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Molecular gas in strongly interacting galaxies: II. Global
properties of the sample
Ming Zhu 1 , Howard A. Bushouse 2 , E. R. Seaquist 1 , Emmanuel Davoust 3
and
David T. Frayer 4
Received ; accepted
1 Dept. of Astronomy, U. of Toronto, 60 St. George St. Toronto, ON M5S\Gamma3H8, Canada
2 Space Telescope Science Institute
3 UMR 5572, Observatoire Midi­Pyr'en'ees, 14 Avenue Edouard Belin, 31400 Toulouse, France
4 Astronomy Department, California Institute of Technology, 105­24, Pasadena, CA 91125 (USA)

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ABSTRACT
We have collected CO data on a sample of 95 strongly interacting galaxies
(SIG) and on comparison samples of 59 weakly interacting and 69 isolated spiral
galaxies. The statistical analysis of the samples shows that the SIGs, especially
the colliding and merging systems, have higher CO luminosity (per unit optical
area or luminosity) than isolated spiral galaxies. If this excess is interpreted as
excess H 2 , then we find no significant difference in the molecular to atomic gas
mass ratio between the samples; this indicates that the excess molecular gas is
not due to conversion of HI to H 2 or to the removal of HI gas from the galaxies
by interaction. Another possible interpretation of the excess CO luminosity and
of the normal gas ratio is that the CO­to­H 2 conversion factor is lower in SIGs
than in isolated starbursting galaxies. In agreement with previous studies, we
find that the star formation rate (estimated by the far infrared luminosity) is
higher in the SIGs than in the isolated galaxies. The star formation efficiency
(measured by the ratio of far infrared luminosity to inferred molecular gas mass)
is higher than average in the mergers and colliding systems only. Our results are
in agreement with a scenario in which gravitational interaction produces intense
star formation which becomes highly efficient in the late stages of evolution.
Subject headings: galaxies: interacting---galaxies: ISM: molecules :
ISM: galaxies --- starburst

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1. Introduction
The influence of gravitational perturbations on spiral galaxies has drawn much attention
in recent years. Galaxy collisions and mergers can have a dramatic impact upon the
morphology and subsequent dynamical evolution of galaxies. Interactions are thought
to be the major cause of the extensive starburst phenomenon seen at high redshift (e.g.
Ivison et al. 2000; Scoville 2000). In the local universe, starbursts induced by galaxy
interactions are believed to be the major source of the tremendous energy output from the
so called ultra­luminous IRAS galaxies (ULIRGs). Indeed, virtually all the ULIRGs with
L IR ? 10 12 L fi are found to be mergers (Sanders et al. 1991). The star formation activity
indicators such as Hff and far­infrared and radio continuum emission all point to a higher
level of star formation activity in interacting galaxies (IG's) compared to isolated spiral
galaxies (ISG's) (e.g, Kennicutt et al. 1987; Xu & Sulentic 1991; Hummel et al. 1990;
Bushouse 1987, 1988).
Molecular clouds are the birth places for star formation. Previous CO surveys of
external galaxies have established a close correlation between the CO luminosity and the
total far infrared (FIR) luminosity (c.f. Young & Scoville 1991). If this correlation holds for
IG's, their FIR enhancement would suggest excess CO emission in IGs compared to ISG's.
This prediction seemed to be confirmed by some CO studies of external galaxies which
reported that the MH2 =LB ratio is enhanced in IGs compared to isolated galaxies (Solomon
1988; Young 1996; Brains, 1993; Combes et al. 1994). However, most previous studies
were based on a relatively small sample and the interacting systems were often selected
according to their IRAS flux, which biases these studies toward gas­rich galaxies. The
reason for this selection is the relatively low sensitivity of millimeter radio telescopes which
made it difficult to detect the CO emission in a large number of gas­poor galaxies within
a reasonable observing time. Nevertheless, the CO data for individual interacting systems

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has been accumulating in the literature over the years. Ironically, interacting systems in the
southern hemisphere are studied more systematically and more comprehensively than the
IGs in the northern sky (e.g., Combes et al. 1994; Horellou & Booth 1997).
We have conducted observations with the NRAO 12m telescope and the IRAM 30m
telescope (Zhu et al. 1999, Paper I), and by combining these data with those obtained
by other investigators, we have compiled a large CO database comprising a complete
optically­selected sample of strongly interacting galaxies (SIGs) in the northern hemisphere.
Here we focus our study on SIGs for mainly two reasons: (1) SIGs can be easily identified
with minimal error because they exhibit obvious disturbances in morphology and most of
them have close companions (except for mergers); (2) the number of SIGs is relatively small
in the local universe so we can observe and sample them completely. Our sample contains
154 IGs (including 95 SIGs) and covers different galaxy progenitor types, interaction phases,
and encounter geometries. This enables us to arrive at a statistically meaningful conclusion
on the influence of galaxy interaction on the molecular gas properties and induced star
formation activity.
This is the second paper in a series devoted to a statistical analysis of the molecular gas
properties in IGs. Our goal is to make clear whether the CO emission is enhanced in IGs
and whether the enhanced star formation activity in IGs is due to a higher abundance of
molecular gas or to a higher star formation efficiency. In Paper I we presented most of the
CO data obtained by us. In this paper, we combine the available CO, HI and IRAS data
in the literature for our sample galaxies and conduct an analysis of the CO, HI and FIR
properties, including a comparison between SIGs and ISGs.

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2. The sample
2.1. A complete SIG sample in the northern sky
The complete sample of SIGs was compiled from the UGC (Nilson 1973), ARP catalog
(Arp 1966) and the Catalogue of Isolated Pairs of Galaxies in the Northern Sky (CPG)
(Karachentsev 1972). Bushouse (1986) has identified all the SIGs from galaxies that
have been labeled as disturbed, distorted, or with bridges or warped disks in the UGC.
By searching the Third Reference Catalogue of Bright Galaxies (RC3) (de Vaucouleurs
1991), we obtained the optical (blue) magnitude for each galaxy in Bushouse's complete
sample. Then all systems with at least one member brighter than B T = 14:5 were chosen,
yielding 164 systems. The ARP catalog contains 338 systems, but it also includes some
elliptical galaxies, isolated peculiar galaxies, and some galaxies in the southern sky. The
CPG contains 1206 galaxies in 603 isolated pairs (including elliptical galaxy pairs) and is
complete to m pg = 15:0 (Stocke 1977). Its selection criteria exclude compact group members
and also under­sample merger types. For our purpose, only the spiral pairs were selected
from the CPG. These three sources overlap with one another and together they cover most
of the field IGs in the northern sky (DEC ? \Gamma2 ffi ). In order to select our SIG sample, we
further inspected all the optical images from the Digitized Sky Survey and classified these
systems according to the following criteria:
ffl IG1: Weakly interacting galaxies (WIG's), including all the interacting galaxies with
component separation ? 1.5 D 25 and with little or no morphology disturbance. Here
D 25 is the diameter of the larger component in the pair, using the brightness contour
B T = 25 mag arcsec \Gamma2 .
ffl IG2: Galaxies with component separation ! 1.5 D 25 showing obvious morphology
disturbances. Those with strong morphology disturbances but with a component

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separation up to 2 D 25 were also included.
ffl IG3: Galaxies in projected contact with their companions and with severe morphological
disturbances, i.e. colliding galaxies. We also include in this category the so called ``ring
galaxies'' which are believed to be the remnants of a head­on collision.
ffl IG4: Merging systems, with a single amorphous body but double nucleus, and a pair of
remnant tidal tails.
ffl IG5: Late­stage mergers, with one or two tidal tails but only a single nucleus.
In this paper SIGs are defined as those in classes IG2--IG4. Since our main interest is the
SIGs, detailed classifications for the SIGs are given and all the paired galaxies not belonging
to SIGs were put into the WIG group (IG1) for comparison. The very late­stage mergers
were put into another group (IG5) because galaxies of this type should have consumed most
of their gas and thus there should be no significant CO enhancement. The very late­stage
mergers can be easily confused with isolated elliptical galaxies, so only those that have
some tidal features such as tails were retained in the sample.
Finally, our sample galaxies were selected according to the following criteria:
ffl (1) must be a SIG (IG2--IG4)
ffl (2) at least one member in the interacting system is brighter than B T = 14:5
ffl (3) declination DEC? \Gamma2 ffi .
Our final sample of SIGs contains 126 galaxies in 92 systems. This is by far the most
complete sample of SIGs in the northern sky with DEC ? \Gamma2 ffi . Table 1 lists the general
properties of our sample galaxies and their IG type. It is ordered by interaction type (IG),
and, within the type, by right ascension (RA). The names are in cols. 1, 2 and 3, the
interaction type (IG) in col 4, the coordinates (RA and DEC) in cols. 5 and 6, the optical

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radial velocity (cz, in km s \Gamma1 ) in col. 7, the corrected blue apparent magnitude (B 0
T ) in col.
8, the apparent diameter (D 25 in arcmin.) in col. 9, the minor axis (r b , in arcmin.) in col.
10 and the morphological type in col. 11. The data are from RC3.
To evaluate the completeness of our sample, we searched the literature for any obvious
SIGs with B ! 14:5 that had not been included in our sample, and none were found.
For example, Dahari (1985) compiled 209 peculiar galaxies with B ! 14:4 from the Atlas
and Catalog of Interacting Galaxies (Vorontsov­Velyaminov 1959, 1977). He also divided
his sample galaxies into six interaction classes (IAC). All IAC3­IAC6 galaxies in Dahari's
sample were inspected, and it was found that all the IAC4 to IAC6 galaxies had already
been included in our sample. All the IAC1, IAC2 and some IAC3 galaxies were not
included because their morphology distortion is not strong. Although the decision to call an
individual system ``strongly interacting'' is a subjective one, the IGs with close components
and severe morphology disturbance are easy to identify and their chance of being missed by
the UGC, ARP or CPG catalogue is small. Hence our SIG sample should be essentially
complete for paired galaxies. Most of the IGs in compact groups were not included because
they are not included in either the UGC or the CPG.
The only source of confusion in classifying the galaxies may be in the merger types.
For some mergers it is difficult to tell whether they contain double nuclei. To distinguish
between IG4 and IG5, use was made of the work by Keel & Wu (1995) who listed the
number of nuclei and tails for a sample of disk­disk merger remnants in the local universe.
We also used near IR images available in the literature (e.g. Stanford & Bushouse 1991;
Bushouse & Stanford 1992). It is possible that some IG4 mergers have been mis­classified
as IG5 if their double nucleus has not been detected. However, this should not seriously
affect our statistical study since the number of such ambiguous cases is very limited.
In summary, our SIG sample contains most of the strongly interacting systems in

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the northern sky. This sample is selected strictly according to optical brightness and
morphology, and without any a priori knowledge of the gas content or FIR emission level.
Thus it is unbiased and ideal for studying the molecular gas properties of IGs.
2.2. The control sample
The control sample of isolated spiral galaxies was obtained from the FCRAO CO survey
made by Young et al. (1996), which is the most homogeneous sample for nearby spiral
galaxies. Since the IGs are generally more luminous than the ISGs in the FCRAO sample
(Young et al. 1995), only the more luminous (LB ? 2 \Theta 10 9 L fi ) spiral galaxies were selected
for the control sample. We also excluded the galaxies belonging to the Virgo cluster. The
final control sample comprises 69 galaxies. Figure 1a shows the distribution of MB for the
SIG and control samples. The average LB of the SIG sample is 2:7 \Theta 10 10 L fi , which is
about 1.8 times higher than that of the control sample. This is primarily because the SIG
sample as a whole is biased towards more distant, and hence more luminous, galaxies, since
there are relatively few SIGs nearby. Moreover, some merging systems include emission
from more than one galaxy.
Figure 1b compares the morphological types of the SIG and control samples. There are
many more Sbc and Sc galaxies in the control sample than in the SIG sample, where the
galaxies are more uniformly distributed among all types. It is not surprising that there
are many Irr and Peculiar galaxies in the SIG sample. Furthermore, about 20% of the
SIGs, especially the merging systems, have a morphology so distorted that they cannot be
classified as any specific Hubble type. These spiral galaxies, which are designated S?, are
not shown in the figure. Since the primary concern is with the distorted galaxies, we have
not attempted to compare the morphology distributions of these two samples, nor did we
compare any parameters that are sensitive to galaxy morphology.

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3. CO fluxes and molecular gas masses
Since our earlier results on 80 interacting galaxies were reported in Paper I (Zhu et
al. 1999), another 30 galaxies were observed with NRAO 12m in March, May, and June
1999. The observing procedures used are identical to that outlined in Paper I. Another 66
galaxies with good CO measurements were obtained from the literature. In addition, all the
available CO data for weak IGs (IG1) and late stage mergers (IG5) were collected providing
a consistent CO data set for all types of IGs suitable for the statistical study.
In order to make comparisons between the IG sample and the ISG sample of Young et
al. (1995), we derived the global CO flux and the inferred total mass of the molecular
gas MH 2
using the same method as that of Young et al. (1995) (see also Paper I). For all
the galaxies larger than twice the telescope beam, several measurements along the major
axis were made and the global flux was derived with a Gaussian or exponential brightness
distribution model which best fit the data. For galaxies only slightly larger than the
telescope beam, measurements were made at a single position only and an exponential
model was used to derive the integrated flux from the observed flux. Similar procedures
were applied to the galaxies observed by other investigators to derive the global CO fluxes.
There are a total of 92 SIGs for which the total CO flux is derived consistently.
In Paper I it was shown that the uncertainty in the global fluxes derived in this way are
usually less than 40%, and the data taken by the IRAM 30 telescope and the NRAO 12m
telescope are generally consistent with each other after correcting for the flux outside the
telescope beam. Uncertainties for the data in the literature are more difficult to estimate.
Some galaxies, especially the typical interacting galaxies such as ARP220, IC883 have been
observed by different groups. Higher priority was given to data with high S/N ratio and
with larger beam sizes compared to the source.
The total MH2 may be derived from the global CO(1­0) flux by assuming a standard

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Galactic CO­to­H 2 conversion factor, (i.e. the X factor = N(H 2 )=I CO ). However, there is
strong observational evidence that the X factor is significantly different from one galaxy
to another (e.g Crawford et al. 1985; Stacey et al. 1991; Solomon et al. 1997; Downes &
Solomon 1998; Wilson et al. 1995), and a radial variation of more than a factor of 10 has
been reported in our Galaxy (Sodroski et al. 1995). Recent studies of the molecular gas in
ULIRGs have shown that the MH2 =L CO ratio may be 3­5 times lower in the centers of these
galaxies than in the Galactic molecular clouds (Solomon et al. 1997; Downes & Solomon
1998). Hence the application of the standard X factor leads to significant overestimates of
the MH2 in these galaxies. There is also a physical basis for such a variation in the X factor,
since the CO flux per unit mass of molecular gas is expected to be a function of gas kinetic
temperature, density, and chemical composition (cf. Young 1991; Wilson et al. 1995). In
particular, it has been shown to increase strongly with decreasing metallicity (Israel 2000),
which is consistent with the tenfold increase outward in our Galaxy (Sodroski et al. 1995).
Accordingly, we introduce a parameter H \Lambda
2 to distinguish the H 2 content derived from
the standard conversion factor from the actual H 2 content, bearing in mind that it is in
reality a reflection of the CO flux rather than a precise measure of the mass of molecular
gas. The total H \Lambda
2 mass may be derived from the CO fluxes using the conversion factor
X = N(H \Lambda
2 )=Ico = 2:8 \Theta 10 20 cm \Gamma2 [Kkms \Gamma1 ] \Gamma1 (Bloemen et al. 1986). Kenney & Young
(1989) have shown that this value of the conversion factor leads to H \Lambda
2 masses in solar units
given by MH \Lambda
2
= 1:1 \Theta 10 4 D 2 SCO , where D is the distance in Mpc and SCO is the CO flux
in Jy km s \Gamma1 .
Tables 2 and 3 lists the interacting galaxies with CO measurements made by us (Table
2) or available in the literature (Table 3). For consistency, the global masses MH \Lambda
2
based on
CO measurements in the literature were re­determined using Young's method as described
above. We list not only the SIGs (IG2 -- IG4), but also the available data for weak IGs
(IG1) and late stage mergers (IG5), in order to provide a comparison sample of weakly

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interacting systems. A total of 154 IGs are listed in Tables 2 and 3, including 95 SIGs, 44
IG1 and 15 IG5 galaxies. Out of the 126 SIGs in the complete SIG sample (Table 1), 75%
of them have CO data available (including non­detections).
4. HI fluxes and IRAS flux densities
HI fluxes for the galaxies in our sample have been taken from the RC3 and the catalogue
of Huchtmeier & Richter (1989), or from Bushouse (1987) and Davis & Seaquist (1981).
Since it has been shown that a significant amount of atomic gas in IGs may have been
dragged out of the optical disk along the tidal tails (Hibbard et al. 1999, 2000), we
preferentially selected measurements of the global HI flux made using telescopes with a
larger beam size, such as the NRAO 92m telescope. We also gave most weight to the more
recent and more sensitive measurements.
The mass of HI is given by:
MHI = 2:36 \Theta 10 5 D 2 S(HI) (1)
where D is the distance in Mpc and S(HI) is the HI flux in Jy km s \Gamma1 . The global HI
fluxes and the derived atomic gas masses are listed in Table 4. The uncertainties in these
values are approximately 30­40%.
The infrared luminosities (in solar units) integrated from 1 to 500 ¯m are based on IRAS
flux densities at 60 and 100 ¯m, denoted as S 60 and S 100 , and were computed using the
formula given by Lonsdale et al. (1985):
L(IR) = 3:94 \Theta 10 5 D 2 [2:58S 60 + S 100 ]; (2)

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where S 60 and S 100 are In Jy.
Xu and Sulentic (1991) and Bushouse (1988) have re­processed the IRAS data and have
derived the IRAS flux densities for their sample galaxies. Their data were used when
available. Other data have been obtained from the NED database. The IRAS 60 and 100
¯m flux densities and the derived L IR for the CO sample SIGs are given in Table 4.
Since the resolution is not high enough to resolve the interacting galaxies, except for a few
widely separated pairs of type IG1, all the IRAS measurements are totals for each system,
as are most of the HI fluxes. Therefore there is only one entry in Table 4 for most paired
systems.
5. Comparison with isolated spiral galaxies
Table 5 lists the statistical properties of all the interacting and isolated galaxies in our
samples. The mean properties were determined using the KMESTM program in the package
ASURV Rev 1.2 (see LaValley et al. 1992 and references therein). This program computes
the Kaplan­Meier estimator (Kaplan & Meier 1958) of a randomly censored distribution
allowing for the non­detections in the sample. The second row for each ratio in Table 5
lists the median properties. These two statistical properties have different advantages. The
median is less sensitive than the means to a few extreme values in the sample. However,
unlike the mean, it is impossible to estimate the uncertainties in the median. Knowing
the uncertainties allows us to test the statistical significance of any difference between the
samples. Therefore we based our discussion mainly on the mean values. However both the
mean and median properties point to the same trend.
The galaxies in the SIG sample on average have a higher LB and larger physical size. In
order to compare the gas content of galaxies between the two samples, we need to eliminate

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the effect of the size. This can be achieved by normalizing the molecular gas mass MH \Lambda
2
by
the optical lumminosity or by the optical area of the galaxy. We adopt the same normalized
quantities as Casoli (1998), namely the MH \Lambda
2
=LB and MH \Lambda
2
=D 2
25 ratios.
5.1. The MH \Lambda
2
=LB ratio
Figure 2 shows the distributions of the MH \Lambda
2
=LB ratio for the two samples. Most of the
normal spirals have a MH \Lambda
2
=LB ratio smaller than 0.3. M82 and NGC1055 are the only two
galaxies in Young's sample with a MH \Lambda
2
=LB ratio higher than 0.31, but they are not really
isolated galaxies. M82 is gravitationally interacting with M81, and N1055 is about 130
kpc from NGC 1068. Therefore essentially no strictly isolated galaxies have been found to
possess a MH \Lambda
2
=LB ratio higher than 0.31. On the other hand, high ratios of MH \Lambda
2
=LB are
frequently seen in IGs and mergers. Merger remnants such as ARP220 have MH \Lambda
2
=LB ratios
as high as 1.1. The average MH \Lambda
2
=LB ratio of the SIG sample is 0.22, which is 1.7 times
higher than that of the ISG sample.
The high MH \Lambda
2
=LB ratios suggest that SIGs, especially the colliding and merging systems,
have more molecular gas than isolated spirals. However, Perea et al. (1997) have argued
that the relation between LB and MH \Lambda
2
is not linear, indicating that normalizing the mass by
LB does not completely remove the size effect. To evaluate the latter, we plot MH \Lambda
2
against
LB in Fig. 3. The merging systems (IG3 and IG4) are shown as solid circles and IG2
systems are shown as open circles while stars represent the isolated spirals in the control
sample. The values of logL B range from 9:6 to 11. A linear regression applied to the ISG
sample, using the ASURV V1.2 BIVAR method which takes into account the upper limits,
yields

-- 14 --
logMH \Lambda
2
= (1:30 \Sigma 0:10)logL B \Gamma (4:04 \Sigma 0:98) (3)
This fit is indicated by a dashed line in Fig. 3. This line fits the ISG and IG2 systems
reasonally well, but fails to fit the merging systems. Examining the deviation of each galaxy
from the dashed line and applying a Peto­Prentice generalized Wilcoxon test (using the
ASURV package) for each subsample, we found that the merger type (IG3 + IG4) has a
residual significantly different from zero with a confidence level of 96%. Not surprisingly,
the residual for the IG2 subsample and the ISG control sample does not significantly deviate
from zero. Hence the high MH \Lambda
2
=LB ratios in merging galaxies are significant and cannot be
accounted for by the size effect. The fact that no isolated galaxies are found in the high
MH \Lambda
2
=LB regime strongly suggests that the high MH \Lambda
2
=LB ratio is the result of strong galaxy
interactions. However, not every interacting or merging system is associated with a high
MH \Lambda
2
=LB ratio.
5.2. The MH \Lambda
2
=D 2
25
ratio
Another way to eliminate the size effect is normalization by the optical surface area of
the galaxies. The mean and median values of MH \Lambda
2
=D 2
25 for the SIG sample as well as
for the individual IG subsamples are given in Table 5. The result is similar to that from
normalization by LB . The average MH \Lambda
2
=D 2
25 ratio is about twice as high in SIGs. The
most significant differences come from the merger types. A two­sample Wilcoxon test shows
that the IG3 and IG4 subsamples have a mean MH \Lambda
2
=D 2
25 ratio significantly higher than
that of the ISGs with a confidence level of 99%, while the confidence level drops to 94% for
the IG2 subsample. The histogram in Fig. 4 shows that the MH \Lambda
2
=D 2
25 distribution for the
SIG sample is more populated at the higher end, though the dispersion is also larger in the
SIG sample.

-- 15 --
Figure 5 shows a plot of MH \Lambda
2
versus D 2
25 for the subsample of IG2 (open circles), merger
type (solid circles) and ISGs (stars). A linear regression fit to the isolated galaxy sample
(dashed line) yields
logMH \Lambda
2
= (1:81 \Sigma 0:24)logD 25 + (6:85 \Sigma 0:32) (4)
or
logMH \Lambda
2
= (0:91 \Sigma 0:12)logD 2
25 + (6:85 \Sigma 0:32) (5)
This relation is similar to that derived by Casoli (1998) from a sample of more than 500
isolated galaxies. From this we can see that the relation between MH \Lambda
2
and D 25 is consistent
with a linear one and that the ratio MH \Lambda
2
=D 2
25 should completely remove the size effect.
Therefore, the high MH \Lambda
2
=D 2
25 ratio for IG3 and IG4 is obviously not due to their larger
size.
On Fig. 5 we can also see that the dashed line which best fits the ISG sample fails to
fit the SIG sample. Statistical tests show that the residuals are significantly different from
zero value at the confidence level of 99% and 95% for the merger type and IG2 subsamples,
respectively.
Due to the morphology distortion in virtually all SIGs, the size (D 25 ) may have a large
uncertainty. However, one is more likely to overestimate D 25 of a disturbed disk, and thus
underestimate the MH \Lambda
2
=D 2
25 ratio. Even if the errors on D 25 were random, this should
not result in a higher average MH \Lambda
2
=D 2
25 ratio for SIGs. Therefore the high MH \Lambda
2
=D 2
25 ratio
in some SIGs must be genuine and indicates an CO enhancement as a result of galaxy
interaction.

-- 16 --
Casoli (1998) showed that the MH \Lambda
2
=D 2
25 ratio depends on the Hubble type, with the Sa,
Sb and Sbc types containing ¸ 6 times more molecular gas than late­type spirals and Irr
galaxies. If a galaxy sample includes a majority of Sa­Sbc galaxies, it would have a higher
than normal average MH \Lambda
2
=D 2
25 ratio. However this is not the case for our SIG sample. In
Figure 1b we have shown that the SIG sample is not overpopulated with Sa­Sbc galaxies
compared to the control sample.
In summary, both the MH \Lambda
2
=LB and MH \Lambda
2
=D 2
25 ratios suggest either an enhancement in
molecular gas content or in the luminous efficiency of CO for some SIGs, especially for
IG3 and IG4. The choice between these two options will depend on careful modeling of a
number of individual systems.
6. Molecular versus atomic gas
Assuming MH \Lambda
2
represents the actual molecular mass MH 2
, we may compare the total
masses of molecular and atomic gases in our samples, in order to investigate the possibility
that some systems are enriched in molecular gas, either by phase conversion of HI to H 2 , or
by removal of HI from the system by tidal forces.
Figure 6 shows the distributions of the MH \Lambda
2
=MHI ratio for the SIG and ISG samples.
Both samples peak at MH \Lambda
2
=MHI ¸ 1 and there is no significant difference between them.
The average MH \Lambda
2
=MHI ratio for SIGs is only 1.3 times higher than that of the ISG sample
(Table 5). The merger type subsample has a higher MH \Lambda
2
=MHI mean value, but this is not
statistically significant, as the confidence level is only 83%.
Table 5 shows that the MHI =LB ratios of the SIGs and ISG samples are similar. This
means that the high MH \Lambda
2
=MHI ratios in some systems corresponds to an enhancement of
MH \Lambda
2
rather than to a depletion of MHI . This is a different conclusion from that derived for

-- 17 --
the IRAS­selected IG sample (e.g. Martin et al. 1991), which shows a depletion of HI gas for
the ULIRGs. This result also indicates that the single­dish HI data do not systematically
underestimate the HI gas even for the late­stage mergers (IG3 and IG4) which could have a
significant amount of HI gas ejected from the galaxy disks by tidal forces (e.g. Hibbard et
al. 1999).
Figure 7 shows plots of MH \Lambda
2
=LB versus MHI =LB for the SIGs. No correlation is found
between these two quantities (correlation coefficient = 0.23). This provides no support for
the idea that the CO enhancement corresponds to a phase conversion from HI to H 2 as
there is no anti­correlation between MH \Lambda
2
=LB and MHI =LB .
Figure 8a shows that MH \Lambda
2
=MHI is correlated with MH \Lambda
2
=LB for the SIGs with a correlation
coefficient of 0.76, confirming that the high MH \Lambda
2
=MHI ratios in some systems are due to
an excess of molecular gas rather than to a deficiency of atomic gas. In contrast, there is
almost no correlation between MH \Lambda
2
=MHI and MH \Lambda
2
=LB for the ISG sample (correlation
coefficient = 0.33) (see Fig. 8b). Therefore the extremes in the MH \Lambda
2
=MHI ratio in normal
galaxies may reflect variations in the richness of atomic, rather than molecular gas. The fact
that the MH \Lambda
2
=MHI -- MH \Lambda
2
=LB correlation is only seen in SIGs indicates that the observed
enhancement in CO emission is a an effect of galaxy interactions on the properties of the
ISM.
Casoli et al. (1998), using a hybrid sample of more than 500 spiral galaxies, reported an
average MH \Lambda
2
=MHI ratio of 0.2­0.34, which is not only one magnitude lower than that of
our SIG sample, but also much lower than that reported by Young et al. (1996). They
argued that Young's sample contains more FIR­bright and optically bright galaxies, while
their own sample contains more dwarf galaxies. They also showed that MH \Lambda
2
=MHI is weakly
correlated with LB , suggesting that there might be a size effect for the high MH \Lambda
2
=MHI ratio
in more luminous galaxies. However, this correlation is very weak in Casoli's sample. We

-- 18 --
also plot log(MH \Lambda
2
=MHI ) versus log(L B ) for the ISG control sample in Fig. 9, and see
virtually no correlation between MH \Lambda
2
=MHI and LB . Hence there is no evidence that the
high MH \Lambda
2
=MHI ratio in SIGs is the result of a size effect.
7. Dependence of statistical properties on interaction type and separation
Table 5 gives the statistical properties of each IG subsample and Fig. 12 presents these
properties graphically. Both show that the average MH \Lambda
2
=LB , MH \Lambda
2
=D 2
25 , MH \Lambda
2
=MHI ,
L IR =LB , L IR =MGAS \Lambda and S 60 =S 100 ratios all systematically increase with interaction
strength. Here MGAS \Lambda is defined as MH \Lambda
2
+MHI .
7.1. The star formation rate
The FIR luminosity L IR of a galaxy is a good measure of its global star formation rate
(SFR) (see Kennicutt 1998 and references therein). The FIR emission is mainly the energy
reradiated by dust that has been heated by young massive stars in star forming regions.
There is a well established correlation between FIR and CO emission for both spiral and
interacting galaxies (Young et al. 1996 and references therein). In agreement with previous
studies, we find a good correlation between L IR =LB and MH \Lambda
2
=LB for both the interacting
(Fig. 10) and isolated galaxy samples. The average L IR =LB ratio is about twice as high in
SIGs than in ISGs (Table 5), suggesting a higher star formation rate.
Although the IG3 and IG4 subsamples include a relatively small number of systems
and there may be errors in our IG type classification, both subsamples show significant
enhancement in the FIR and CO emission. If we group IG3 and IG4 into a merging­system
category, this class forms a larger subsample that is not affected by errors in classification

-- 19 --
between the two subsamples. Therefore, it is safe to say that merging systems have an
enhancement in CO and FIR luminosity. We have also performed the same statistical
analysis for an enlarged IG3 and IG4 subsample by including all the known merging systems
in the southern sky such as NGC1614, ARP236, NGC4038/39, NGC7592, NGC7252, and
the conclusion remains the same.
The variations in the L IR =LB and MH \Lambda
2
=LB ratios within each IG subsample are quite
large however. For example, in several IG4 systems the MH \Lambda
2
=LB ratio is as low as 0.02 (for
NGC3239) and as high as 1.09 (for ARP220). Thus not every merging process produces a
CO and FIR enhancement.
NGC3239 is an example of a merger with low MH \Lambda
2
=LB . It is a small galaxy with LB less
than 10 10 L fi ; it is rich in HI, but has little H \Lambda
2 . Thus the merging of dwarf galaxies may
not produce a CO enhancement. Another example is NGC4194, which is classified as IG5
because it has only one nucleus. However, it has FIR and CO properties similar to most
IG4 galaxies. It may be at a stage later than ARP220 but still earlier than most of the
other IG5 galaxies. Starburst activity is still quite dramatic in these systems because the
gas has not yet been consumed, even though the nuclei have merged.
Figure 11 shows the correlation between L IR =LB and MH \Lambda
2
=MHI for the merger type
subsample (solid circles), IG2 (open circles) and ISG sample (stars). This correlation is
strong for the merger type subsample, but very weak for IG2 and ISGs. This suggests that
molecular gas­rich mergers are more likely to have a high SFR, but the tidal force may not
be strong enough in interacting pairs (IG2) to trigger a high level of star formation even if
there is plenty of molecular gas available in the system.

-- 20 --
7.2. The star formation efficiency
If MH \Lambda
2
is indicative of the true mass of molecular gas, then the L IR =MH \Lambda
2
ratio can be
used as an indicator of the star formation efficiency (SFE) which is the number of newly
formed stars per unit mass of molecular gas. This ratio is not significantly enhanced for
strongly interacting pairs (IG2 and IG3) compared to the weak tidal types (IG1). Only
mergers (IG4 and IG5) have a significantly higher SFE. In the case of IG5 this is mostly
due to the depletion of molecular gas at a late stage of merging rather than to a high L IR
emission. Only mergers (IG4 and IG5) have a higher SFE than ISGs by a factor of ¸ 2.
Therefore the high SFR in most SIGs is mainly due to more molecular gas available to fuel
the gas.
7.3. The dependence on projected separation
The projected separation is also a statistical indicator of the interaction strength. Figure
13 shows the relationship between Sep=D 25 and several derived quantities of the interacting
systems, where Sep=D 25 is the pair separation normalized by the primary's diameter D 25 .
We see a weak correlation between Sep=D 25 and L IR =MGAS \Lambda , L IR =LB , L IR =MH \Lambda
2
, S 60 =S 100
and MH \Lambda
2
=LB , but there are almost no correlations between Sep=D 25 and MHI =LB ,
MH \Lambda
2
=MHI . In all cases, the scatter is quite large. This is not surprising, because the
projected pair separation is not a measure of the true linear separation. Sep=D 25 provides
only a lower limit on the true separation due to projection effects, and it also depends on the
evolutionary stage of the encountering pair. Therefore, the IG classification according to the
disturbed morphology may be a better measure of the interaction strength. Nevertheless,
the correlation between Sep=D 25 and L IR =LB or L IR =MGAS \Lambda confirms the results found by
other investigators that closer interactions increase the chances of producing enhanced star
formation and hence more FIR emission. The MHI =LB ratio is not anti­correlated with

-- 21 --
Sep=D 25 , and provides no evidence for conversion of atomic gas to molecular gas along the
merging sequence.
8. Discussion: More molecular gas in SIGs?
Here we return to the issue of whether the quantity MH \Lambda
2
truly represents the mass of
molecular gas. As shown in the previous sections, SIGs on average have higher MH \Lambda
2
=LB and
MH \Lambda
2
=D 2
25 ratios than ISGs, which suggests that SIGs have more molecular gas than ISGs.
However, as discussed in section 6, there is no evidence for HI depletion in SIGs. HI appears
deficient in the late­stage mergers (IG5) only. If the extra molecular gas is not produced by
conversion of HI to H 2 , where does it come from? There is a priori no reason why galaxies
about to interact initially have more than average amounts of molecular gas. It is thus
necessary to examine whether the global CO flux or the X factor of the SIGs have been
overestimated before stating that SIGs have more molecular gas than average.
It is unlikely that the excess CO emission seen in SIGs is caused by a significant
overestimate of the derived global CO flux. As discussed in Paper I, if the molecular gas
is more centrally distributed in SIGs, the global CO flux might be overestimated to some
degree when a correction factor is applied to take into account emission outside the beam.
However this correction factor is small for most of the sample. Thus any overestimate
should be less than 40%, and therefore cannot account for all of the CO enhancement in
most SIGs. Also there are a number of SIGs, such as ARP220, ARP193, and UGC8874
that have been resolved by interferometers, and the derived MH \Lambda
2
=LB and MH \Lambda
2
=D 2
25 ratios
are still significantly higher than those of normal ISGs. Thus the CO enhancement is real,
and the question is then indeed whether the X factor is larger in SIGs, or whether there is
really more MH \Lambda
2
in SIGs than in ISGs.

-- 22 --
Figure 10 shows that MH \Lambda
2
=LB is strongly correlated with the SFR indicator L IR =LB , and
thus many of the galaxies with high MH \Lambda
2
=LB are FIR­luminous systems. If their X factor
is on average reduced by a factor of two, the resulting average MH \Lambda
2
=LB and MH \Lambda
2
=MHI for
SIGs would become comparable to that of ISGs. Thus the higher MH \Lambda
2
=LB ratio in the SIG
sample may be the result of the dependence of the X factor on physical and/or chemical
conditions in the clouds. From Table 5 we can see S 60 =S 100 is slightly higher in SIGs than in
ISGs, suggesting a higher dust temperature in SIGs. However, this increment is marginal.
In particular, there is no significant correlation between MH \Lambda
2
=LB and S 60 =S 100 indicating a
weak dependence at best of the X factor on temperature.
We have further investigated this issue in a detailed case study of the nearest prototype
colliding pair NGC4038/9, in which we can obtain relatively high resolution multi­transition
data from both 12 CO and 13 CO to get a better estimate of the physical parameters and the
molecular gas mass (Zhu et al. in preparation). We found that the 12 CO is ``over luminous''
with respect to its H 2 mass, based on modeling the CO emission with a Large Velocity
Gradient (LVG) model. The derived X factor in this system is approximately one order
of magnitude lower than the conventional value if the Galactic [H 2 /CO] ratio is assumed.
Furthermore, according to the LVG model, there is an excessive velocity gradient within
the molecular clouds, considerably exceeding that for virialized clouds. Thus the excess
CO emission in many SIGs may be the result of a large velocity dispersion, which could
increase the radiative efficiency and luminosity of the CO emitting gas. High resolution
interferometer data shows that some high MH \Lambda
2
=LB mergers, such as ARP220 and ARP299,
have relatively broad spectral profiles within the central 1 kpc region, indicating a large
velocity gradient (e.g. Casoli et al. 1999; Downes & Solomon 1998; Scoville et al. 1997;
Aalto et al. 1997; Gao et al. 1997). However, we caution that these observations may not
relate directly to the velocity dispersion inferred from an LVG analysis since the latter refers
to conditions localized within the molecular clouds. In any event, a lower than normal X

-- 23 --
factor may very well be the cause of the excess CO emission in these systems.
9. Summary
We have compiled a complete sample of strongly interacting galaxies in the northern sky
which contains 126 galaxies in 92 systems. CO data have been collected for 75% of these
systems including non­detections (for the subsamples IG3 and IG4 the coverage is more
than 90%). In deriving estimates of the molecular gas mass from the CO luminosity using
the traditional Galactic X factor, we have introduced the parameter MH \Lambda
2
to represent the
H 2 mass in order to emphasize that it is a direct reflection of the CO luminoity. The
principal findings are as follows:
(1) Mergers (IG4) and colliding systems (IG3) on average have significantly higher
MH \Lambda
2
=LB and MH \Lambda
2
=D 2
25 ratios than normal spiral galaxies. If the CO­to­H 2 conversion
factor is the same for all types of galaxies, the enhancement in the CO emission indicates
more molecular gas in merging galaxies.
(2) Strongly interacting pairs (IG2) on average have MH \Lambda
2
=LB and MH \Lambda
2
=D 2
25 ratios
comparable to those of ISGs. But there are still some peculiar IG2 systems that show
evidence of enhanced CO emission. The MH \Lambda
2
=LB and MH \Lambda
2
=D 2
25 ratios are correlated with
interaction strengths.
(3) There is no evidence that enhanced H 2 content in SIGs is produced by the conversion
of atomic gas to molecular gas. Note that this result may be the strongest argument to date
that the enhancement is most readily accounted for by a significantly higher value of the X
factor in SIGs.
(4) The high L IR =LB ratios seen in most interacting pairs, especially those involved in
only mild to moderate strength interactions (groups IG2 and IG3), appears to be due to an

-- 24 --
enhancement in the amount of molecular gas available to fuel star formation, rather than
an increase in the SFE (a conventional X factor is assumed). Only in the most strongly
interacting and merging systems (IG4) does it appear that the SFE is enhanced, indicating
that the strong tidal forces experienced by these systems affects the process of star
formation. This enhancement would be even more pronounced if MH \Lambda
2
is an overestimate of
the true molecular mass.
(5) The primary issue now is whether molecular masses can be reliably estimated for
interacting and star­forming galaxies using a universal X factor. The possibility exists that
this approach yields a significant overestimate. Future effort will have to concentrate on
deriving more precise estimates of the molecular mass, and on understanding the origin of
the suspected enhancement in CO luminosity for these systems.
10. Acknowledgements
This research was supported by a research grant to E.R.S. from the Natural Sciences and
Engineering Research Council of Canada.

-- 25 --
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-- 29 --
Figure Captions
Figure 1a: Distribution of the MB for the SIG sample (upper panel) and the control sample
(lower panel).
Figure 1b: Hubble type distribution of the SIG sample (upper panel) and the control
sample (lower panel).
Figure 2: Histogram of the MH \Lambda
2
=LB ratio for the SIG sample (upper panel) and the control
sample (lower panel).
Figure 3: Correlation between MH \Lambda
2
and LB . Solid circles are merger types (IG3+IG4) and
open circles are IG2 and starts are for isolated galaxies in Young's sample.
Figure 4: Histogram of the MH \Lambda
2
=D 2
25 ratio for the SIG sample (upper panel) and the control
sample (lower panel).
Figure 5: Correlation between MH \Lambda
2
and D 25 . Solid circles are mergers (IG3+IG4), open
circles are IG2 and starts are isolated galaxies in Young's sample.
Figure 6: Histogram of the MH \Lambda
2
=MHI ratio for the SIG sample (upper panel) and the
control sample (lower panel).
Figure 7: Relation between MHI =LB and MH \Lambda
2
=LB for the SIG sample.
Figure 8: Relation between MH \Lambda
2
=MHI and MH \Lambda
2
=LB for (a) the SIG sample; (b) the control
sample.
Figure 9: Relation between MH \Lambda
2
=MHI and LB for the SIG sample.
Figure 10: Relation between L IR =LB and MH \Lambda
2
=LB for the SIG sample.
Figure 11: Relation between L IR =LB and MH \Lambda
2
=MHI for the SIG sample and the control
sample. Solid circles represent mergers (IG3+IG4), open circles represent IG2, and stars

-- 30 --
represent ISGs.
Figure 12: Relation between statistics properties and IG types.
Figure 13: Relation between statistics properties and pair separation Sep=D 25 .

-- 31 --
Table 1. SIG Sample
Name Other Name ARP IG R.A. DEC cz B 0
T D 25 r b Hubble Type
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
NGC317A UGC593 2 00 54 49.8 + 43 31 51 5293 14.60 1.4 1.3 S?
NGC317B UGC594 2 00 54 51.2 + 43 31 19 5334 12.80 1.1 0.5 SB?
UGC816 2 01 13 24.3 + 46 29 01 5188 13.16 1.9 0.9 S?
UGC01449 126 2 01 55 31.1 + 02 50 35 5555 13.58 1.2 0.7 SBm pec
UGC1555 290 2 02 01 02.0 + 14 28 08 3648 13.78 1.5 0.8 S0
UGC1556 290 2 02 01 07.4 + 14 30 00 3534 13.05 2.8 1.4 S0­
NGC0935 UGC1937 276 2 02 25 23.1 + 19 22 35 4142 12.93 1.7 1.1 Scd:
NGC1633 UGC3125 2 04 37 27.3 + 07 15 09 4989 13.73 1.0 0.9 SAB(s)ab
UGC3706N 2 07 06 09.0 + 47 59 41 6090 14.10 0.4 0.2 Sa
UGC3706S 2 07 06 10.3 + 47 59 21 6090 14.10 0.4 0.2 Sa
NGC2744 UGC4757 2 09 01 49.4 + 18 39 25 3428 13.55 1.7 1.1 SBab(s) pec:
NGC2798 283 2 09 14 09.5 + 42 12 37 1739 12.48 2.6 1.0 SB(s)a pec
NGC2799 283 2 09 14 18.1 + 42 12 15 1755 13.72 1.9 0.5 SB(s)m?
NGC2820 UGC4961 2 09 17 43.7 + 64 28 16 1579 11.73 4.1 0.5 SBc(s) pec sp
UGC5304 2 09 50 30.1 + 08 06 08 12308 13.98 1.1 0.9 S?
NGC3226 UGC5617 94 2 10 20 43.6 + 20 09 07 1322 12.45 3.2 2.8 E2: pec
NGC3227 UGC5620 94 2 10 20 47.6 + 20 07 00 1157 11.18 5.4 3.6 SAB(s)a pec
UGC5643 NGC 3212 181 2 10 23 12.1 + 80 04 42 9769 13.73 1.5 1.1 SB?
UGC05832 291 2 10 40 09.6 + 13 43 18 1216 13.58 1.1 1.0 SB?
NGC3395 UGC5931 270 2 10 47 02.3 + 33 14 45 1620 12.09 2.1 1.2 SAB(rs)cd pec
NGC3396 UGC5935 270 2 10 47 09.0 + 33 15 16 1625 12.32 3.1 1.2 IBm pec
UGC5984 2 10 49 29.6 + 30 19 26 10423 14.10 1.9 1.2 S?
NGC3561N UGC6224N 105 2 11 08 31.3 + 28 59 03 8549 14.40 0.7 0.7 SA(r)a pec

-- 32 --
Table 1---Continued
Name Other Name ARP IG R.A. DEC cz B 0
T D 25 r b Hubble Type
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
NGC3561S UGC6224S 105 2 11 08 31.5 + 28 58 07 8811 14.40 0.9 0.9 S0 0 : pec
NGC3786 294 2 11 37 04.7 + 32 11 13 2723 12.97 2.2 1.3 SAB(rs)a pec
NGC3788 294 2 11 37 06.4 + 32 12 35 2699 12.80 2.1 0.7 SAB(rs)ab pec
NGC3808 UGC6643 2 11 38 07.9 + 22 42 18 7078 13.66 1.7 0.9 SAB(rs)c: pec
UGC6854 KGC307A 2 11 50 9.7 + 02 01 10 6128 14.06 1.0 0.9 SB(rs)bc pec
UGC6865 62 2 11 51 02.1 + 43 43 53 5886 13.96 1.1 0.4 SB
NGC3985 UGC6921 2 11 54 06.8 + 48 36 44 950 12.89 1.3 0.8 SB(s)m:
NGC3994 313 2 11 55 01.5 + 32 33 26 3096 13.18 1.0 0.6 SA(r)c pec?
NGC3995 UGC6944 313 2 11 55 09.9 + 32 34 20 3254 12.32 2.8 1.0 SA(m)pec
NGC4211N UGC7277N 106 2 12 13 04.4 + 28 27 20 6599 14.30 1.0 1.0 S0/a pec
NGC4211S UGC7277S 106 2 12 13 05.8 + 28 26 51 6670 15.20 1.4 1.4 S0/a pec
NGC4435 UGC7575 120 2 12 25 08.6 + 13 21 23 781 11.61 2.8 2.0 SB(s)0 0
NGC4438 UGC7574 120 2 12 25 13.5 + 13 17 11 69 10.49 8.5 3.2 SA(s)0/a pec
NGC4485 269 2 12 28 05.5 + 41 58 35 493 12.22 2.3 1.6 IB(s)m pec
NGC4490 269 2 12 28 10.5 + 41 54 56 578 9.81 6.3 3.1 SB(s)0+? sp
NGC4496A UGC7668A 2 12 29 06.2 + 04 12 56 1730 11.78 4.0 3.2 SBm(rs)
UGC7905S UGC7905S 2 12 41 31.9 + 55 10 11 4933 13.97 1.0 0.6 Pec
UGC7905N UGC7905N 2 12 41 33.4 + 55 10 44 4875 13.97 0.0 1.0 S? pec
NGC4676A UGC7938 242 2 12 43 44.2 + 31 00 23 6613 14.10 2.3 0.7 S0 pec?
NGC4676B UGC7939 242 2 12 43 45.3 + 30 59 51 6607 14.70 2.2 0.8 SB(s)0/a pec
NGC4809 277 2 12 52 18.0 + 02 55 27 915 14.04 1.7 0.7 Im pec
NGC4810 277 2 12 52 18.2 + 02 54 40 876 14.75 0.8 0.5 Im pec
UGC8335 2 13 13 41.8 + 62 23 16 9243 14.10 0.9 0.8 S?

-- 33 --
Table 1---Continued
Name Other Name ARP IG R.A. DEC cz B 0
T D 25 r b Hubble Type
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
UGC8357S 2 13 15 25.1 ­ 00 02 56 9944 13.84 1.2 0.5 SB(s)b pec?
NGC5257 UGC8641 240 2 13 37 19.7 + 01 05 40 6798 13.05 1.8 0.9 Sb? pec
NGC5258 UGC8645 240 2 13 37 24.7 + 01 05 10 6784 13.19 1.7 1.1 Sb? pec
NGC5278 UGC8677 239 2 13 39 47.2 + 55 55 19 7541 13.4 1.3 1.0 SA(s)b? pec
NGC5279 UGC8678 239 2 13 39 51.8 + 55 55 29 7580 -- 0.6 0.4 SB(s)a pec
NGC5331S UGC8774S 2 13 49 43.6 + 02 20 53 9906 14.85 1.1 0.7 Sb pec?
NGC5331N UGC8774N 2 13 49 43.7 + 02 21 17 9906 13.83 1.1 0.7 Sb pec?
NGC5394 UGC8898 84 2 13 56 25.2 + 37 41 51 3427 13.29 1.7 1.0 SB(s)b pec
NGC5395 UGC8900 84 2 13 56 29.7 + 37 40 05 3487 12.01 2.9 1.5 SA(s)b pec
NGC5410 UGC8931 2 13 58 49.9 + 41 13 40 3738 13.53 1.5 0.8 SB?
NGC5421N UGC8941N 111 2 13 59 30.0 + 34 04 10 7889 14.10 1.2 0.9 SB?
NGC5421S UGC8941S 2 13 59 30.9 + 34 03 45 7868 15.00 1.2 0.9 E
UGC8972 KPG411B 2 14 01 12.0 + 11 37 03 11338 14.14 1.2 0.7 Sb
UGC8973 KPG411A 2 14 01 12.0 + 11 38 00 11629 14.15 1.2 0.6 Sb
UGC9098W KPG415A 2 14 10 30.4 + 45 55 32 8387 14.2 0.45 0.3 Sb
UGC9098E KPG415B 2 14 10 33.1 + 45 55 32 8303 14.3 1.0 0.4 Sb
NGC5544 UGC9142 199 2 14 14 56.5 + 36 48 11 3077 13.2 1.0 1.0 (R)SB(rs)0/a
NGC5545 UGC9143 199 2 14 14 59.5 + 36 48 25 3071 -- 1.3 0.4 (R)SB(rs)0/a
UGC9221 KPG425A 2 14 21 34.9 + 34 14 06 3860 14.38 0.8 0.7 S?
UGC9222 KPG425B 2 14 21 38.9 + 34 14 35 3539 14.42 0.9 0.3 S0/a
NGC5613 178 2 14 21 59.7 + 35 07 08 8746 -- 1.0 0.8 (R)SAB(r)0+
NGC5614 UGC9226 178 2 14 22 01.1 + 35 05 04 3892 12.37 2.5 2.0 SA(r)ab pec
NGC5679 274 2 14 32 38.7 + 05 34 40 8654 -- 1.1 0.7 Sb

-- 34 --
Table 1---Continued
Name Other Name ARP IG R.A. DEC cz B 0
T D 25 r b Hubble Type
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
UGC9622 2 14 55 04.0 + 19 52 16 4833 13.68 1.2 0.8 S?
NGC5953 UGC9903 91 2 15 32 13.2 + 15 21 40 1965 13.0 1.6 1.3 SAa: pec
NGC5954 UGC9904 91 2 15 32 15.7 + 15 22 10 1959 13.1 1.3 0.6 SAB(rs)cd pec
UGC10770 2 17 12 27.4 + 59 23 23 1108 -- 1.2 0.6 Im pec
ARP311 311 2 17 26 31.4 + 58 31 32 7928 14.18 0.9 0.7 Sab?
UGC10923 2 17 36 22.7 + 86 46 39 7721 13.70 1.2 0.7 S?
NGC6636 UGC11221 2 18 22 03.5 + 66 35 23 4226 -- 2.2 0.4 S?
NGC6745 UGC11391 2 19 00 03.4 + 40 40 20 4545 -- 1.4 0.7 S?
UGC11657 2 20 57 11.7 ­ 02 04 57 5836 14.10 1.1 1.1 Pec
UGC11658 2 20 57 12.8 ­ 02 04 07 5843 13.90 1.4 0.9 SAB(rs)a pec?
NGC7253A UGC11984 278 2 22 17 08.6 + 29 08 48 4583 13.3 1.7 0.5 SB?
NGC7253B UGC11985 278 2 22 17 11.3 + 29 08 06 4493 13.3 1.6 0.5 SB?
UGC12011A 2 22 20 45.5 + 30 40 15 6715 14.0 0.3 0.2 Sb
UGC12011B 2 22 20 48.0 + 30 40 18 6702 14.0 0.2 0.2 Sa
NGC7550 99 2 23 12 46.8 + 18 41 25 5101 13.11 1.4 1.2 SA0­
UGC12856 262 2 23 54 11.9 + 16 32 09 1777 13.80 1.6 1.0 IB(s)m
NGC7805 UGC12908 112 2 23 58 52.7 + 31 09 20 4850 13.90 1.2 0.9 SAB0 0 : pec
NGC7806 UGC12911 112 2 23 58 56.4 + 31 09 51 4768 13.88 1.1 0.8 SA(rs)bc? pec
UGC5615N 2 10 20 36.6 + 53 21 38 9586 13.8 0.8 0.7 Sb
UGC5615S 2 10 20 37.0 + 53 21 05 10028 13.5 0.7 0.5 Sb
UGC9618b 302 2 14 54 48.1 + 24 48 21 9776 15.1 0.6 0.6 Sc
UGC9618a 302 2 14 54 48.6 + 24 49 02 10094 14.2 0.9 0.3 Sb
UGC1720 3 02 11 28.8 + 04 56 33 9061 14.16 0.8 0.6 I?

-- 35 --
Table 1---Continued
Name Other Name ARP IG R.A. DEC cz B 0
T D 25 r b Hubble Type
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
ARP145 145 3 02 20 00.3 + 41 08 35 5425 13.61 1.5 1.3 Ring B
NGC1143 UGC2388 118 3 02 52 36.2 ­ 00 22 47 8459 14.06 0.9 0.8 SAB0­:pec/Rng
NGC1144 UGC2389 118 3 02 52 38.5 ­ 00 23 07 8647 13.21 1.1 0.7 Im pec/Ring B
UGC03730 141 3 07 08 12.1 + 73 33 58 2709 12.88 2.8 1.5 Ring
UGC3829 3 07 20 28.5 + 33 32 24 4028 13.40 1.0 0.8 S?
UGC3852 KPG135A 3 07 23 38.5 + 72 13 50 3605 13.82 0.8 0.6 I?
NGC2444 143 3 07 43 30.6 + 39 09 24 4048 13.91 1.2 0.8 Ring A
NGC2445 UGC4017 143 3 07 43 32.3 + 39 08 25 4002 13.58 1.4 1.1 Ring B
NGC2936 UGC5130 142 3 09 35 08.5 + 02 59 11 6989 13.75 2.1 0.7 I, Ring B
NGC2937 UGC5131 142 3 09 35 09.3 + 02 58 23 6839 14.41 1.3 1.1 E, Ring A
ARP299 UGC6471/72 299 3 11 25 44.2 + 58 50 23 3033 11.85 2.0 1.5 IBm pec
UGC6767 KPG301A 3 11 45 00.0 + 57 55 00 9223 13.90 0.4 0.3 Sb
NGC4410A UGC7535 3 12 23 55.8 + 09 17 48 7219 14.01 1.3 0.8 Sab? pec
NGC4922 UGC8135 3 12 59 00.2 + 29 34 36 7071 13.60 1.8 1.3 I0 pec
NGC5514 UGC9102 3 14 11 10.6 + 07 53 32 7300 13.19 2.2 1.1 SA
NGC5929 UGC9851 90 3 15 24 18.3 + 41 50 43 2561 -- 1.0 0.9 SAB(s)bc I­II
NGC5930 UGC9852 90 3 15 24 20.6 + 41 51 05 2672 12.56 1.7 0.9 Sab: pec
NGC6621 UGC11175 81 3 18 13 14.5 + 68 20 15 6284 13.40 2.1 0.8 Sb: pec
NGC7714 UGC12699 284 3 23 33 41.2 + 01 52 42 2799 12.62 1.9 1.4 SB:(s)b? pec
UGC12914 3 23 59 04.1 + 23 12 23 4371 12.51 2.3 1.3 (R)Scd(r):pec
UGC12915 3 23 59 08.6 + 23 12 59 4336 13.01 1.5 0.5 S?
NGC0520 157 4 01 21 59.4 + 03 32 13 2217 11.97 4.5 1.8 S pec
NGC0523 UGC979 158 4 01 22 29.9 + 33 45 53 4750 12.60 2.5 0.7 Pec

-- 36 --
Table 1---Continued
Name Other Name ARP IG R.A. DEC cz B 0
T D 25 r b Hubble Type
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
NGC2623 UGC4509 243 4 08 35 25.3 + 25 55 35 5535 13.19 2.4 0.7 P: Trip Coll
NGC3239 263 4 10 22 23.3 + 17 24 50 753 11.54 5.0 3.3 Ib(s)m pec
NGC3303 UGC5773 192 4 10 34 17.9 + 18 23 48 6165 14.00 3.0 2.1 Pec
Mrk231 UGC8058 4 12 54 05.0 + 57 08 38 12642 13.97 1.3 1.0 SA(rs)c?:pec
ARP193 193 4 13 18 17.0 + 34 24 04 6892 14.29 1.5 1.1 Im: pec
NGC5256 UGC8632 4 13 36 14.9 + 48 31 47 8353 13.90 0.4 0.3 Compact pec
NGC5860 UGC9717 4 15 04 44.4 + 42 50 00 5398 14.20 0.7 0.7 S?
ARP220 220 4 15 32 47.3 + 23 40 06 5434 13.55 1.5 1.2 S?
NGC6052 UGC10182 209 4 16 03 01.3 + 20 40 43 4716 13.40 0.9 0.7 Sc
NGC6090 UGC10267 4 16 10 24.0 + 52 35 00 8785 13.89 1.7 0.7 S0?
NGC6240 UGC10592 4 16 50 27.8 + 02 28 58 7339 13.17 2.1 1.1 I0:pec
Note. --- (1) The most common name (2) Other Name (3) ARP number (4) IG type as defined
in x2.1 (5) R.A. (1950) (6) DEC (1950) (7) Optical radial velocity cz from RC3 (8) B 0
T (9) D 25 in
arcmin 10) r b in arcmin (11) Hubble Type

-- 37 --
Table 2. CO Data and Derived Quantities for IGs Observed by Us
Name D SCO log(L B ) log(M H \Lambda
2
) MH \Lambda
2
=LB log [ M(H \Lambda
2
)
D 2
25
] IG type
(1) (2) (3) (4) (5) (6) (7) (8)
ARP205 18.9 96 9.95 8.58 0.04 5.60 1
UGC0717 148.6 28 10.96 9.83 0.07 6.21 1
UGC0365 63.6 240 10.84 10.03 0.15 6.66 1
UGC1063 78.0 48 10.48 9.51 0.11 6.79 1
UGC1810 102.4 64 11.05 9.87 0.07 6.24 1
UGC4264 53.8 66 10.60 9.32 0.05 6.14 1
UGC5265 22.9 !23 9.47 !8.12 !0.04 !6.32 1
UGC5269 22.9 30 9.76 8.24 0.03 5.67 1
UGC5600 39.5 43 9.73 8.87 0.14 6.45 1
UGC5609 39.5 !31 9.69 !8.69 !0.10 !6.38 1
UGC8529 39.2 367 10.24 9.79 0.36 7.17 1
UGC10033 45.3 89 10.44 9.30 0.07 6.60 1
UGC11137 32.5 90 10.29 9.02 0.05 6.56 1
UGC11414 103.1 !31 10.94 !9.56 !0.04 !6.52 1
UGC11695 130.9 104 10.80 10.29 0.31 6.84 1
UGC12066 77.2 37 10.44 9.38 0.09 6.60 1
UGC12456 67.8 31 10.64 9.23 0.04 6.31 1
UGC12457 67.8 215 10.75 10.00 0.18 6.55 1
ARP126N 74.5 69 10.50 9.63 0.13 6.80 2
ARP126S 74.5 45 10.74 9.43 0.05 7.36 2
ARP262N 25.8 !53 9.42 ! 8.59 ! 0.15 ! 6.43 2
ARP308N 72.6 !63 10.71 ! 9.55 ! 0.07 ! 6.15 2
ARP308S 72.6 !41 10.76 ! 9.39 ! 0.04 ! 6.50 2
ARP311 108.5 48 10.59 9.79 0.16 6.89 2
UGC0593 73.2 !14 10.08 !8.91 !0.07 !5.97 2

-- 38 --
Table 2---Continued
Name D SCO log(L B ) log(M H \Lambda
2
) MH \Lambda
2
=LB log [ M(H \Lambda
2 )
D 2
25
] IG type
(1) (2) (3) (4) (5) (6) (7) (8)
UGC0594 73.2 208 10.80 10.09 0.19 7.35 2
UGC0813 70.2 63 10.40 9.52 0.13 6.76 2
UGC0816 70.2 161 10.64 9.96 0.21 6.76 2
UGC3706N 80.8 12 10.37 8.94 0.04 6.99 2
UGC3706S 80.8 10 10.37 8.86 0.03 6.91 2
UGC5304 154.1 14 10.98 9.56 0.04 6.18 2
UGC5931 21.4 99 10.01 8.69 0.05 6.47 2
UGC5935 21.4 112 9.93 8.76 0.07 6.18 2
UGC5984 138.6 46 10.84 9.99 0.14 6.22 2
UGC6643 93.8 62 10.67 9.78 0.13 6.45 2
UGC6865 78.4 154 10.40 10.02 0.42 7.22 2
UGC7277N 88.3 !9 10.36 !8.88 !0.03 !6.07 2
UGC7277S 88.3 47 10.01 9.61 0.40 6.43 2
UGC7905N 67.0 !40 10.26 !9.30 !0.11 !8.11 2
UGC7905S 67.0 !24 10.26 !9.07 !0.07 !6.49 2
UGC8357S 132.6 122 10.90 10.37 0.30 0.00 2
UGC8677 102.6 76 10.85 9.94 0.12 6.70 2
UGC8774N 133.8 70 11.00 10.23 0.17 7.75 2
UGC8774S 133.8 151 10.42 10.39 0.93 7.74 2
UGC8898 46.9 201 10.21 9.68 0.29 6.96 2
UGC8900 46.9 548 10.74 10.13 0.25 6.93 2
UGC8931 51.1 !43 10.20 !9.09 !0.08 !6.40 2
UGC8941N 106.0 93 10.60 10.06 0.29 6.92 2
UGC8941S 106.0 7 10.24 8.94 0.05 5.80 2
UGC9142 42.3 40 10.16 8.90 0.05 6.72 2

-- 39 --
Table 2---Continued
Name D SCO log(L B ) log(M H \Lambda
2
) MH \Lambda
2
=LB log [ M(H \Lambda
2 )
D 2
25
] IG type
(1) (2) (3) (4) (5) (6) (7) (8)
UGC9226 52.9 490 10.69 10.18 0.31 7.01 2
UGC9903 27.1 365 9.86 9.47 0.41 7.27 2
UGC9904 27.1 73 9.82 8.77 0.09 6.75 2
UGC10923 105.4 129 10.76 10.20 0.28 7.07 2
UGC11657 79.8 !9 10.35 !8.80 !0.03 !5.99 2
UGC11658 79.8 25 10.44 9.24 0.06 6.22 2
UGC11984 63.8 180 10.48 9.91 0.27 6.91 2
UGC11985 63.8 47 10.47 9.31 0.07 6.38 2
UGC12011A 89.5 !15 10.50 ! 9.12 !0.04 !7.34 2
UGC12011B 89.5 !15 10.50 ! 9.12 !0.04 !7.69 2
UGC12908 66.6 !36 10.29 ! 9.25 !0.09 !6.51 2
UGC12911 66.6 46 10.28 9.34 0.12 6.69 2
ARP141 37.8 !56 10.19 ! 8.94 ! 0.06 ! 5.97 3
ARP142S 90.3 !45 10.33 ! 9.60 ! 0.19 ! 6.13 3
ARP142N 90.3 335 10.65 10.49 0.68 7.41 3
ARP90 37.4 148 10.31 9.36 0.11 6.82 3
UGC11175 86.2 255 10.70 10.32 0.41 6.88 3
UGC12914 60.4 420 10.75 10.23 0.30 7.01 3
UGC12915 60.4 381 10.55 10.18 0.43 7.34 3
UGC3829 53.3 92 10.29 9.46 0.15 7.08 3
UGC3852 49.4 23 10.05 8.79 0.06 6.67 3
UGC7535 95.7 !34 10.55 !9.53 !0.10 !6.42 3
UGC9102 97.4 111 10.89 10.06 0.15 6.47 3
NGC3239 8.9 ! 66 9.47 !7.76 !0.02 !5.54 4
UGC5773 81.9 !7 10.42 !8.71 !0.02 !6.40 4

-- 40 --
Table 2---Continued
Name D SCO log(L B ) log(M H \Lambda
2
) MH \Lambda
2
=LB log [ M(H \Lambda
2
)
D 2
25
] IG type
(1) (2) (3) (4) (5) (6) (7) (8)
UGC0979 65.5 90 10.78 9.63 0.07 6.27 4
UGC9717 73.7 37 10.25 9.34 0.12 6.99 4
ARP137 41.9 !49 9.95 !8.98 !0.11 !6.80 5
ARP162 18.3 !75 9.97 !8.44 !0.03 !5.90 5
ARP165 66.3 !64 10.65 !9.49 !0.07 !6.41 5
NGC3509 101.2 91 11.09 10.01 0.08 6.43 5
NGC3921 79.1 80 10.83 9.74 0.08 6.37 5
ARP164 71.3 !54 10.65 ! 9.48 ! 0.07 ! 6.29 5
ARP222 25.3 !157 10.49 ! 9.04 ! 0.04 ! 5.97 5
ARP223 47.5 !49 10.65 ! 9.08 ! 0.03 ! 6.08 5
Note. --- Col.(1) Galaxy name. Col.(2) Distance in Mpc, computed using the mean
optical radial velocity of each system from Table 1 and assuming H 0 = 75 km s \Gamma1 Mpc \Gamma1 .
Col.(3) Global CO flux, in units of Jy km/sec. Corrections for flux outside the telescope
beam are applied. Upper limits are 3 oe. Col.(4) Logarithm of the blue luminosity in units of
L fi , computed from values of B 0
T in Table 1 and assuming MB fi = +5:48. Col.(5) Logarithm
of the global molecular gas mass in the units of M fi , computed using MH \Lambda
2
= 1:1 \Theta 10 4 D 2 SCO .
Col.(6) Ratio of the global molecular gas mass to the blue luminosity in the units of L fi =M fi .
Col.(7) Logarithm of the ratio of the global molecular gas mass to the optical area D 2
25 in
units of M fi /kpc 2 . The optical area is computed using the D 25 from Table 1. Col.(8) IG
type as defined in x2.1.

-- 41 --
Table 3. CO Data for IGs in the Literature
Name D SCO log(L B ) log(MH \Lambda
2
) MH \Lambda
2
=LB log [ M(H \Lambda
2 )
D25 ] IG type Reference
(1) (2) (3) (4) (5) (6) (7) (8) (9)
ARP295a 91.5 43 10.47 9.60 0.13 6.19 1 16
ARP295b 91.5 89 10.81 9.91 0.13 7.15 1 16
NGC1055 13.3 2800 10.15 9.74 0.38 6.85 1 15
NGC2276 32.1 800 10.51 9.96 0.28 7.25 1 15
NGC3169 14.6 1300 10.14 9.48 0.22 6.83 1 15
NGC3718 13.3 !315 9.96 !8.79 !0.06 !5.82 1 15
NGC3800 43.1 200 10.39 9.61 0.17 6.81 1 5
NGC3893 13.0 540 10.15 9.00 0.07 6.61 1 15
NGC4017 45.8 83 10.31 9.28 0.09 6.52 1 3
NGC4298 14.5 660 9.87 9.18 0.21 6.85 1 15
NGC4302 14.1 620 9.95 9.13 0.15 6.43 1 15
NGC4567 29.5 500 10.42 9.68 0.18 6.17 1 15
NGC4568 29.4 1050 10.66 10.00 0.22 6.12 1 15
NGC4647 18.2 600 9.99 9.34 0.22 6.67 1 15
NGC5000 75.1 60 10.57 9.57 0.10 6.60 1 3
NGC5194 6.4 9450 10.34 9.63 0.20 7.01 1 15
NGC5195 6.4 240 9.64 8.03 0.02 6.03 1 15
NGC6285 74.7 83 10.36 9.71 0.22 7.12 1 9
NGC6286 74.7 290 10.31 10.25 0.86 7.35 1 9
NGC7603 119.7 249 10.91 10.59 0.48 7.16 1 8
NGC7674 118.8 350 10.90 10.74 0.69 7.51 1 15
NGC7771 59.3 805 10.74 10.49 0.57 7.24 1 13
UGC480W 151.2 160 11.36 10.60 0.17 6.97 1 2

-- 42 --
Table 3---Continued
Name D S CO log(L B ) log(MH \Lambda
2
) MH \Lambda
2
=LB log [ M(H \Lambda
2 )
D 25
] IG type Reference
(1) (2) (3) (4) (5) (6) (7) (8) (9)
UGC3031 62.3 !28 9.94 !9.08 !0.14 !6.15 1 2
UGC3032 62.2 !36 10.12 ! 9.19 !0.12 !6.59 1 2
UGC4718 41.7 !48 9.99 !8.96 !0.09 !6.50 1 2
UGC4744 31.8 55 9.84 8.79 0.09 6.70 1 2
ARP302a 132.0 332 10.94 10.08 0.78 7.73 2 9
ARP302b 132.0 59 10.58 10.94 0.32 7.33 2 9
NGC0935 56.3 170 10.52 9.77 0.18 6.88 2 5
NGC2798 23.2 440 9.93 9.42 0.31 6.91 2 15
NGC2799 23.9 60 9.46 8.58 0.13 6.29 2 15
NGC2820 22.2 100 10.19 8.73 0.03 5.89 2 15
NGC3226 16.6 !135 9.65 !8.62 !0.09 !6.24 2 15
NGC3227 14.4 1418 10.04 9.51 0.30 6.80 2 10
NGC3786 36.2 157 10.12 9.35 0.17 6.62 2 10
NGC3788 36.2 !49 10.18 8.84 0.046 !6.16 2 10
NGC3995 43.4 55 10.54 9.06 0.03 5.96 2 3
NGC4438 13.3 357 10.63 9.61 0.10 5.81 2 6
NGC4490 10.9 480 10.11 8.56 0.03 6.43 2 15
NGC4676A 88.5 71 10.45 9.79 0.22 6.24 2 3
NGC4676B 88.2 51 10.20 9.64 0.27 6.14 2 3
UGC4757 44.4 !55 10.07 !9.08 !0.10 !6.39 2 2
UGC6224N 117.1 125 10.57 10.28 0.51 7.52 2 2
UGC8335E 124.1 51 10.74 9.93 0.16 7.28 2 2
UGC8641 90.1 258 10.88 10.36 0.30 7.01 2 2

-- 43 --
Table 3---Continued
Name D S CO log(L B ) log(MH \Lambda
2
) MH \Lambda
2
=LB log [ M(H \Lambda
2 )
D 25
] IG type Reference
(1) (2) (3) (4) (5) (6) (7) (8) (9)
UGC8645 89.6 296 10.82 10.42 0.39 7.12 2 2
ARP143 53.6 95 10.22 9.48 0.18 6.80 3 10
ARP145 69.0 44 10.43 9.36 0.09 6.41 3 10
ARP299 42.1 900 10.69 10.23 0.36 7.45 3 15
NGC1144 115.3 318 11.03 10.67 0.43 7.53 3 9
NGC5929 35.8 104 10.14 9.17 0.11 7.13 3 10
NGC6090 119.4 200 10.79 10.50 0.51 6.91 3 15
NGC7714 39.0 130 10.33 9.34 0.10 6.67 3 15
UGC1720 121.3 88 10.70 10.15 0.29 7.25 3 2
UGC8135 94.7 107 10.22 10.02 0.63 6.63 3 2
ARP193 92.5 220 10.40 10.32 0.83 7.10 4 15
ARP220 73.5 556 10.48 10.52 1.09 7.51 4 4
Mrk 231 165.7 100 11.04 10.48 0.28 6.65 4 15
NGC520 30.3 1260 10.53 10.10 0.37 6.91 4 15
NGC5256 111.4 210 10.73 10.46 0.54 8.23 4 15
NGC6052 64.3 190 10.44 9.94 0.31 7.49 4 11
NGC6240 101.3 313 10.88 10.55 0.46 6.93 4 4
NGC828 71.6 430 10.87 10.38 0.33 6.82 5 12
NGC2623 72.9 170 10.64 10.00 0.23 6.53 5 15
NGC3310 13.9 140 10.10 8.47 0.02 6.30 5 15
NGC3597 46.8 141 10.20 9.53 0.22 6.70 5 14
NGC3656 39.6 162 10.09 9.45 0.23 6.92 5 14
NGC4194 34.5 251 10.12 9.52 0.25 7.00 5 4

-- 44 --
Table 3---Continued
Name D S CO log(L B ) log(MH \Lambda
2
) MH \Lambda
2
=LB log [ M(H \Lambda
2
)
D25 ] IG type Reference
(1) (2) (3) (4) (5) (6) (7) (8) (9)
Note. --- Col.(1) Galaxy name. Col.(2) Distance in Mpc, computed using the mean optical
radial velocity of each system from Table 1 and assuming H 0 = 75 km s \Gamma1 Mpc \Gamma1 . Col.(3)
Global CO flux, in units of Jy km/sec. Corrections for flux outside the telescope beam are
applied. Upper limits are 3 oe. Col.(4) Logarithm of the blue luminosity in units of L fi ,
computed from values of B 0
T in Table 1 and assuming MB fi = +5:48. Col.(5) Logarithm of
the global molecular gas mass in the units of M fi , computed using MH \Lambda
2
= 1:1 \Theta 10 4 D 2 S CO .
Col.(6) Ratio of the global molecular gas mass to the blue luminosity in the units of L fi =M fi .
Col.(7) Logarithm of the ratio of the global molecular gas mass to the optical area D 2
25 in units
of M fi /kpc 2 . The optical area is computed using the D 25 from Table 1. Col.(8) IG type as
defined in x2.1. Col.(9) Reference for the CO data.
References. --- (1) Braine et al. 1993; (2) Bushouse et al. 1999; (3) Casoli et al. 1996 ; (4)
Casoli et al. 1992; (5) Chini et al. 1996; (6) Combes et al. 1988; (7) Dupraz et al. 1990; (8)
Elfhag et al. 1996; (9) Gao 1997; (10) Maiolino et al. 1997; (11) Sage 1993; (12) Sanders et al.
1991; (13) Solomon et al. 1988; (14) Wiklind et al. 1995 (15) Young et al. 1996; (16) Hibbard
1995.

-- 45 --
Table 4. Global Quantities for Interacting Systems
Name IG Sep=D 25 S CO HI S 60 S 100 log(L B ) log(MH \Lambda
2
) log(MHI ) log(L IR )
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
ARP205 1 -- 96 64.80 5.87 10.65 9.95 8.58 9.74 9.56
NGC2276 1 2.60 800 24.50 14.20 29.80 10.56 10.02 9.84 10.49
NGC3169 1 1.53 1300 96.00 8.06 22.16 10.10 9.45 9.65 9.52
NGC3718 1 1.48 !315 151.10 0.78 3.42 10.05 !8.79 9.88 8.66
NGC3800 1 0.65 200 4.80 -- -- 10.39 9.61 9.32 --
NGC4017 1 3.23 83 25.80 1.89 5.41 10.31 9.28 10.11 9.93
NGC4298/302 1 0.42 1280 64.00 5.74 24.03 10.15 9.40 9.43 9.43
NGC4647 1 0.63 600 8.20 5.63 16.69 9.72 9.07 8.54 9.34
NGC5000 1 -- 60 5.79 0.96 2.39 10.57 9.57 9.89 10.03
NGC5194 1 0.42 9690 220.00 84.48 233.32 10.42 9.64 9.33 9.86
NGC6285/6 1 1.15 373 -- 9.50 23.32 10.64 10.36 -- 11.02
NGC7603 1 -- 249 6.80 0.85 2.04 10.91 10.59 10.36 10.38
NGC7674 1 1.99 350 7.38 5.56 8.79 10.91 10.75 10.40 11.12
NGC7771 1 2.22 805 15.20 20.15 46.64 10.75 10.50 10.11 11.15
UGC0365 1 0.14 240 10.10 3.00 -- 10.84 10.03 9.98 --
UGC0717 1 1.38 28 0.68 0.45 1.60 10.96 9.83 9.55 10.38
UGC1063 1 1.12 48 -- 1.48 3.73 10.48 9.51 -- 10.26
UGC1810 1 0.08 64 16.03 2.00 -- 11.05 9.87 10.60 --
UGC3031/2 1 0.77 !64 2.42 0.35 0.93 10.34 !9.44 9.34 9.44
UGC4264 1 0.70 66 25.30 3.14 6.61 10.60 9.32 10.24 10.22
UGC4718 1 0.30 !48 16.70 1.97 2.71 9.99 !8.96 9.84 9.73
UGC4744 1 0.64 55 15.90 1.22 3.33 9.85 8.79 9.58 9.41
UGC480W 1 0.94 160 1.93 1.69 4.00 11.36 10.61 10.02 10.88

-- 46 --
Table 4---Continued
Name IG Sep=D 25 S CO HI S 60 S 100 log(L B ) log(MH \Lambda
2
) log(MHI ) log(L IR )
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
UGC5265 1 2.29 !23 30.20 2.27 4.54 9.46 !8.12 9.57 9.33
UGC5269 1 0.95 30 18.80 2.27 4.54 9.76 8.24 9.37 9.33
UGC5600 1 0.94 43 10.30 3.46 14.96 9.73 8.87 9.58 10.17
UGC5609 1 1.02 !31 11.40 2.00 -- 9.69 !8.69 9.59 --
UGC8529 1 2.27 367 3.88 6.91 13.57 10.24 9.79 9.15 10.28
UGC8929S 1 0.91 6 -- 0.47 0.71 10.40 8.91 -- 9.97
UGC10033 1 0.89 89 20.20 4.43 8.05 10.44 9.30 9.99 10.20
UGC11137 1 -- 90 11.89 2.55 6.25 10.29 9.02 9.47 9.73
UGC11414 1 1.12 !31 -- 7.38 11.36 10.94 !9.56 -- 11.10
UGC11695 1 0.90 104 6.99 0.56 1.12 10.80 10.29 10.45 10.24
UGC12066 1 0.20 37 11.90 1.10 2.37 10.44 9.38 10.22 10.09
UGC12457 1 1.80 215 4.57 1.73 4.79 10.75 10.04 9.70 10.22
ARP262 2 0.50 !53 -- 0.32 !1.10 9.42 !8.73 -- !8.84
ARP126 2 0.35 114 9.53 4.80 8.22 11.04 9.84 10.10 10.65
ARP302 2 0.77 391 4.66 6.94 15.52 11.10 10.90 10.31 11.39
ARP311 2 -- 48 -- 0.04 0.19 10.59 9.79 -- 9.13
NGC0935 2 0.62 170 18.20 4.41 11.72 10.52 9.77 10.13 10.46
NGC2798/9 2 0.60 500 9.60 22.55 28.28 10.03 9.46 9.07 10.25
NGC2820 2 0.49 100 58.50 4.07 10.43 10.19 8.73 9.83 9.61
NGC3226/7 2 0.40 1508 23.14 8.60 17.98 10.19 9.54 9.05 9.51
NGC3786 2 0.63 157 15.40 -- -- 10.45 9.35 9.68 --
NGC3788 2 0.66 ! 49 12.91 3.00 -- 10.18 !8.84 9.59 --
NGC3995 2 0.71 55 46.40 7.57 14.27 10.54 9.06 10.31 10.40

-- 47 --
Table 4---Continued
Name IG Sep=D 25 S CO HI S 60 S 100 log(L B ) log(MH \Lambda
2
) log(MHI ) log(L IR )
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
NGC4438 2 0.52 357 10.20 4.24 12.46 10.44 8.84 8.63 9.21
NGC4490 2 0.78 480 335.00 46.94 78.37 10.34 8.80 9.97 9.97
NGC4567/8 2 0.13 1550 33.30 21.00 55.93 10.86 9.48 9.15 9.89
NGC4676 2 0.25 122 6.68 2.81 5.12 10.64 10.02 10.09 10.58
UGC4757 2 -- !55 9.54 1.45 2.65 10.07 !9.08 9.65 9.70
NGC5929 2 0.46 104 3.10 9.22 13.68 10.14 9.17 8.97 10.27
UGC0594 2 0.52 208 6.90 9.02 13.54 10.80 10.09 9.94 10.89
UGC0813/6 2 0.46 224 10.50 2.74 6.53 10.84 10.10 10.10 10.44
UGC3706 2 0.97 22 1.61 0.34 0.62 10.67 9.20 9.39 9.58
UGC5304 2 0.53 14 -- 0.98 1.60 10.98 9.56 -- 10.58
UGC5931/5 2 0.40 211 37.80 10.29 17.43 10.27 9.02 9.61 9.89
UGC5984 2 0.60 46 2.66 0.15 0.77 10.84 9.99 10.08 9.94
UGC6224 2 1.33 125 1.65 4.30 7.80 10.57 10.27 9.73 11.01
UGC6643 2 0.61 62 1.12 3.36 6.24 10.67 9.78 9.37 10.71
UGC6865 2 0.53 154 8.04 2.45 5.77 10.40 10.02 10.07 10.46
UGC7277S 2 0.58 47 1.22 0.66 1.03 10.01 9.61 9.36 9.93
UGC7905 2 0.59 !64 7.58 2.18 2.92 10.56 !9.50 9.90 10.18
UGC8335 2 0.67 119 1.20 10.66 11.80 10.74 10.30 9.64 11.38
UGC8357S 2 0.82 122 -- 2.87 4.39 10.90 10.37 -- 10.91
UGC8641/5 2 0.73 554 8.97 9.39 19.42 11.15 10.69 10.23 11.14
UGC8677 2 0.40 76 1.71 1.86 4.70 10.85 9.94 9.63 10.59
UGC8774 2 0.86 221 5.23 5.63 10.75 11.10 10.73 10.43 11.34
UGC8898/900 2 0.65 749 51.40 9.92 19.16 10.85 10.27 10.44 10.60

-- 48 --
Table 4---Continued
Name IG Sep=D 25 S CO HI S 60 S 100 log(L B ) log(MH \Lambda
2
) log(MHI ) log(L IR )
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
UGC8931 2 0.73 !43 11.90 0.75 2.34 10.20 !9.09 9.87 9.64
UGC8941 2 0.26 100 2.22 0.82 1.69 10.76 10.09 9.77 10.22
UGC9142 2 0.52 40 4.03 0.82 1.82 10.16 8.90 9.23 9.44
UGC9226 2 0.38 490 4.20 1.48 5.50 10.69 10.18 9.44 10.01
UGC9903/4 2 0.49 438 9.70 11.58 20.51 10.14 9.55 9.23 10.16
UGC10923 2 0.60 129 4.33 4.95 10.01 10.76 10.20 10.06 11.00
UGC11657/8 2 0.66 9 7.24 1.09 1.18 10.70 8.80 10.04 10.00
UGC11984/5 2 0.40 227 7.79 6.70 12.90 10.77 10.00 9.86 10.67
UGC12011 2 1.09 !30 1.47 1.50 2.37 10.80 !9.42 9.44 10.29
UGC12911 2 0.77 46 12.69 0.38 0.81 10.28 9.34 10.12 9.49
ARP141 3 -- !56 10.68 0.19 0.81 10.19 !8.94 9.53 !8.83
ARP142N 3 0.64 335 -- 2.00 4.70 11.03 10.49 -- 10.51
ARP143 3 0.69 95 14.07 3.25 6.08 10.22 9.48 9.98 10.21
ARP145 3 0.43 44 1.55 0.77 1.39 10.43 9.36 9.24 9.80
ARP299 3 0.20 900 16.10 120.81 115.33 10.69 10.23 9.81 11.46
ARP90 3 0.27 148 4.10 9.42 13.80 10.31 9.36 9.13 10.32
NGC1143/4 3 0.32 318 2.83 5.43 11.53 11.03 10.67 9.95 11.12
NGC7714 3 1.04 130 26.00 11.22 11.66 10.33 9.34 9.97 10.38
UGC11175 3 0.32 255 9.04 6.60 11.61 10.70 10.32 10.20 10.92
UGC12914/5 3 0.48 1012 15.30 6.13 13.25 10.96 10.61 10.12 10.62
UGC1720 3 0.25 88 2.96 5.32 8.15 10.69 10.15 10.01 11.10
UGC3829 3 -- 92 2.30 7.15 8.48 10.29 9.46 9.19 10.48
UGC3852 3 0.27 23 9.50 2.34 3.15 10.05 8.79 9.74 9.94

-- 49 --
Table 4---Continued
Name IG Sep=D 25 S CO HI S 60 S 100 log(L B ) log(MH \Lambda
2
) log(MHI ) log(L IR )
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
UGC7535 3 0.25 !34 0.74 3.00 -- 10.55 !9.53 9.20 --
UGC8135 3 0.20 107 0.70 5.73 7.54 10.71 10.02 9.17 10.89
UGC9102 3 0.13 111 1.27 1.39 2.16 10.89 10.06 9.45 10.33
ARP193 4 0.11 220 1.30 16.46 24.72 10.40 10.32 9.42 11.35
ARP220 4 0.14 556 3.80 107.43 119.34 10.48 10.52 9.69 11.92
Mrk231 4 0.03 100 3.10 9.14 13.69 11.04 10.48 10.30 11.60
NGC0520 4 0.14 1260 38.80 31.70 47.32 10.54 10.10 9.92 10.67
NGC2623 4 0.07 170 3.90 24.02 26.08 10.62 9.98 9.67 11.25
NGC3239 4 0.13 66 79.50 4.09 6.32 9.48 7.76 9.17 8.72
NGC5256 4 0.40 210 -- 7.19 10.76 10.73 10.46 -- 11.16
NGC6052 4 -- 190 10.10 7.04 10.96 10.44 9.94 9.99 10.67
NGC6240 4 0.08 313 4.18 22.80 27.03 10.88 10.55 10.01 11.54
UGC5773 4 0.39 ! 7 1.79 0.34 1.05 10.42 !8.71 9.45 9.71
UGC9717 4 0.46 37 -- 1.64 3.02 10.25 9.34 -- 10.19
UGC979 4 -- 90 16.57 2.00 4.57 10.78 9.63 10.22 10.21
UGC10267 4 0.08 179 5.29 6.48 9.42 10.79 10.45 10.25 11.16
ARP137 5 -- !49 1.60 -- -- 9.95 !8.98 8.82 --
ARP162 5 -- !75 1.01 0.19 !0.91 9.97 !8.44 7.90 !8.27
ARP164 5 -- !54 !3.00 -- -- 10.65 !9.47 !9.54 --
ARP165 5 -- !64 -- -- -- 10.65 !9.49 -- --
NGC0828 5 -- 430 13.87 11.58 25.50 10.87 10.38 10.22 11.05
NGC3310 5 -- 140 63.06 34.91 45.24 10.12 8.49 9.48 10.03
NGC3509 5 -- 91 8.70 1.09 3.41 11.09 10.01 10.32 10.40

-- 50 --
Table 4---Continued
Name IG Sep=D 25 S CO HI S 60 S 100 log(L B ) log(MH \Lambda
2
) log(MHI ) log(L IR )
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)
NGC3597 5 -- 141 -- 13.14 16.54 10.20 9.53 -- 10.64
NGC3656 5 -- 162 4.90 2.47 5.27 10.09 9.45 9.26 9.86
NGC3921 5 -- 80 5.65 0.77 1.23 10.83 9.74 9.92 9.90
NGC4194 5 -- 251 9.20 23.79 25.38 10.12 9.52 9.41 10.61
NGC7252 5 -- 82 5.42 3.97 7.19 10.82 9.57 9.72 10.45
Note. --- Col.(1) Galaxy name. Col.(2) IG type as defined in x2.1. Col.(3) Pair separations
Sep normalized by the primary's diameter D 25 . Col.(4) Total CO flux for the whole system, in
units of Jy km/sec. For close pairs the sum of the two galaxies are used. Upper limits are 3 oe.
Col.(5) HI flux for the system, in units of Jy km/sec. Col.(6) FIR flux at 60 ¯m, in units of Jy.
Col.(7) FIR flux at 100 ¯m, in units of Jy. Col.(8) Logarithm of the total blue luminosity in
units of L fi , computed from values of B 0
T in Table 1 and assuming MB fi = +5:48. For pairs this
is the sum of both galaxies. Col.(9) Logarithm of the global molecular gas mass in the units of
M fi , computed using MH \Lambda
2
= 1:1 \Theta 10 4 D 2 SCO . Col.(10)Logarithm of the global HI mass in the
units of M fi , computed using MHI = 2:36 \Theta 10 5 D 2 SHI . Col.(11) Logarithm of FIR luminosity,
computed from the 60 ¯m and 100 ¯m flux densities using L IR = 3:94 \Theta 10 5 D 2 [2:58S 60 +S 100 ]

-- 51 --
Table 5. Mean and Median Properties of IG subsamples
Quanty a IG0 IG1 IG2 IG3 IG4 IG5 SIG
N b 59 42 61 18 13 11 92
D (Mpc) 23.9 \Sigma 1.0 54.6 \Sigma 6.8 71.9 \Sigma 5.8 71.8 \Sigma 7.2 81.5 \Sigma10.9 54.6 \Sigma 8.0 73.3 \Sigma 7.6
logL B (L fi ) 10.18 \Sigma 0.04 10.31 \Sigma 0.07 10.38 \Sigma 0.05 10.52 \Sigma 0.06 10.53 \Sigma 0.11 10.45 \Sigma 0.14 10.43 \Sigma 0.04
Size (Kpc) 23.69 \Sigma 1.31 29.60 \Sigma 2.28 27.80 \Sigma 1.83 33.12 \Sigma 3.38 37.60 \Sigma 6.29 33.30 \Sigma 6.28 30.30 \Sigma 1.67
MH \Lambda
2
=LB 0.13 \Sigma 0.01 0.18 \Sigma 0.03 0.19 \Sigma 0.02 0.23 \Sigma 0.04 0.37 \Sigma 0.08 0.11 \Sigma 0.03 0.22 \Sigma 0.02
(M fi =L fi ) 0.12 0.13 0.14 0.21 0.31 0.08 0.15
MH \Lambda
2
=D 2
25 5.29 \Sigma 0.46 6.10 \Sigma 1.03 8.35 \Sigma 1.50 10.73 \Sigma 2.38 22.42 \Sigma12.14 4.20 \Sigma 1.01 10.71 \Sigma 2.04
(M fi =Kpc 2 ) 4.04 3.54 5.03 6.39 8.02 2.96 5.44
MH \Lambda
2
=MHI 1.34 \Sigma 0.20 1.12 \Sigma 0.21 1.31 \Sigma 0.21 2.18 \Sigma 0.52 2.28 \Sigma 0.80 1.23 \Sigma 0.32 1.72 \Sigma 0.22
0.81 0.66 0.89 1.68 1.51 1.27 1.36
L IR =MH \Lambda
2
4.34 \Sigma 0.63 5.65 \Sigma 1.24 5.51 \Sigma 0.71 6.53 \Sigma 1.32 9.76 \Sigma 1.75 9.76 \Sigma 3.85 6.54 \Sigma 0.63
(M fi =L fi ) 3.30 3.80 4.14 5.54 9.11 6.09 4.97
L IR =LB 0.54 \Sigma 0.12 0.66 \Sigma 0.12 0.80 \Sigma 0.13 1.31 \Sigma 0.37 4.53 \Sigma 2.05 1.18 \Sigma 0.41 1.62 \Sigma 0.43
0.41 0.38 0.55 1.01 2.37 0.70 0.87
L IR =MGAS \Lambda c 1.90 \Sigma 0.29 1.13 \Sigma 0.16 2.24 \Sigma 0.33 3.57 \Sigma 0.89 6.39 \Sigma 1.98 2.70 \Sigma 0.81 3.26 \Sigma 0.47
(L fi =M fi ) 1.48 0.86 1.46 2.19 3.18 2.71 2.09
MGAS \Lambda =LB 0.30 \Sigma 0.03 0.50 \Sigma 0.05 0.39 \Sigma 0.03 0.39 \Sigma 0.06 0.58 \Sigma 0.09 0.27 \Sigma 0.05 0.43 \Sigma 0.03
(M fi =L fi ) 0.31 0.45 0.37 0.48 0.59 0.25 0.44
MHI =LB 0.19 \Sigma 0.02 0.35 \Sigma 0.05 0.21 \Sigma 0.03 0.18 \Sigma 0.05 0.22 \Sigma 0.04 0.14 \Sigma 0.03 0.21 \Sigma 0.02
(L fi =M fi ) 0.15 0.29 0.16 0.08 0.18 0.15 0.16
S60=S100 d 0.43 \Sigma 0.02 0.43 \Sigma 0.02 0.53 \Sigma 0.05 0.64 \Sigma 0.07 0.66 \Sigma 0.05 0.52 \Sigma 0.11 0.58 \Sigma 0.02
0.42 0.41 0.53 0.64 0.67 0.55 0.55
a Each ratio has two rows, the first is the mean value and the second is the median. The errors in the means are
the 1oe uncertainties.
b N is the number of sytems in the subsample.
c MGAS \Lambda is the total mass of the HI and H \Lambda
2
gas.
d S 60 and S 100 are the IRAX flux densities at 60 and 100 ¯m.