presumably the interaction agitates the interstellar medium to make the large
velocity dispersions. The orbital motions are forced to be non-circular and then the gaseous
orbits cross, converting orbital energy into turbulent energy and shocks. Similar evidence
for high velocity dispersions was found in the masses and spacings of star forming complexes
in clump cluster galaxies (Elmegreen & Elmegreen 2005) and in spectral line widths (Genzel
et al. 2006; Weiner et al. 2006).
3.3. Tail Properties
Figure 14 shows the average tail surface brightness as a function of (1+ z)4 for galaxies
in Figures 1-4. Some systems have more than one tail. Cosmological dimming causes a fixed
surface brightness to get fainter as (1+z)−4, so there should be an inverse correlation in this
diagram. Clearly, the tails are brighter for the more nearby galaxies, and they decrease out
to z ∼ 1, where they are fairly constant. This constant limit is at the 2σ detection limit of
25 mag arcsec−2. Antennae galaxies with average tail surface brightnesses fainter than this
limit have patchy tails with no apparent emission between the patches. Only the brightest
high redshift tails can be observed in this survey.
Simulations by Mihos (1995) suggested that tidal tails are observable for a brief time in
the early stages of a merger, corresponding to ∼ 150 Myr at a redshift z = 1 and 350 Myr
at z = 0.4. The difference is the result of surface brightness dimming as tails disperse. A
nearby galaxy merger, Arp 299, has a 180 kpc long tail encompassing 2 to 4% of the total
galaxy luminosity, with an interaction age of 750 Myr, but its low surface brightness of 28.5
mag arcsec−2 (Hibbard & Yun 1999) would be below the GOODS/GEMS detection limit.
The ratio of the luminosity of the combined tails and bridges to the luminosity of the
disk (the luminosity fraction) is shown in Figure 15. The luminosity fraction in the tidal
debris ranges from 10% to 80%, averaging about 30% regardless of redshift. This range is
consistent with that of local galaxies in the Arp atlas and Toomre sequence (e.g., Schombert,
Wallin, & Struck-Marcell 1990; Hibbard & van Gorkom 1996).
Interaction models with curled tails, as in our shrimp galaxies, were made by Bournaud
et al. (2003). Their models had dark matter halos with masses ∼ 10 times the disk mass
and extents less than 12 disk scale lengths. Some of our shrimp galaxies have one prominent
curved arm that is pulled out from the main disk but not very far, resulting in a lopsided
galaxy. Simulations indicate that such lopsidedness may be the result of a recent minor
merger (Bournaud et al. 2005). In some of our cases, a nearby companion is obvious.
The linear sizes of the tidal tails in our sample are shown in Figure 16. They range from
– 13 –
2 to 60 kpc, and are typically a few times the disk diameter, as shown in Figure 17, which
plots this ratio versus redshift. The average tail to diameter ratio is 2.9±1.7 for diffuse tails,
2.5± 1.3 for antennae, 2.5± 1.1 for M51-types and 1.5± 1.4 for shrimps, so the shrimps are
about 60% as extended as the antennae types. There is no apparent dependence of these
ratios on redshift in Figure 17. Projection effects make these apparent ratios smaller than
the intrinsic ratios.
For comparison, the ratio of tail length to disk diameter versus the tail length for local
galaxies is shown in Figure 18 based on measurements of antennae-type systems in the Arp
atlas (1966) and the Vorontsov-Velyaminov atlas (1959). Our galaxies are also shown. The
average tail length for the local galaxies in this figure is 72 ± 48 kpc, while the average tail
length for the GEMS and GOODS antennae is 37% as much, 27±16 kpc. The diameters for
these two groups are 20±12 kpc and 11±5 kpc, and the ratios of tail length to diameter are
4.5± 3.7 and 2.5± 1.3, respectively. Thus the local antennae mergers are larger in diameter
by a factor of 2 than the GEMS and GOODS antennae, and the tails for the locals are larger
by a factor of 2.7. These results for the diameters are consistent with other indicators that
galaxies are smaller at higher redshift, although usually this change does not show up until
z > 1 (see observations and literature review in Elmegreen et al. 2007).
3.4. Tidal dwarf galaxy candidates
Three antennae galaxies at the top of Figure 2, numbers 15, 16, and 18, have long
straight tidal arms with large star-forming regions at the ends. These clumps are possibly
tidal dwarf galaxies. The clump diameters and restframe B magnitudes are listed in Table
2, along with the clump in diffuse galaxy number 1 discussed in Sect. 2. Listed are their
V606 and V606 − z850 magnitudes and associated masses, calculated as in Sect. 3.2. The
masses range from 0.2× 108 to 4.6× 108 M⊙ for the star-forming dwarfs, but for the stellar
condensation in the diffuse-tail galaxy 1 (Fig. 1), the mass is 50×108 M⊙. The star-forming
dwarf masses are similar to or larger than those found for the tidal object at the end of the
Superantennae (Mirabel et al. 1991) as well as the tidal object at the end of the tidal arm in
the IC 2163/NGC 2207 interaction (Elmegreen et al. 2001) and at the end of the Antennae
tail (Mirabel et al. 1992). The HI dynamical masses for these local tidal dwarfs are ∼ 109
Simulations of interacting galaxies that form tidal dwarf galaxies require long tails and
a dark matter halo that extends a factor of 10 beyond the optical disk (Bournaud et al.
2003). If one or both galaxies contain an extended gas disk before the interaction, then more
massive, 109 M⊙ stellar objects can form at the tips of the tidal arms from the accumulated
– 14 –
pool of outer disk material (Elmegreen et al. 1993; Bournaud et al. 2003). Observations of
nearby interactions show clumpy regions of tidal condensations with masses of ∼ 108 − 109
M⊙ (Bournaud et al. 2004; Weilbacher et al. 2002, 2003; Knierman et al. 2003; Iglesias-
Paramo & Vilchez 2001), like what is observed in our high redshift tidal dwarfs.
No well-resolved models have yet formed tidal dwarfs from stellar debris. Wetzstein,
Naab, & Burkert (2007) considered this possibility and found collapsing gas more likely. Yet
the condensed object in the tail of galaxy 1 could have formed there and it is interesting
to consider whether the Jeans mass in such an environment is comparable to the observed
mass. If, for example, the tidal arm surface density corresponds to a value typical for the
outer parts of disks, ∼ 10 M⊙ pc
−2, and the stellar velocity dispersion is comparable to that
required in Sect. 3.2 for the gas to give the giant star forming regions, ∼ 40 km s−1, then the
Jeans mass is M ∼ a4/ (G2Σ) ∼ 1010 M⊙. This is not far from the value we observe, 5× 10
M⊙, so the diffuse clump could have formed by self-gravitational collapse of tidal tail stars.
The timescale for the collapse would be a/ (πGΣ) ∼ 300 Myr, which is not unreasonable
considering that the orbit time at this galactocentric radius is at least this large.
4. Dark Matter Halo Constraints
Models of interacting galaxies have been used to place constraints on dark halo poten-
tials. Springel & White (1999) and Dubinski, Mihos, & Hernquist (1999) found that tidal
tail lengths can be long compared to the disk if the ratio of escape speed to rotation speed
at 2 disk scale lengths is small, ve/Vr < 2.5, and the rotation curve is falling in the outer
disk. In a series of models, Dubinski et al. showed that this condition may result from either
disk-dominated rotation curves where the halo is extended and has a low concentration, or
halo-dominated rotation curves where the halo is compact and low mass. Dubinski et al.
point out that the latter possibility is inconsistent with observed flat or rising disk rotation
curves, but the first is compatible if the disk is massive and dominant in the inner regions.
The first case also gives prominent bridges. In addition, Springel & White (1999) found that
CDM halo models with embedded disks allow long tidal tails, but Dubinski et al. noted that
most of those which do are essentially low surface brightness disks in massive halos, and not
normal bright galaxies. Galaxies without dark matter halos are not capable of generating
long tidal tails (Barnes 1988). In all cases, longer tails develop in prograde interactions.
The smooth diffuse types and antenna types in Figures 1 and 2 have relatively long tails,
so the progenitors were presumably disks of early and late types, respectively, with falling
rotation curves in their outer parts. These long-tail cases are relatively rare, comprising
only about 8% and 9%, respectively, of our original (300 galaxy) interacting sample from
– 15 –
GEMS and GOODS. The more compact M51 types and shrimps represent 9% and 12% of
the sample. Short tail interactions could be younger, less favorably projected, or have a
more steeply rising rotation curve than long tail interactions. The M51 types have clear
companions, so the prominent features are bridges. According to Dubinski et al. (1999),
bridging requires a prograde interaction with a maximum-disk galaxy, that is, one with a
low-mass, extended halo.
5. Conclusions
Mergers and interactions out to redshift z = 1.4 have tails, bridges, and plumes that are
analogous to features in local interacting galaxies. Some interactions have only smooth and
red features, indicative of gas-free progenitors, while others have giant blue star-formation
clumps. The tail luminosity fraction has a wide range, comparable to that found locally.
A striking difference arises regarding the tail lengths, however. The tails in our antenna
sample, at an average redshift of 0.7, are only one-third as long as the tails in local antenna
mergers, and the disk diameters are about half the local merger diameters. This difference is
consistent with the observations that high redshift galaxies are smaller than local galaxies,
although such a drop in size has not yet been seen for galaxies at redshifts this low. The
implication is that dark matter halos have not built up to their full sizes for typical galaxies
in GEMS and GOODS.
Star formation is strongly triggered by the interactions observed here, as it is locally.
The star-forming clumps tend to be much more massive than their local analogs, however,
with masses between ∼ 106 M⊙ and a few ×10
8 M⊙, increasing with redshift. This is not
merely a selection effect, since the massive clumps seen at high redshift would show up at
lower redshift, although of course smaller clumps would not be resolved at high redshift. The
clump spacings were measured along the tidal arms of the most prominent one-arm type of
interaction, the shrimp-type, and found to be 2.20 ± 0.94, which is typical for the spacing
between beads on a string of star formation in local spiral arms. If both types of arms form
clumps by gravitational instabilities, then the turbulent speed of the interstellar medium in
the GEMS and GOODS sample has to be larger than it is locally by a factor of ∼ 5 or more;
the gas mass column density has to be larger by this factor squared.
Some interactions have tidal dwarf galaxies at the ends of their tidal arms, similar to
those found in the Superantennae galaxy and other local mergers. One diffuse interaction
with red stellar tidal debris has a large stellar clump that may have formed by gravitational
collapse in a stellar tidal arm; the clump mass is 5 × 109 M⊙. Long-arm interactions are
relatively rare, comprising only ∼ 17% of our total sample of ∼ 300 interacting systems (only
– 16 –
a fraction of which were discussed here). For those with long arms, numerical models suggest
the dark matter halos must be extended, so that the rotation curves are falling in the outer
disks. Most interactions are not like this, however, so the rotation curves are probably still
rising in their outer disks, like most galaxies locally.
We gratefully acknowledge summer student support for B.M. and T.F. through an REU
grant for the Keck Northeast Astronomy Consortium from the National Science Foundation
(AST-0353997) and from the Vassar URSI (Undergraduate Research Summer Institute) pro-
gram. D.M.E. thanks Vassar for publication support through a Research Grant. We thank
the referee for useful comments. This research has made use of the NASA/IPAC Extragalac-
tic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute
of Technology, under contract with the National Aeronautics and Space Administration.
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– 21 –
Table 1. Interacting Galaxies in GEMS and GOODS
Type, Figure Number COMBO 17 z R mag.
Diffuse (Fig. 1) 1 6423 0.15 16.572
2 12639 0.154 16.678
3 11538 0.134 17.713
4 53129 0.171 16.968
5 57881 0.118 17.552
6 28509 0.093 18.79
7 17207 0.69 19.742
8 30824 0.341 19.755
9 25874 0.262 19.757
Diffuse (other) 10 22588 0.684 21.263
11 21990 0.429 21.243
12 46898 0.617 20.794
13 49709 0.302 20.23
14 15233 0.304 18.857
Antennae (Fig. 2) 15 61546 0.552 20.41
16 45115 0.579 21.275
17 20280 0.555 21.653
18 41907 0.702 22.66
19 35611 1.256 22.655
20 10548 0.698 22.43
21 33650 0.169 18.86
22 42890 0.421 20.68
23 49860 1.169 23.632
24 34926 0.779 -19.69
Antennae (other) 25 14829 0.219 21.429
26 18588 0.814 22.748
27 46738 1.204 20.65
28 7551 1.162 25.926
29 20034 1.326 21.932
30 33267 0.067 23.112
31 38651 0.988 23.89
32 55495 1.00 24.261
M51-type (Fig. 3) 33 5640 0.204 19.477
34 9415 0.523 21.16
35 40901 0.193 19.751
36 17522 0.82 23.103
37 6209 1.187 22.723
38 23667 1.151 23.514
39 37293 0.274 20.533
40 39805 0.557 20.089
41 53243 0.698 21.683
42 15599 0.56 21.381
43 25783 0.663 20.732
44 39228 0.117 18.031
M51-type (other) 45 1984 0.762 22.855
– 22 –
Table 1—Continued
Type, Figure Number COMBO 17 z R mag.
46 2760 1.281 23.202
47 15040 0.667 22.392
48 18502 0.228 21.942
49 14959 0.306 19.581
50 16023 0.668 21.887
51 30226 0.509 22.689
52 40744 0.292 21.119
53 45102 0.857 22.514
54 60582 0.946 22.54
Shrimp (Fig. 4) 55 40198 0.201 20.55
56 14373 0.795 23.183
57 12222 1.004 22.417
58 28344 0.257 19.509
59 56284 0.657 21.667
60 2385 0.283 21.334
61 54335 0.892 22.824
62 28841 0.673 20.971
63 6955 0.983 22.24
Shrimp (other) 64 34244 0.999 22.504
65 48298 0.429 21.663
66 37809 0.357 20.667
67 25316 0.985 23.717
68 49595 0.663 21.939
69 59467 0.487 21.568
70 9062 0.854 23.82
71 30076 0.832 22.672
72 2760 1.281 23.202
73 54335 0.892 22.824
Assembly (Fig. 5) 74 28751 0.093 23.506
75 4728 0.702 22.799
76 23187 1.183 23.565
77 45309 1.061 22.916
78 41835 0.098 19.134
79 61945 1.309 21.813
80 62605 1.011 23.143
81 44956 0.506 22.494
82 4546 0.809 22.163
83 23000 0.132 22.951
84 63112 0.499 22.273
85 43975 1.059 22.878
Equal (Fig. 6) 86 40813 0.182 19.983
87 8496 0.354 22.415
88 13836 0.661 21.054
89 11164 0.464 19.351
90 39877 0.493 22.142
– 23 –
Table 1—Continued
Type, Figure Number COMBO 17 z R mag.
91 40598 0.263 20.128
92 51021 0.743 20.96
93 35317 0.671 20.755
94 56256 0.502 20.309
95 47568 0.649 20.206
96 40766 0.46 19.997
97 24927 0.524 19.647
98 15233 0.304 18.857
99 18663 1.048 24.011
100 43242 0.657 21.177
– 24 –
Table 2: Tidal Dwarf Galaxy Candidates
Galaxy z Galaxy Diam. Dwarf V606 V606 − z850 Clump Mass
(COMBO17 #) MB,rest (mag) (kpc) MB,rest (mag) mag mag x10
1 (6423) 0.15 -20.64 13.9 -17.55 20.83 0.90 50
15 (61546) 0.552 -20.77 5.5 -16.67 26.12 0.78 1.2
16 (45115) 0.579 -20.17 4.7 -17.72 25.42 1.1 4.6
18 (41907) 0.702 -19.21 1.9 -16.03 27.45 0.55 0.24
17 (20280) 0.555 -19.56 6.2 -17.06 25.77 0.73 1.4
– 25 –
Fig. 1.— Color images of galaxies in the GEMS and GOODS fields with smooth diffuse tidal
debris. The galaxy at the top right, number 3 in Table 1, is only partially covered by the
GEMS field; the right-hand portion of the image is from ground-based observations. The
smooth debris is presumably from old stars that were spread out during the interaction. A
few small star-formation patches are evident in some cases. The clump in the upper right
corner of the galaxy 1 image could be a rare example of a gravitationally driven condensation
in a pure-stellar arm. The smooth arcs and spirals in this and other images are probably
a combination of orbital debris and flung-out tidal tails. The galaxy numbers, as listed in
Table 1, are 1 through 9, as plotted from left to right and top to bottom. (Image quality
degraded for astroph.)
– 26 –
Fig. 2.— Color images of interacting antennae galaxies with long and structured tidal arms.
Galaxy numbers, in order, are 15 through 24. Several have dwarf galaxy-like condensations
at the arm tips or broad condensations midway out in the arms. The dwarf elliptical at the
tip of the tidal arm in galaxy 20 might have existed before the interaction and been placed
there by tidal forces; the main body of this system has a double nucleus from the main
interaction. (Image quality degraded for astroph.)
– 27 –
Fig. 3.— M51-type galaxies are shown as logarithmic grayscale V-band images. In order,
the galaxy numbers are 33 through 44. The linear streak in galaxy 44 could be orbital debris
from the small companion on the right. (Image quality degraded for astroph.)
– 28 –
Fig. 4.— Shrimp galaxies, named because of their curved tails, are shown as logarithmic
V-band images. In order, they are numbers 55 through 63. (Image quality degraded for
astroph.)
– 29 –
Fig. 5.— Assembly galaxies look like they are being assembled through mergers. In order:
galaxy 74 through 85.
– 30 –
Fig. 6.— Galaxies with approximately equal-mass grazing companions, in order, are 86
through 100.
– 31 –
0 0.2 0.4 0.6 0.8 1
Restframe Johnson B–V
Equal
Assembly
Shrimp
M51 type
Antenna
Diffuse
Fig. 7.— Restframe (U-B) and (B-V) integrated colors for interacting galaxies in the GEMS
and GOODS fields, from COMBO-17. The reddest tend to be the diffuse types, which
are presumably dry mergers, and the bluest are the assembly types, which could be young
proto-galaxies. Crosses indicate standard Hubble types, measured by Fukugita et al. (1995).
– 32 –
–24 –22 –20 –18 –16 –14 –12
MB (mag)
Diffuse
Antenna
M51 type
Shrimp
Assembly
Equal
Fig. 8.— Restframe Johnson U-B integrated color versus absolute restframe MB, from
COMBO-17. The solid line separates the red sequence and blue cloud (Conselice 2006b).
Color limits for local galaxies are indicated by the horizontal short-dashed lines; local galaxies
are brighter than the vertical line. The local blue cloud galaxies are approximately delimited
on the left side of the diagram by the long-dashed lines. Thus, most of our observed galaxies
fall near the local galaxy colors and magnitudes.
– 33 –
–17 –18 –19 –20 –21 –22
Galaxy MB
Diffuse
Antenna
M51 type
Shrimp
Assembly
Tidal Dwarf
Fig. 9.— Restframe B absolute magnitudes of star-forming clumps versus integrated galaxy
restframe magnitudes. The correlation is also found for local galaxies.
– 34 –
0.01 0.1 1 10
Duration of SF (Gyr)
0.01 0.1 1 10
Duration of SF (Gyr)
–1 0 1 2 3
V606–z850
–1 0 1 2 3
V606–z850
Fig. 10.— Models at z = 1 for clump color (bottom left) and clump mass at an apparent
V606 magnitude of 27 (lower right) are shown in the bottom panels versus the duration of star
formation in 6 models with exponentially decaying star formation. Five lines are for decay
times of 107, 3×107, 108, 3×108, and 109 years, and the sixth line represents a constant rate.
Shorter decay times correspond to redder color (upper lines) and higher masses (upper lines).
In the top panels, the clump mass at V606 = 27 (top left) and the clump apparent magnitude
at 108 M⊙ masses (top right) are shown versus the clump color. The correspondence between
color and mass gives a degeneracy to plots of mass versus color at a fixed apparent magnitude
(top left) and apparent magnitude versus color at a fixed mass (top right). Thus the masses
of clumps can be derived approximately from their V606 − z850 colors and V606 magnitudes
for each redshift.
– 35 –
–1 0 1 2 3
V606–z850
0–0.125
Diffuse
Antenna
M51 type
Shrimp
Assembly
Equal
Tidal Dw.
0 1 2 3
V606–z850
0.125–0.375
0.375–0.625
0.625–0.875
z=0.875–1.125
z=1.125–1.375
Fig. 11.— The masses of the clumps can be estimated from this figure. Each curve in a
cluster of curves is a different model for color-magnitude evolution of a star-forming region,
with the age of the region changing along the curve and the exponential decay rate of the
star formation changing from curve to curve. The different clusters of curves correspond to
different total masses for the star-forming regions (mass in M⊙ is indicated to the right of each
curve). The symbols represent observations of apparent magnitude and color. Bandshifting
and absorption are considered by plotting the observations and models in redshift bins.
The mass scales shift slightly with redshift. The mass of each star-forming region can be
determined by interpolation between the curves. Typical masses are 106 M⊙ for low z and
108 M⊙ for high z. The circle near the 10
10 M⊙ curves in the z = 0.125 − 0.375 interval
corresponds to the diffuse clump in the tidal debris of galaxy 1 in Fig. 1.
– 36 –
D A M S AS E T
0–0.125
D A M S AS E T
0.125–0.375
log M
= 9.7
0.375–0.625 0.625–0.875
z=0.875–1.125 z=1.125–1.375
Fig. 12.— Clump masses (left axis) are plotted versus galaxy type in order of Figs. 1-6:
Diffuse, Antenna, M51-type, Shrimp, Assembly, and Equal, with T representing the tidal
dwarfs. The method of Fig. 11 is used. The rms deviations among the six star formation
decay times are shown as plus-symbols using the right-hand axes.
– 37 –
0.01 0.1 1 10
Duration of SF (Gy)
50 z=0.25
0.01 0.1 1 10
Duration of SF (Gy)
z=0.5
z=0.75 z=1
z=1.25 z=1.5
Fig. 13.— The apparent color of a star forming region is shown versus the duration of star
formation for an exponentially decaying star formation law. The decay times are as in Fig.
10, with short decay times the upper lines and continuous star formation the lower lines.
Using the observed clump colors, the durations of star formation are found to range between
107 and 3× 108 yrs for short decay times.
– 38 –
0 0.5 1 1.5
log(1+z)4
Diffuse
Antenna
M51 type
Shrimp
Fig. 14.— V-band surface brightness of tidal tails for galaxies in Figures 1-4 plotted as a
function of (1 + z)4 for redshift z. Some systems have more than one tail. Cosmological
dimming causes a decrease with redshift equal to 2.5 magnitudes for each factor of 10 in
(1 + z)
; this decrease is consistent with the dimming seen here. The observable 2σ limit
for these fields is ∼ 25 mag arcsec−2. Some antenna galaxies have patchy tails with fainter
average surface brightnesses.
– 39 –
0 0.5 1
Redshift, z
Diffuse
Antenna
M51 type
Shrimp
Fig. 15.— Fraction of V-band luminosity in antennae tidal tails relative to their integrated
galaxy luminosity, as a function of redshift.
– 40 –
0 5 10 15 20 25 30 35
Disk Diameter (kpc)
Diffuse
Antenna
M51 type
Shrimp
Fig. 16.— Tail length versus disk diameter from Figs. 1-4, based on the V-band images.
Conversions to linear size assumed a standard ΛCDM cosmology applied to the photometric
redshifts.
– 41 –
0 0.5 1 1.5
Redshift, z
Diffuse
Antenna
M51 type
Shrimp
Fig. 17.— Tail length/disk diameter as a function of redshift for shrimps and antennae,
measured from the V-band images. There is no obvious trend.
– 42 –
0 50 100 150 200
Tail Length (kpc)
GEMS GOODS Antennae
Local Antennae
Superantennae
Arp 299
NGC 3628Arp 241
VV109
NGC 3256
Antennae
Arp 243
Arp 242
Arp 226
Arp 157 Arp 75
Arp 35 NGC 5548
Arp 33
Arp 102
Fig. 18.— Tail length/disk diameter versus the tail length for antenna galaxies in our sample
as well as for local antennae, whose names are indicated. The GEMS and GOODS systems
are significantly smaller than the local antenna galaxies, even if the two extreme local cases,
the Superantennae and Arp 299, are excluded.
	Introduction
	The Sample of Interactions and Mergers
	Photometric Results
	Global galaxy properties
	Clump properties
	Tail Properties
	Tidal dwarf galaxy candidates
	Dark Matter Halo Constraints
	Conclusions

  GEMS and GOODS fields were examined to z~1.4 for galaxy interactions and
mergers. The basic morphologies are familiar: antennae with long tidal tails,
tidal dwarfs, and merged cores; M51-type galaxies with disk spirals and tidal
arm companions; early-type galaxies with diffuse plumes; equal-mass
grazing-collisions; and thick J-shaped tails beaded with star formation and
double cores. One type is not common locally and is apparently a loose
assemblage of smaller galaxies. Photometric measurements were made of the tails
and clumps, and physical sizes were determined assuming photometric redshifts.
Antennae tails are a factor of ~3 smaller in GEMS and GOODS systems compared to
local antennae; their disks are a factor of ~2 smaller than locally. Collisions
among early type galaxies generally show no fine structure in their tails,
indicating that stellar debris is usually not unstable. One exception has a
5x10**9 Msun smooth red clump that could be a pure stellar condensation. Most
tidal dwarfs are blue and probably form by gravitational instabilities in the
gas. One tidal dwarf looks like it existed previously and was incorporated into
the arm tip by tidal forces. The star-forming regions in tidal arms are 10 to
1000 times more massive than star complexes in local galaxies, although their
separations are about the same. If they all form by gravitational
instabilities, then the gaseous velocity dispersions in interacting galaxies
have to be larger than in local galaxies by a factor of ~5 or more; the gas
column densities have to be larger by the square of this factor.

Introduction
Galaxy interactions and mergers are observed at all redshifts and play a key role in
galaxy evolution. Two percent of local galaxies are interacting or merging (Athanassoula &
Bosma 1985; Patton et al. 1997), and this fraction is larger at high redshift (e.g., Abraham
et al. 1996b; Neuschaefer et al. 1997 ; Conselice et al. 2003; Lavery et al. 2004; Straughn
et al. 2006; Lotz et al. 2006, and others). Conselice (2006a) estimates that massive galaxies
have undergone about 4 major mergers by redshift 1. Toomre (1977) described a sequence
of merger activity ranging from separated galaxies with tails and a bridge between them,
to double nuclei in a common envelope with tails, to merged nuclei with tails. Ground-
based (Hibbard & van Gorkom 1996) and space-based (Laine et al. 2003; Smith et al. 2007)
observations of this sequence show optical, infrared, and radio activity in the tails and nuclei.
High resolution images and numerical simulations of nearby interactions demonstrate
how star formation and morphology are affected. General reviews of interaction simulations
are given by Barnes & Hernquist (1992) and Struck (1999). The initial galaxy properties,
such as mass, rotational velocity, gas content and dark matter content, and their initial sep-
arations and velocity vectors, all play a role in generating structure. The viewing angle also
affects the morphology. Early-type galaxies with little gas are expected to display smooth
plumes and shells, while spiral interactions and mergers should exhibit clumpy star forma-
tion along tidal tails, and condensations of material at the tail ends. Equal mass companions
may show bridges between them. A prominent example of a tidal interaction is the Antennae
(NGC4038/9), a merging pair of disk galaxies with rampant star formation in the central
regions, including young globular clusters (Whitmore et al. 2005). Its interaction was first
modeled by Toomre & Toomre (1972). The Cartwheel galaxy is a collisional ring system
rimmed with star formation from a head-on collision (Struck et al. 1996). Sometimes polar-
ring or spindle galaxies are the result of perpendicular collisions (Struck 1999). The Mice
(NGC 4676) has a long narrow straight tail and a curved tidal arm (Vorontsov-Velyaminov
1957; Burbidge & Burbidge 1959); numerical simulations reproduce both features well in a
model with a halo:(disk+bulge) mass ratio of 5 (Barnes 2004). The Superantennae (IRAS
19254-7245) is a pair of infrared-luminous merging giant galaxies having Seyfert and star-
– 3 –
burst nuclei and ∼ 200 kpc tails with a tidal tail dwarf (Mirabel, Lutz, & Maza, 1991). The
Leo Triplet includes NGC 3628 with an 80 kpc stellar tail containing star-forming complexes
with masses up to 106 M⊙ (Chromey et al. 1998). The Tadpole galaxy UGC10214 (Tran et
al. 2003; de Grijs et al. 2003; Jarrett et al. 2006), the IC2163/NGC2207 pair (Elmegreen
et al. 2001, 2006), and Arp 107 (Smith et al. 2005) are all interacting systems observed
with HST and SST and modeled in simulations. Many local mergers have intense nuclear
activity, such as the Seyfert galaxy NGC 5548, which also has an 80 kpc long, low surface
brightness (V=27-28 mag arcsec−2) tidal tail and a 1-arm diffuse spiral (Tyson et al. 1998).
The GEMS (Galaxy Evolution from Morphology and SEDs; Rix et al. 2004), GOODS
(Great Observatories Origins Deep Survey; Giavalisco et al. 2004), and UDF (Ultra Deep
Field; Beckwith et al. 2006) surveys done with the HST ACS (Hubble Space Telescope
Advanced Camera for Surveys) have enabled high resolution studies of the morphology of
intermediate and high redshift galaxies. Light distribution parameters such as the Gini co-
efficient (Lotz et al. 2006) and concentration index, asymmetry, and clumpiness (CAS;
Conselice 2006) have been applied to galaxies in these fields to study possible merger
systems. For GEMS and GOODS, John Caldwell of the GEMS team has posted images
(archive.stsci.edu/prepds/gems/datalist.html) of several galaxies from each field, including
peculiar and interacting systems with tails and bridges. Here we examine the entire GEMS
and GOODS fields systematically for such galaxies and study their tails, bridges, and star-
forming regions. Their properties are useful for understanding interactions and interaction-
triggered star formation, and for probing the relative dark matter content (e.g., Dubinski,
Mihos, & Hernquist 1999).
2. The Sample of Interactions and Mergers
The GOODS and GEMS images from the public archive were used for this study. They
include exposures in 4 filters for GOODS: F435W (B435), F606W (V606), F775W (i775), and
F850LP (z850); and 2 filters (V606 and z850) for GEMS. The public images were drizzled to
produce final archival images with a scale of 0.03 arcsec per px. GEMS, which incorporates
the southern GOODS survey (Chandra Deep Field South, CDF-S) in the central quarter of
its field, covers 28 arcmin x 28 arcmin; there are 63 GEMS and 18 GOODS images that
make up the whole field. The GOODS images have a limiting AB mag of V606= 27.5 for
an extended object, or about two mags fainter than the GEMS images. There are over
25,000 galaxies catalogued in the COMBO-17 survey (Classifying Objects by Medium-Band
Observations, a spectrophotometric 17-filter survey; Wolf et al. 2003), and 8565 that are
cross-correlated with the GEMS survey (Caldwell et al. 2005).
– 4 –
Interacting galaxies with tails, bridges, diffuse plumes and other features were identified
by eye on the online Skywalker images and examined on high resolution V606 fits images.
The lower limit to the length of detectable tails is about 20 pixels. Snapshots of several
different morphologies for interacting galaxies are shown in Figures 1-6. Out of an initial
list of about 300 galaxies, a total of 100 best cases are included in our sample: 14 diffuse
types, 18 antennae types, 22 M51 types, 19 shrimp types, 15 equal mass interactions, and
12 assemblies, as we describe below.
GEMS and GOODS galaxy redshifts were obtained from the COMBO-17 list (Wolf et
al. 2003). Our sample ranges from redshift z = 0.1 to 1.4 in an area of 2.8×106 square arcsec.
The linear diameters of the central objects were determined from their angular diameters
and redshifts using the appropriate conversion for a ΛCDM cosmology (Carroll et al., 1992;
Spergel et al., 2003). The range is ∼ 3 to 33 kpc. Projected tail lengths were measured in
a straight line from the galaxy center to the 2σ noise limit (25.0 mag arcsec−2) in the outer
tail.
Photometry was done on the whole galaxies, on each prominent star-forming clump, and
on the tails using the IRAF task imexam. A box of variable size was defined around each
feature; the outer limits of the boxes were chosen to be where the clump brightness is about
3 times the surrounding region. Sky subtraction was not done because the background
is negligible. The photometric errors are ∼0.1-0.2 mag for individual clumps. The V606
surface brightnesses of the tidal tails were determined using imagej (Rasband 1997) to trace
freehand contours around the tails, so that they could be better defined than with rectangular
or circular apertures.
Figure 1 shows galaxies with diffuse plumes and either no blue star formation patches
or only a few tiny patches (e.g., galaxies number 5 and 6); we refer to these interactions as
diffuse types. The colors of the plumes match the colors of the outer parts of the central
galaxies, indicating the plumes are tidally shorn stars with little gas. There is structure in
most of the plumes consisting of arcs or sharp edges. This is presumably tidal debris from
early type galaxies with little or no gas (e.g. Larson & Tinsley 1978; Malin & Carter 1980;
Schombert et al. 1990). This type of interaction is relatively rare in the GEMS and GOODS
images, perhaps because the tidal debris is faint. The best cases are shown here and they
all have relatively small redshifts compared to the other interaction types (the average z is
0.23 and the maximum z is 0.69).
The image in the top left panel of Figure 1 (galaxy 1) has a giant diffuse clump in the
upper right corner. This could be a condensation in the tidal arm, or it could be another
galaxy. In either case, it has the same color as the rest of the tidal arm nearby. That is,
V606 − z850 = 0.90 ± 0.5 for the clump and also in six places along the tail; the color is
– 5 –
essentially the same, 0.94 ± 0.05, in the core of the galaxy. The absolute magnitude of the
clump is MV = −18.41 for redshift z = 0.15. The mass is ∼ 5 × 10
9 M⊙ (Sect. 3.2). If
this clump is a condensation in the tail, then it could be a rare case where a pure stellar arc
has collapsed gravitationally into a gas-free tidal dwarf. The final result could be a dwarf
elliptical. Usually tidal dwarfs form by gaseous condensations in tidal arms (Wetzstein,
Naab, & Burkert 2007).
Figure 2 shows interactions that resemble the local Antennae pair, so we refer to them as
antennae types. These types have long tidal tails and double nuclei or highly distorted centers
that appear to be mergers of disk galaxies. Note that antennae are not the same as “tadpole”
galaxies (Elmegreen et al. 2005a; de Mello et al. 2006; Straughn et al. 2006), which have one
main clump and a sometimes wiggly tail that may contain smaller clumps. Some antennae
have giant clumps near the ends of the tails which could have formed there (galaxies 16 and
17) and are analogous to the clump at the end of the Superantennae (Mirabel et al. 1991).
Galaxy 18 is in a crowded field with at least two long tidal arms; here we consider only the
tail system in the north, which is in the upper part of the figure. These long-tail systems are
relatively rare and all the best cases are shown in the figure; their average redshift, 0.70, is
typical for GEMS and GOODS fields. Galaxy 24 is somewhat like a tadpole galaxy, but its
very narrow tail and protrusion on the anti-tail side of the main clump are unlike structures
seen in tadpoles of the Ultra Deep Field.
For the antennae galaxies in Figure 2, the tails have an average (V-z) color that is
negligibly bluer, 0.10±0.25 mag, than the central disks. In a study of tidal features in local
Arp atlas galaxies, Schombert et al. (1990) also found that the tail colors are uniform and
similar to those of the outer disks. They noted that the most sharply-defined tails are with
spiral systems and the diffuses plumes are with ellipticals. This correlation may be true
here also, but it is difficult to tell from Figure 1 whether the smooth distorted systems are
intrinsically disk-like.
Galaxy 20 in Figure 2 is an interesting case. It has an elliptical clump at the end of its
tail that could be one of the collision partners. There are two central galaxy cores, however,
and their interaction may have formed the tidal arms without this companion. Furthermore,
the clump at the tip is aligned perpendicular to the tail, which is unusual for a tidal dwarf.
Thus it is possible that the clump was a pre-existing galaxy lying in the orbital plane of one
of the larger galaxies now at the center. Presumably this former host is the galaxy currently
connected to the dwarf by the tidal arm. The interaction could have swung it around to
its current position at the tip. A similar case occurs for the local IC 2163/NGC 2207 pair,
which has a spheroidal dwarf galaxy at the tip of its tidal arm (Elmegreen et al. 2001). Such
swing-around dwarfs should have the same dynamical origin as the large pools of gas and
– 6 –
star formation that are at the tips of superantenna-type galaxies; i.e. the whole outer disk
moves to this position during the interaction (Elmegreen et al. 1993; Duc et al. 1997).
Figure 3 shows examples of interactions that we refer to as M51-type galaxies, where
the tidal arms can be bridges that connect the main disk galaxy to the companion (galaxy
33), or tails on the opposite side of the companion (e.g., galaxies 34 and 35), or both (galaxy
36). In galaxy 44, the tidal arm looks like the debris path of a pre-existing galaxy that lies at
the right; the orbit path apparently curves around on the left. The M51-types usually have
strong spirals in the main disk. In the top row, the tails and bridges are thin and diffuse.
The galaxy on the left in the lower row (galaxy 42) has a thick, fan-shaped tail opposite
the companion. Some bridges have star formation clumps (galaxy 40) and others appear
smooth (galaxy 33). Interactions like this, especially those with small companions, are more
common than the previous two types and only a few best cases are shown in Figure 3 and
discussed in the rest of this paper.
Figure 4 shows examples of galaxies dominated by one highly curved, dominant arm and
large, regularly-spaced clumps of star formation. We call these “shrimp” galaxies because
of their resemblance to the tail of a shrimp. Although their star formation indicates they
contain gas and therefore are disk systems, there are no well-defined spirals (except for the
prominent arm), merging cores, or obvious central nuclei. The clumps resemble the beads-
on-a-string star formation in spiral density waves and probably have the same origin, a
gravitational instability (Elmegreen & Elmegreen 1983; Kim & Ostriker 2006; Bournaud,
Duc, & Masset 2003). The J-shaped morphology is reminiscent of the 90 kpc gas tail of M51
(Rots et al. 1990) and the 48 kpc gas tail observed in NGC 2535 (Kaufman et al. 1997).
Rots el al. point out that the M51 gas tail is much broader (10 kpc) than the narrow tails
seen in merging systems like the Antennae. The broad tail in galaxy 42 (Fig. 3) is similar
to the M51 tail. Sometimes there is a bright tail with no obvious companion (galaxies 56,
57, and 60); one of these, galaxy 56, was in our ring galaxy study (Elmegreen & Elmegreen
2006). Asymmetric, strong arm galaxies like this are not common in GEMS and GOODS;
this figure shows the best cases.
Figure 5 has a selection of irregular galaxies that appear to be interactions. Most of
them suggest an assembly of small pieces, so we refer to them as assembly types. If they were
slightly more round in overall shape, with more obvious interclump emission, then we would
classify them as clump-clusters, as we did in the UDF (Elmegreen et al. 2005a). The galaxy
in the lower left (galaxy 83) is like this. The resemblance of these types to clump-clusters
suggests that some of the clumps are accreted from outside the disk and others form from
gravitational instabilities in a pre-existing gas disk, as suggested previously (Elmegreen &
Elmegreen 2005). The system in the lower right (galaxy 85) could be interacting spirals, or a
– 7 –
triple system, or a bent chain (as studied in Elmegreen & Elmegreen 2006). There are many
examples of highly irregular galaxies like these in the GEMS and GOODS fields; indeed most
galaxies at z > 1.5 are peculiar in this sense (Conselice 2005). In what follows, we discuss
only these 12 galaxies.
Figure 6 has samples of grazing or close interactions, with spirals at the top of the
page (numbers 86-93), ellipticals lower down (numbers 95-97) and two polar-ring galaxies
(numbers 99 and 100) in the lowest row at the middle and right bottom. We refer to these
paired systems as “equals” because their distinguishing feature is that the two galaxies have
comparable size. The pair number 89 has a bright oval in the smaller galaxy, which is char-
acteristic of recent tidal forces for an in-plane, prograde encounter such as IC2163/NGC2207
(Sundin 1993; Elmegreen et al. 1995 ). There is a spiral-elliptical pair on the right in the
middle row (galaxy 94). Double ellipticals in the UDF were studied previously (Elmegreen
et al. 2005a, Coe et al. 2006). Near neighbors like this have been studied previously in the
GEMS field; 6 double systems out of 379 red sequence galaxies were identified as being dry
merger candidates, as reproduced in simulations (Bell et al. 2006). The models of mergers
of early-type systems by Naab et al. (2006) apparently account for kinematic and isophotal
properties of ellipticals better than the formation of ellipticals through late-type mergers
alone. For the pairs in our figure, both components have the same COMBO17 redshift.
There are many other examples of close galaxy groups and near interactions in the GEMS
and GOODS surveys. In what follows we discuss only the properties of those shown in Figure
The interacting types shown in the figures are meant to be as distinct as possible. These
and other good cases are listed in Table 1 by running number, along with their COMBO-
17 catalog number, redshift, and R magnitude. There is occasionally some ambiguity and
overlap in the interaction types, particularly between M51-types and shrimps when the M51-
types have small or uncertain companions at the ends of their prominent tails. Projection
effects can lead to uncertainties in the classifications as well, particularly for antennae whose
tails may be foreshortened. Nevertheless, these divisions serve as a useful attempt to sort out
the most prominent features among interacting galaxies. There are numerous other galaxies
in GEMS and GOODS that are apparently interacting, but most of them are too highly
distorted to indicate the particular physical properties of interest here, namely, disk-to-halo
mass ratio and star formation scale.
– 8 –
3. Photometric Results
3.1. Global galaxy properties
The integrated Johnson restframe (U-B) and (B-V) colors from COMBO-17 for the
observed galaxies with measured redshifts are shown in a color-color diagram in Figure 7.
The crosses in the diagram are Johnson colors for standard Hubble types (Fukugita et al.
1995). Our sample of galaxies spans the range of colors from early to late Hubble types,
although the bluest are bluer than standard irregular galaxies (a typical Im has U-B= −0.35,
B-V= 0.27). The reddest galaxies tend to be the diffuse types, thought to originate with
ellipticals involved in interactions. The two reddest galaxies in our sample are the diffuse
types number 1 and 2 in Figure 1. The bluest tend to be the assemblies, consistent with
their having formed recently.
Figure 8 shows a restframe color-magnitude diagram. Early and late type galaxies
usually separate into a “red sequence” and a “blue cloud” on such a diagram (Baldry et
al. 2004; Faber et al. 2005). The solid line indicates the boundary between these two
regions from a study of 22,000 nearby galaxies (Conselice 2006b). The short-dashed lines
are the limits of the Conselice (2006) survey; local galaxies are brighter than the vertical
short-dashed line and their colors lie between the horizontal short-dashed lines. The long-
dashed lines approximately outline the bright limit for the local blue cloud galaxies. Our
galaxies fall in both the red sequence and the blue cloud. The restframe colors in Figure 8
are consistent with their morphological appearances. The red sequence galaxies in the figure
usually appear smooth (the diffuse types) or lack obvious huge star formation clumps (the
equal mass mergers), while the blue cloud galaxies usually have patches of star formation
(the M51-types, shrimps, assemblies, and many antennae). We see now why the redshifts of
the diffuse galaxies (z < 0.3) are much lower than the others: this is a selection effect for the
ACS camera. These tails comprise old stellar populations without star-forming clumps, and
their intrinsic redness makes them difficult to see at high redshifts. Also, they tend to have
intrinsically low surface brightnesses because of a lack of star formation, and cosmological
dimming makes them too faint to see at high redshift. Hibbard & Vacca (1997) note that it
is difficult to detect tidal arms beyond z ∼ 1.5.
3.2. Clump properties
Prominent star-forming clumps are apparent in many of the interacting galaxies. Their
sizes and magnitudes were measured using rectangular apertures. The observed magnitudes
were converted to restframe B magnitudes whenever possible, using linear interpolations
– 9 –
between the ACS bands. For example, GEMS observations are at two filters, V606 and z850.
GEMS galaxies with redshifts z between 0.39 (= 606/435 − 1) and 0.95 (= 850/435 − 1)
were assumed to have restframe blue luminosities given by LB,rest = LV,obs(0.95− z)/(0.95−
0.39) + Lz,obs(z − 0.39)/(0.95 − 0.39). The restframe B magnitude is then −2.5 logLB,rest.
For GOODS galaxies, the conversions were divided into 3 redshift bins to make use of
the 4 available filters, and a linear interpolation was again applied to get restframe clump
magnitudes. For the GOODS galaxies, the restframe magnitudes determined by interpolation
between the nearest 2 filters among the 4 filters are within ±0.2 mag of the restframe
magnitudes determined from only the V and z filters. Thus, the GEMS interpolations are
accurate to this level. (We do not include corrections for intergalactic absorption in these
colors, because we are comparing them directly with their parent galaxy properties. Below,
when we convert the colors and magnitudes to masses and ages, absorption corrections are
taken into account.)
The apparent restframe B magnitudes of the clumps were converted to absolute rest-
frame B magnitudes using photometric redshifts and the distance modulus for a ΛCDM
cosmology. These absolute clump magnitudes are shown as a function of absolute galaxy
magnitude in Figure 9. The clump absolute B magnitudes scale linearly with the galaxy
magnitudes. The clumps are typically a kpc in size (∼ 3 to 8 pixels across), comparable
to star-forming complexes in local galaxies (Efremov 1995), which also scale with galaxy
magnitude (Elmegreen et al. 1996; Elmegreen & Salzer 1999).
Clump ages and masses were estimated by comparing observed clump colors, magni-
tudes, and redshifts with evolutionary models that account for bandshifting and intergalac-
tic absorption and that assume an exponential star formation rate decay (see Elmegreen &
Elmegreen 2005). Internal dust extinction as a function of redshift is taken from Rowan-
Robinson (2003). The GEMS galaxy clumps only have (V606-z850) colors, so the ages are not
well constrained. For the GOODS galaxies, the additional B and I filters help place better
limits on the ages, although there is still a wide range of possible fits.
Figure 10 shows sample model results for redshift z = 1. The different lines in each
panel correspond to different decay times for the star formation rate, in years: 107, 3× 107,
108, 3×108, and 109, and the sixth line represents a constant rate. Generally the shorter the
decay time, the redder the color and higher the mass for a given duration of star formation.
This correspondence between color and mass gives a degeneracy to plots of mass versus color
at a fixed apparent magnitude (top left) and apparent magnitude versus color at a fixed mass
(top right). Thus the masses of clumps can be derived approximately from their colors and
magnitudes, without needing to know their ages or star formation histories.
Figure 11 shows observations and models in the color-magnitude plane for 6 redshift in-
– 10 –
tervals spanning our galaxies. Each curve represents a wide range of star formation durations
that vary along the curve as in the top right panel of Fig. 10; each curve in a set of curves is
a different decay time. The different sets of curves, shifted vertically in the plots, correspond
to different clump masses, as indicated by the adjacent numbers, which are in M⊙. Each
different point is a different clump; many galaxies have several points. Only clumps with
both V606 and z850 magnitudes above the 2σ noise limit are plotted in Figure 11. The clump
(V606 − z850) colors range from 0 to 1.5. The magnitudes tend to be about constant for each
redshift because of a selection effect (brighter magnitudes are rare and fainter magnitudes
are not observed).
Figure 11 indicates that the masses of the observable clumps are between 106 and 109
M⊙ for all redshifts, with higher masses selected for the higher redshifts. The masses for all
of the clumps are plotted in Figure 12 versus the galaxy type (types 1 through 6 are in order
of Figs. 1 through 6 above). The masses are obtained from the observed values of V606 and
V606 − z850 using the method indicated in Figure 11. The different mass evaluations for the
six decay times are averaged together in the log to give the log of the mass plotted as a dot
in Figure 12. The rms values of log-mass among these six evaluations are shown in Figure
12 as plus-symbols, using the right-hand axes. These rms deviations are less than 0.2, so
the uncertainties in star formation decay times and clump ages do not lead to significant
uncertainties in the clump mass. (Systematic uncertainties involving extinctions, stellar
evolution models, photometric redshifts, and so on, would be larger.)
The clump ages cannot be determined independently from the star formation decay
times with only the few passbands available at high angular resolution. Figure 13 shows
model results that help estimate the clump ages. As in the other figures, each line is a
different exponential decay time for the star formation rate. If we consider the two extreme
decay times in this figure (continuous star formation for the bottom lines in each panel and
107 years for the top lines), then we can estimate the age range for each decay time from
the observed color range. For V606 − z850 colors in the range from 0 to 0.5 at low z (cf. Fig
11), the clump ages range from 107 to 1010 yr with continuous star formation and from 107
to 3 × 108 yr with a decay time of 107 yrs. For colors in the range from 0 to 1.5 at higher
redshifts, the age ranges are about the same in each case. For intermediate decay times, the
typical clump ages are between ∼ 107 years for the bluest clumps and ∼ 109 years for the
reddest clumps. These are reasonable ages for star formation regions, and consistent with
model tail lifetimes.
The star-forming complexes in the GEMS and GOODS interacting galaxies are 10 to
1000 times more massive than the local analogs seen in non-interacting late-type galaxies
(Elmegreen & Salzer 1999), but the low mass end in the present sample is similar to the
– 11 –
high mass end of the complexes measured in local interacting galaxies. For example, the
Tadpole galaxy, UGC 10214, contains 106 M⊙ complexes along the tidal arm (Tran et al.
2003; Jarrett et al. 2006). The interacting galaxy NGC 6872 has tidal tails with 109 M⊙ HI
condensations (Horellou & Koribalski 2007), but the star clusters have masses only up to
106 M⊙ (Bastian et al. 2005). The most massive complexes in the tidal tail of NGC 3628
in the Leo Triplet are also ∼ 106 M⊙ (Chromey et al. 1998). The NGC 6872 clusters differ
qualitatively from those in our sample in being spread out along a narrow arm; ours are big
round clumps spaced somewhat evenly along the arm. Small star clusters are also scattered
along the tidal arms the Tadpole and Mice systems; they typically contain less than 106 M⊙
(de Grijs et al. 2003). The NGC 3628 clusters are also faint with surface brightnesses less
than 27 mag arcsec−2; they would not stand out at high redshift.
It is reasonable to consider whether the observed increase of complex mass with increas-
ing redshift is a selection effect. Our clumps are several pixels in size, corresponding to a
scale of ∼ 1 kpc. Individual clusters are not resolved and we only sample the most massive
conglomerates. These kpc sizes are comparable to the complex sizes in local galaxies, but the
high redshift complexes are much brighter and more massive. They would be observed easily
in local galaxies. The massive complexes in our sample are more similar to those measured
generally in UDF galaxies (Elmegreen & Elmegreen 2005).
Clump separations were measured for clumps along the long arms in the shrimp galaxies
of Figure 4. They average 2.20±0.94 kpc for 49 separations. This is about the same separa-
tion as that for the largest complexes in the spiral arms of local spiral galaxies (Elmegreen
& Elmegreen 1983, 1987), and comparable to the spacing between groups of dust-feathers
studied by La Vigne et al. (2006). Yet the clumps in shrimp galaxies and others studied here
are much more massive than the complexes in local spiral arms, which are typically < 106
M⊙ in stars and ∼ 10
7 M⊙ in gas. This elevated mass can be explained by a heightened
turbulent speed for the gas, combined with an elevated gas density. Considering that the
separation is about equal to the two-dimensional Jeans length, λ ∼ 2a2/ (GΣ) for velocity
dispersion a and mass column density Σ, and that the mass is the Jeans mass, λ2Σ, the
mass scales with the square of the velocity dispersion, M = M0 (a/a0)
for fixed length
λ0 = 2a
0/ (GΣ0) and M0 = λ
0Σ0. The mass column density also scales with the square of
the dispersion, Σ = Σ0 (a/a0)
to keep λ constant. Thus the interacting tidal arm clumps
are massive because the velocity dispersions and column densities are high. Another way
to derive this is to note that for regular spiral arm instabilities, 2Gµ/a2 is about unity at
the instability threshold, where µ is the mass/length along the arm (Elmegreen 1994). Thus
cloud mass scales with a2 for constant cloud separation. High velocity dispersions for neutral
hydrogen, ∼ 50 km s−1, are also observed in local interacting galaxies (Elmegreen et al. 1993;
Irwin 1994; Elmegreen et al. 1995; Kaufman et al. 1997; Kaufman et al. 1999; Kaufman
– 12 –
et al. 2002). Presumably the interaction agitates the interstellar medium to make the large
velocity dispersions. The orbital motions are forced to be non-circular and then the gaseous
orbits cross, converting orbital energy into turbulent energy and shocks. Similar evidence
for high velocity dispersions was found in the masses and spacings of star forming complexes
in clump cluster galaxies (Elmegreen & Elmegreen 2005) and in spectral line widths (Genzel
et al. 2006; Weiner et al. 2006).
3.3. Tail Properties
Figure 14 shows the average tail surface brightness as a function of (1+ z)4 for galaxies
in Figures 1-4. Some systems have more than one tail. Cosmological dimming causes a fixed
surface brightness to get fainter as (1+z)−4, so there should be an inverse correlation in this
diagram. Clearly, the tails are brighter for the more nearby galaxies, and they decrease out
to z ∼ 1, where they are fairly constant. This constant limit is at the 2σ detection limit of
25 mag arcsec−2. Antennae galaxies with average tail surface brightnesses fainter than this
limit have patchy tails with no apparent emission between the patches. Only the brightest
high redshift tails can be observed in this survey.
Simulations by Mihos (1995) suggested that tidal tails are observable for a brief time in
the early stages of a merger, corresponding to ∼ 150 Myr at a redshift z = 1 and 350 Myr
at z = 0.4. The difference is the result of surface brightness dimming as tails disperse. A
nearby galaxy merger, Arp 299, has a 180 kpc long tail encompassing 2 to 4% of the total
galaxy luminosity, with an interaction age of 750 Myr, but its low surface brightness of 28.5
mag arcsec−2 (Hibbard & Yun 1999) would be below the GOODS/GEMS detection limit.
The ratio of the luminosity of the combined tails and bridges to the luminosity of the
disk (the luminosity fraction) is shown in Figure 15. The luminosity fraction in the tidal
debris ranges from 10% to 80%, averaging about 30% regardless of redshift. This range is
consistent with that of local galaxies in the Arp atlas and Toomre sequence (e.g., Schombert,
Wallin, & Struck-Marcell 1990; Hibbard & van Gorkom 1996).
Interaction models with curled tails, as in our shrimp galaxies, were made by Bournaud
et al. (2003). Their models had dark matter halos with masses ∼ 10 times the disk mass
and extents less than 12 disk scale lengths. Some of our shrimp galaxies have one prominent
curved arm that is pulled out from the main disk but not very far, resulting in a lopsided
galaxy. Simulations indicate that such lopsidedness may be the result of a recent minor
merger (Bournaud et al. 2005). In some of our cases, a nearby companion is obvious.
The linear sizes of the tidal tails in our sample are shown in Figure 16. They range from
– 13 –
2 to 60 kpc, and are typically a few times the disk diameter, as shown in Figure 17, which
plots this ratio versus redshift. The average tail to diameter ratio is 2.9±1.7 for diffuse tails,
2.5± 1.3 for antennae, 2.5± 1.1 for M51-types and 1.5± 1.4 for shrimps, so the shrimps are
about 60% as extended as the antennae types. There is no apparent dependence of these
ratios on redshift in Figure 17. Projection effects make these apparent ratios smaller than
the intrinsic ratios.
For comparison, the ratio of tail length to disk diameter versus the tail length for local
galaxies is shown in Figure 18 based on measurements of antennae-type systems in the Arp
atlas (1966) and the Vorontsov-Velyaminov atlas (1959). Our galaxies are also shown. The
average tail length for the local galaxies in this figure is 72 ± 48 kpc, while the average tail
length for the GEMS and GOODS antennae is 37% as much, 27±16 kpc. The diameters for
these two groups are 20±12 kpc and 11±5 kpc, and the ratios of tail length to diameter are
4.5± 3.7 and 2.5± 1.3, respectively. Thus the local antennae mergers are larger in diameter
by a factor of 2 than the GEMS and GOODS antennae, and the tails for the locals are larger
by a factor of 2.7. These results for the diameters are consistent with other indicators that
galaxies are smaller at higher redshift, although usually this change does not show up until
z > 1 (see observations and literature review in Elmegreen et al. 2007).
3.4. Tidal dwarf galaxy candidates
Three antennae galaxies at the top of Figure 2, numbers 15, 16, and 18, have long
straight tidal arms with large star-forming regions at the ends. These clumps are possibly
tidal dwarf galaxies. The clump diameters and restframe B magnitudes are listed in Table
2, along with the clump in diffuse galaxy number 1 discussed in Sect. 2. Listed are their
V606 and V606 − z850 magnitudes and associated masses, calculated as in Sect. 3.2. The
masses range from 0.2× 108 to 4.6× 108 M⊙ for the star-forming dwarfs, but for the stellar
condensation in the diffuse-tail galaxy 1 (Fig. 1), the mass is 50×108 M⊙. The star-forming
dwarf masses are similar to or larger than those found for the tidal object at the end of the
Superantennae (Mirabel et al. 1991) as well as the tidal object at the end of the tidal arm in
the IC 2163/NGC 2207 interaction (Elmegreen et al. 2001) and at the end of the Antennae
tail (Mirabel et al. 1992). The HI dynamical masses for these local tidal dwarfs are ∼ 109
Simulations of interacting galaxies that form tidal dwarf galaxies require long tails and
a dark matter halo that extends a factor of 10 beyond the optical disk (Bournaud et al.
2003). If one or both galaxies contain an extended gas disk before the interaction, then more
massive, 109 M⊙ stellar objects can form at the tips of the tidal arms from the accumulated
– 14 –
pool of outer disk material (Elmegreen et al. 1993; Bournaud et al. 2003). Observations of
nearby interactions show clumpy regions of tidal condensations with masses of ∼ 108 − 109
M⊙ (Bournaud et al. 2004; Weilbacher et al. 2002, 2003; Knierman et al. 2003; Iglesias-
Paramo & Vilchez 2001), like what is observed in our high redshift tidal dwarfs.
No well-resolved models have yet formed tidal dwarfs from stellar debris. Wetzstein,
Naab, & Burkert (2007) considered this possibility and found collapsing gas more likely. Yet
the condensed object in the tail of galaxy 1 could have formed there and it is interesting
to consider whether the Jeans mass in such an environment is comparable to the observed
mass. If, for example, the tidal arm surface density corresponds to a value typical for the
outer parts of disks, ∼ 10 M⊙ pc
−2, and the stellar velocity dispersion is comparable to that
required in Sect. 3.2 for the gas to give the giant star forming regions, ∼ 40 km s−1, then the
Jeans mass is M ∼ a4/ (G2Σ) ∼ 1010 M⊙. This is not far from the value we observe, 5× 10
M⊙, so the diffuse clump could have formed by self-gravitational collapse of tidal tail stars.
The timescale for the collapse would be a/ (πGΣ) ∼ 300 Myr, which is not unreasonable
considering that the orbit time at this galactocentric radius is at least this large.
4. Dark Matter Halo Constraints
Models of interacting galaxies have been used to place constraints on dark halo poten-
tials. Springel & White (1999) and Dubinski, Mihos, & Hernquist (1999) found that tidal
tail lengths can be long compared to the disk if the ratio of escape speed to rotation speed
at 2 disk scale lengths is small, ve/Vr < 2.5, and the rotation curve is falling in the outer
disk. In a series of models, Dubinski et al. showed that this condition may result from either
disk-dominated rotation curves where the halo is extended and has a low concentration, or
halo-dominated rotation curves where the halo is compact and low mass. Dubinski et al.
point out that the latter possibility is inconsistent with observed flat or rising disk rotation
curves, but the first is compatible if the disk is massive and dominant in the inner regions.
The first case also gives prominent bridges. In addition, Springel & White (1999) found that
CDM halo models with embedded disks allow long tidal tails, but Dubinski et al. noted that
most of those which do are essentially low surface brightness disks in massive halos, and not
normal bright galaxies. Galaxies without dark matter halos are not capable of generating
long tidal tails (Barnes 1988). In all cases, longer tails develop in prograde interactions.
The smooth diffuse types and antenna types in Figures 1 and 2 have relatively long tails,
so the progenitors were presumably disks of early and late types, respectively, with falling
rotation curves in their outer parts. These long-tail cases are relatively rare, comprising
only about 8% and 9%, respectively, of our original (300 galaxy) interacting sample from
– 15 –
GEMS and GOODS. The more compact M51 types and shrimps represent 9% and 12% of
the sample. Short tail interactions could be younger, less favorably projected, or have a
more steeply rising rotation curve than long tail interactions. The M51 types have clear
companions, so the prominent features are bridges. According to Dubinski et al. (1999),
bridging requires a prograde interaction with a maximum-disk galaxy, that is, one with a
low-mass, extended halo.
5. Conclusions
Mergers and interactions out to redshift z = 1.4 have tails, bridges, and plumes that are
analogous to features in local interacting galaxies. Some interactions have only smooth and
red features, indicative of gas-free progenitors, while others have giant blue star-formation
clumps. The tail luminosity fraction has a wide range, comparable to that found locally.
A striking difference arises regarding the tail lengths, however. The tails in our antenna
sample, at an average redshift of 0.7, are only one-third as long as the tails in local antenna
mergers, and the disk diameters are about half the local merger diameters. This difference is
consistent with the observations that high redshift galaxies are smaller than local galaxies,
although such a drop in size has not yet been seen for galaxies at redshifts this low. The
implication is that dark matter halos have not built up to their full sizes for typical galaxies
in GEMS and GOODS.
Star formation is strongly triggered by the interactions observed here, as it is locally.
The star-forming clumps tend to be much more massive than their local analogs, however,
with masses between ∼ 106 M⊙ and a few ×10
8 M⊙, increasing with redshift. This is not
merely a selection effect, since the massive clumps seen at high redshift would show up at
lower redshift, although of course smaller clumps would not be resolved at high redshift. The
clump spacings were measured along the tidal arms of the most prominent one-arm type of
interaction, the shrimp-type, and found to be 2.20 ± 0.94, which is typical for the spacing
between beads on a string of star formation in local spiral arms. If both types of arms form
clumps by gravitational instabilities, then the turbulent speed of the interstellar medium in
the GEMS and GOODS sample has to be larger than it is locally by a factor of ∼ 5 or more;
the gas mass column density has to be larger by this factor squared.
Some interactions have tidal dwarf galaxies at the ends of their tidal arms, similar to
those found in the Superantennae galaxy and other local mergers. One diffuse interaction
with red stellar tidal debris has a large stellar clump that may have formed by gravitational
collapse in a stellar tidal arm; the clump mass is 5 × 109 M⊙. Long-arm interactions are
relatively rare, comprising only ∼ 17% of our total sample of ∼ 300 interacting systems (only
– 16 –
a fraction of which were discussed here). For those with long arms, numerical models suggest
the dark matter halos must be extended, so that the rotation curves are falling in the outer
disks. Most interactions are not like this, however, so the rotation curves are probably still
rising in their outer disks, like most galaxies locally.
We gratefully acknowledge summer student support for B.M. and T.F. through an REU
grant for the Keck Northeast Astronomy Consortium from the National Science Foundation
(AST-0353997) and from the Vassar URSI (Undergraduate Research Summer Institute) pro-
gram. D.M.E. thanks Vassar for publication support through a Research Grant. We thank
the referee for useful comments. This research has made use of the NASA/IPAC Extragalac-
tic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute
of Technology, under contract with the National Aeronautics and Space Administration.
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This preprint was prepared with the AAS LATEX macros v5.2.
– 21 –
Table 1. Interacting Galaxies in GEMS and GOODS
Type, Figure Number COMBO 17 z R mag.
Diffuse (Fig. 1) 1 6423 0.15 16.572
2 12639 0.154 16.678
3 11538 0.134 17.713
4 53129 0.171 16.968
5 57881 0.118 17.552
6 28509 0.093 18.79
7 17207 0.69 19.742
8 30824 0.341 19.755
9 25874 0.262 19.757
Diffuse (other) 10 22588 0.684 21.263
11 21990 0.429 21.243
12 46898 0.617 20.794
13 49709 0.302 20.23
14 15233 0.304 18.857
Antennae (Fig. 2) 15 61546 0.552 20.41
16 45115 0.579 21.275
17 20280 0.555 21.653
18 41907 0.702 22.66
19 35611 1.256 22.655
20 10548 0.698 22.43
21 33650 0.169 18.86
22 42890 0.421 20.68
23 49860 1.169 23.632
24 34926 0.779 -19.69
Antennae (other) 25 14829 0.219 21.429
26 18588 0.814 22.748
27 46738 1.204 20.65
28 7551 1.162 25.926
29 20034 1.326 21.932
30 33267 0.067 23.112
31 38651 0.988 23.89
32 55495 1.00 24.261
M51-type (Fig. 3) 33 5640 0.204 19.477
34 9415 0.523 21.16
35 40901 0.193 19.751
36 17522 0.82 23.103
37 6209 1.187 22.723
38 23667 1.151 23.514
39 37293 0.274 20.533
40 39805 0.557 20.089
41 53243 0.698 21.683
42 15599 0.56 21.381
43 25783 0.663 20.732
44 39228 0.117 18.031
M51-type (other) 45 1984 0.762 22.855
– 22 –
Table 1—Continued
Type, Figure Number COMBO 17 z R mag.
46 2760 1.281 23.202
47 15040 0.667 22.392
48 18502 0.228 21.942
49 14959 0.306 19.581
50 16023 0.668 21.887
51 30226 0.509 22.689
52 40744 0.292 21.119
53 45102 0.857 22.514
54 60582 0.946 22.54
Shrimp (Fig. 4) 55 40198 0.201 20.55
56 14373 0.795 23.183
57 12222 1.004 22.417
58 28344 0.257 19.509
59 56284 0.657 21.667
60 2385 0.283 21.334
61 54335 0.892 22.824
62 28841 0.673 20.971
63 6955 0.983 22.24
Shrimp (other) 64 34244 0.999 22.504
65 48298 0.429 21.663
66 37809 0.357 20.667
67 25316 0.985 23.717
68 49595 0.663 21.939
69 59467 0.487 21.568
70 9062 0.854 23.82
71 30076 0.832 22.672
72 2760 1.281 23.202
73 54335 0.892 22.824
Assembly (Fig. 5) 74 28751 0.093 23.506
75 4728 0.702 22.799
76 23187 1.183 23.565
77 45309 1.061 22.916
78 41835 0.098 19.134
79 61945 1.309 21.813
80 62605 1.011 23.143
81 44956 0.506 22.494
82 4546 0.809 22.163
83 23000 0.132 22.951
84 63112 0.499 22.273
85 43975 1.059 22.878
Equal (Fig. 6) 86 40813 0.182 19.983
87 8496 0.354 22.415
88 13836 0.661 21.054
89 11164 0.464 19.351
90 39877 0.493 22.142
– 23 –
Table 1—Continued
Type, Figure Number COMBO 17 z R mag.
91 40598 0.263 20.128
92 51021 0.743 20.96
93 35317 0.671 20.755
94 56256 0.502 20.309
95 47568 0.649 20.206
96 40766 0.46 19.997
97 24927 0.524 19.647
98 15233 0.304 18.857
99 18663 1.048 24.011
100 43242 0.657 21.177
– 24 –
Table 2: Tidal Dwarf Galaxy Candidates
Galaxy z Galaxy Diam. Dwarf V606 V606 − z850 Clump Mass
(COMBO17 #) MB,rest (mag) (kpc) MB,rest (mag) mag mag x10
1 (6423) 0.15 -20.64 13.9 -17.55 20.83 0.90 50
15 (61546) 0.552 -20.77 5.5 -16.67 26.12 0.78 1.2
16 (45115) 0.579 -20.17 4.7 -17.72 25.42 1.1 4.6
18 (41907) 0.702 -19.21 1.9 -16.03 27.45 0.55 0.24
17 (20280) 0.555 -19.56 6.2 -17.06 25.77 0.73 1.4
– 25 –
Fig. 1.— Color images of galaxies in the GEMS and GOODS fields with smooth diffuse tidal
debris. The galaxy at the top right, number 3 in Table 1, is only partially covered by the
GEMS field; the right-hand portion of the image is from ground-based observations. The
smooth debris is presumably from old stars that were spread out during the interaction. A
few small star-formation patches are evident in some cases. The clump in the upper right
corner of the galaxy 1 image could be a rare example of a gravitationally driven condensation
in a pure-stellar arm. The smooth arcs and spirals in this and other images are probably
a combination of orbital debris and flung-out tidal tails. The galaxy numbers, as listed in
Table 1, are 1 through 9, as plotted from left to right and top to bottom. (Image quality
degraded for astroph.)
– 26 –
Fig. 2.— Color images of interacting antennae galaxies with long and structured tidal arms.
Galaxy numbers, in order, are 15 through 24. Several have dwarf galaxy-like condensations
at the arm tips or broad condensations midway out in the arms. The dwarf elliptical at the
tip of the tidal arm in galaxy 20 might have existed before the interaction and been placed
there by tidal forces; the main body of this system has a double nucleus from the main
interaction. (Image quality degraded for astroph.)
– 27 –
Fig. 3.— M51-type galaxies are shown as logarithmic grayscale V-band images. In order,
the galaxy numbers are 33 through 44. The linear streak in galaxy 44 could be orbital debris
from the small companion on the right. (Image quality degraded for astroph.)
– 28 –
Fig. 4.— Shrimp galaxies, named because of their curved tails, are shown as logarithmic
V-band images. In order, they are numbers 55 through 63. (Image quality degraded for
astroph.)
– 29 –
Fig. 5.— Assembly galaxies look like they are being assembled through mergers. In order:
galaxy 74 through 85.
– 30 –
Fig. 6.— Galaxies with approximately equal-mass grazing companions, in order, are 86
through 100.
– 31 –
0 0.2 0.4 0.6 0.8 1
Restframe Johnson B–V
Equal
Assembly
Shrimp
M51 type
Antenna
Diffuse
Fig. 7.— Restframe (U-B) and (B-V) integrated colors for interacting galaxies in the GEMS
and GOODS fields, from COMBO-17. The reddest tend to be the diffuse types, which
are presumably dry mergers, and the bluest are the assembly types, which could be young
proto-galaxies. Crosses indicate standard Hubble types, measured by Fukugita et al. (1995).
– 32 –
–24 –22 –20 –18 –16 –14 –12
MB (mag)
Diffuse
Antenna
M51 type
Shrimp
Assembly
Equal
Fig. 8.— Restframe Johnson U-B integrated color versus absolute restframe MB, from
COMBO-17. The solid line separates the red sequence and blue cloud (Conselice 2006b).
Color limits for local galaxies are indicated by the horizontal short-dashed lines; local galaxies
are brighter than the vertical line. The local blue cloud galaxies are approximately delimited
on the left side of the diagram by the long-dashed lines. Thus, most of our observed galaxies
fall near the local galaxy colors and magnitudes.
– 33 –
–17 –18 –19 –20 –21 –22
Galaxy MB
Diffuse
Antenna
M51 type
Shrimp
Assembly
Tidal Dwarf
Fig. 9.— Restframe B absolute magnitudes of star-forming clumps versus integrated galaxy
restframe magnitudes. The correlation is also found for local galaxies.
– 34 –
0.01 0.1 1 10
Duration of SF (Gyr)
0.01 0.1 1 10
Duration of SF (Gyr)
–1 0 1 2 3
V606–z850
–1 0 1 2 3
V606–z850
Fig. 10.— Models at z = 1 for clump color (bottom left) and clump mass at an apparent
V606 magnitude of 27 (lower right) are shown in the bottom panels versus the duration of star
formation in 6 models with exponentially decaying star formation. Five lines are for decay
times of 107, 3×107, 108, 3×108, and 109 years, and the sixth line represents a constant rate.
Shorter decay times correspond to redder color (upper lines) and higher masses (upper lines).
In the top panels, the clump mass at V606 = 27 (top left) and the clump apparent magnitude
at 108 M⊙ masses (top right) are shown versus the clump color. The correspondence between
color and mass gives a degeneracy to plots of mass versus color at a fixed apparent magnitude
(top left) and apparent magnitude versus color at a fixed mass (top right). Thus the masses
of clumps can be derived approximately from their V606 − z850 colors and V606 magnitudes
for each redshift.
– 35 –
–1 0 1 2 3
V606–z850
0–0.125
Diffuse
Antenna
M51 type
Shrimp
Assembly
Equal
Tidal Dw.
0 1 2 3
V606–z850
0.125–0.375
0.375–0.625
0.625–0.875
z=0.875–1.125
z=1.125–1.375
Fig. 11.— The masses of the clumps can be estimated from this figure. Each curve in a
cluster of curves is a different model for color-magnitude evolution of a star-forming region,
with the age of the region changing along the curve and the exponential decay rate of the
star formation changing from curve to curve. The different clusters of curves correspond to
different total masses for the star-forming regions (mass in M⊙ is indicated to the right of each
curve). The symbols represent observations of apparent magnitude and color. Bandshifting
and absorption are considered by plotting the observations and models in redshift bins.
The mass scales shift slightly with redshift. The mass of each star-forming region can be
determined by interpolation between the curves. Typical masses are 106 M⊙ for low z and
108 M⊙ for high z. The circle near the 10
10 M⊙ curves in the z = 0.125 − 0.375 interval
corresponds to the diffuse clump in the tidal debris of galaxy 1 in Fig. 1.
– 36 –
D A M S AS E T
0–0.125
D A M S AS E T
0.125–0.375
log M
= 9.7
0.375–0.625 0.625–0.875
z=0.875–1.125 z=1.125–1.375
Fig. 12.— Clump masses (left axis) are plotted versus galaxy type in order of Figs. 1-6:
Diffuse, Antenna, M51-type, Shrimp, Assembly, and Equal, with T representing the tidal
dwarfs. The method of Fig. 11 is used. The rms deviations among the six star formation
decay times are shown as plus-symbols using the right-hand axes.
– 37 –
0.01 0.1 1 10
Duration of SF (Gy)
50 z=0.25
0.01 0.1 1 10
Duration of SF (Gy)
z=0.5
z=0.75 z=1
z=1.25 z=1.5
Fig. 13.— The apparent color of a star forming region is shown versus the duration of star
formation for an exponentially decaying star formation law. The decay times are as in Fig.
10, with short decay times the upper lines and continuous star formation the lower lines.
Using the observed clump colors, the durations of star formation are found to range between
107 and 3× 108 yrs for short decay times.
– 38 –
0 0.5 1 1.5
log(1+z)4
Diffuse
Antenna
M51 type
Shrimp
Fig. 14.— V-band surface brightness of tidal tails for galaxies in Figures 1-4 plotted as a
function of (1 + z)4 for redshift z. Some systems have more than one tail. Cosmological
dimming causes a decrease with redshift equal to 2.5 magnitudes for each factor of 10 in
(1 + z)
; this decrease is consistent with the dimming seen here. The observable 2σ limit
for these fields is ∼ 25 mag arcsec−2. Some antenna galaxies have patchy tails with fainter
average surface brightnesses.
– 39 –
0 0.5 1
Redshift, z
Diffuse
Antenna
M51 type
Shrimp
Fig. 15.— Fraction of V-band luminosity in antennae tidal tails relative to their integrated
galaxy luminosity, as a function of redshift.
– 40 –
0 5 10 15 20 25 30 35
Disk Diameter (kpc)
Diffuse
Antenna
M51 type
Shrimp
Fig. 16.— Tail length versus disk diameter from Figs. 1-4, based on the V-band images.
Conversions to linear size assumed a standard ΛCDM cosmology applied to the photometric
redshifts.
– 41 –
0 0.5 1 1.5
Redshift, z
Diffuse
Antenna
M51 type
Shrimp
Fig. 17.— Tail length/disk diameter as a function of redshift for shrimps and antennae,
measured from the V-band images. There is no obvious trend.
– 42 –
0 50 100 150 200
Tail Length (kpc)
GEMS GOODS Antennae
Local Antennae
Superantennae
Arp 299
NGC 3628Arp 241
VV109
NGC 3256
Antennae
Arp 243
Arp 242
Arp 226
Arp 157 Arp 75
Arp 35 NGC 5548
Arp 33
Arp 102
Fig. 18.— Tail length/disk diameter versus the tail length for antenna galaxies in our sample
as well as for local antennae, whose names are indicated. The GEMS and GOODS systems
are significantly smaller than the local antenna galaxies, even if the two extreme local cases,
the Superantennae and Arp 299, are excluded.
	Introduction
	The Sample of Interactions and Mergers
	Photometric Results
	Global galaxy properties
	Clump properties
	Tail Properties
	Tidal dwarf galaxy candidates
	Dark Matter Halo Constraints
	Conclusions
