A&A 416, 537-553 (2004)
DOI: 10.1051/0004-6361:20034533
S. S. Larsen
European Southern Observatory (ESO), Karl-Schwarzschild-Str. 2, 85748 Garching b. München, Germany
Received 17 October 2003 / Accepted 3 December 2003
Abstract
A search for stellar clusters has been carried out in 18 nearby spiral
galaxies, using archive images from the Wide Field Planetary Camera 2 on
board the Hubble Space Telescope. All of the galaxies
have previously been imaged from the ground in UBVI. A catalogue of
structural parameters, photometry and comments based on visual inspection
of the clusters is compiled and used to investigate correlations between
cluster structure, environment, age and mass.
Least-squares fits to the data suggest correlations between both
the full-width at half-maximum (FWHM) and half-light radius (
)
of
the clusters and their masses (M) at about the
level. Although
both relations show a large scatter, the fits have substantially shallower
slopes than for a constant-density relation (size
M1/3).
However, many of the youngest clusters have extended halos which make the
determinations uncertain.
There is no evidence for galaxy-to-galaxy variations in the mean cluster
sizes. In particular, the mean sizes do not appear to depend on the host
galaxy star formation rate surface density.
Many of the youngest objects (age < 107 years) are located in strongly
crowded regions, and about 1/3-1/2 of them are double or multiple sources.
The HST images are also used to check the nature of cluster candidates
identified in a previous ground-based survey. The contamination rate in
the ground-based sample is generally less than about 20%, but some
cluster identifications remain ambiguous because of crowding even with
HST imaging, especially for the youngest objects.
Key words: galaxies: star clusters - galaxies: spiral - catalogs
In previous papers (Larsen & Richtler 1999, 2000; Larsen 1999; hereafter Papers I-III), we have studied populations of young stellar clusters in the disks of nearby spiral galaxies using ground-based imaging. That work was motivated by a desire to understand why some galaxies host young stellar clusters which are significantly more luminous (and, presumably, more massive) than open clusters in the Milky Way. Well-known examples of galaxies with rich populations of luminous young clusters include a number of merger galaxies and starbursts (see e.g. compilation in Whitmore 2003), but there are also some relatively "normal'' galaxies such as the Large Magellanic Cloud and M 33 which host a number of unusually (by Milky Way standards, at least) bright and massive young clusters (Shapley & Nail 1951; Hodge 1961; Richtler 1993; Christian & Schommer 1982,1988; Chandar et al. 1999). In Paper III we concluded that the main driving factor behind these differences seems to be the star formation rate (SFR) of the host galaxy. Galaxies with high SFRs (per unit disk area) apparently form a larger fraction of their stars in massive, bound clusters. The presence of highly luminous clusters in galaxies with high SFRs may be - at least partially - a size-of-sample effect, due to the rich cluster populations in such galaxies (Billett et al. 2002; Larsen 2002).
The study of stellar clusters is intimately linked to that of
star formation in general.
Observations show that a large fraction, if not the vast majority, of
all stars are born in clusters (e.g. Carpenter 2000;
Lada & Lada 2003). This does not, however, imply that all embedded
clusters are dynamically bound entities which survive emergence from
their native molecular cloud cores and become observable at optical
wavelengths. Lada & Lada (2003) estimated that less
than 4-7% of local embedded clusters survive to become bound
clusters of Pleiades age (108 years), but this number
may depend on environment. In Paper III we noted a steady increase with
host galaxy area-normalised star formation rate (
)
in the fraction
of U-band light originating from clusters, ranging from well below 1% in
galaxies with very low SFRs (like IC 1613) to several percent in starbursts.
Meurer et al. (1995) found that on average about 20% of the UV
light in a sample of starburst galaxies comes from young clusters.
De Grijs et al. (2003) estimated that as much as
70%
of the B-band light in the tidal tails of the "Tadpole'' and "Mice''
galaxies originates from young clusters or compact star-forming regions.
In most "normal'' star-forming galaxies, the young clusters contribute about
1% of the U-band light. Whether or not an embedded cluster remains
gravitationally bound depends on the star formation efficiency within
the proto-cluster cloud, as well as the timescale on which the gas is
expelled (Elmegreen 1983; Kroupa 2001). It may therefore
be more appropriate to
view the fraction of optically visible young stars associated with clusters
as a survival frequency, associated with the star formation
efficiency, than a cluster formation efficiency per se (which
is probably always close to 100%).
While care must be taken when interpreting the above results, due to possible differences in the age distributions of the clusters and/or field stars, completeness limits, etc., it seems clear that the Solar neighbourhood samples only a small part of the conditions under which star formation takes place in the Universe. Fortunately, there are several star-forming galaxies available within a few Mpc, spanning a range in SFRs, morphological type etc., which can be studied in considerable detail with a combination of ground-based and space-based techniques. The original sample of 21 nearby spirals analysed in Papers I-III has since been augmented by an additional handful of galaxies observed with the 3-m Shane telescope at Lick Observatory (see Larsen 2002). However, on the ground-based images, clusters were only marginally resolved, and although significant efforts were made to weed out the most obvious contaminants, the cluster lists in Paper II should only be taken as provisionary. Many of the galaxies have now been imaged with the Wide Field Planetary Camera 2 (WFPC2) on board the Hubble Space Telescope (HST) for a variety of reasons, often with multiple pointings. The WFPC2 imaging not only provides a welcome check of the true nature of the sources identified as cluster candidates from the ground, but also allows the structure and immediate environment of individual clusters to be examined in much greater detail.
Several studies have indicated a remarkable uniformity in the sizes
of stellar clusters over a wide range of masses, environments, ages and
metallicities. The most robust measure of cluster size is the
half-light or "effective'' radius (
)
which is expected to remain
relatively stable over the lifetime of a cluster (Spitzer 1987).
For Galactic globular clusters,
and luminosity are uncorrelated,
although there is a trend of increasing cluster size with galactocentric
distance (
;
van den Bergh et al. 1991). Using data
from the compilation by Harris (1996), the median
is 3.0 pc.
Similarly, the diameters and masses of Galactic open clusters show no
strong correlation, with typical sizes only slightly smaller
than those of globular clusters (Janes et al. 1988). For young
clusters in the "Antennae'' merger, Whitmore et al. (1999) found
mean effective radii of
pc. Zepf et al. (1999) estimated
half-light radii of 5-10 pc for clusters in NGC 3256, perhaps slightly
larger than for the Antennae, but again without any strong size-luminosity
correlation. For globular clusters around early-type galaxies, typical
effective radii are again 3-4 pc with no clear size-mass correlation
(e.g. Kundu & Whitmore 2001). The lack of a significant mass-size
relation is puzzling, since one might a priori expect
a cluster to form once the parent gas cloud reaches a certain density,
independent of the total mass. If this initial gas density is reflected in
the stellar density of the resulting cluster, one might naively expect the
radius to scale with mass (M) roughly as
.
However,
this is
not what has generally been reported. From the above examples, it appears
that star clusters typically have effective radii of a few (
3) pc, with
a scatter of perhaps 1-2 pc. Exceptions are found, however, including the
faint "Palomar''-type globular clusters in the outer part of the Galactic
halo, and the "faint fuzzy'' clusters recently discovered in a couple
of nearby S0-type galaxies (Larsen & Brodie 2000;
Brodie & Larsen 2002), which have larger effective radii (
10 pc).
In this paper, the cluster candidates identified from the ground are first
re-examined on archive WFPC2 images. Additional cluster candidates are then
identified on the WFPC2 images and combined with ground-based photometry
to produce a catalog of structural parameters and photometry for a sample
of clusters.
Relying on ground-based photometry limits the sample to relatively
bright objects, but has the advantage of providing uniform photometric
coverage of all clusters (even if crowding effects are more severe than
in the HST data). In particular, most of the HST datasets do not
include imaging in a U-band equivalent filter, which is essential for
age-dating the clusters. However, the HST photometry may still be
useful for some purposes and aperture photometry in an
aperture is presented for the available bandpasses in a separate table
for each cluster candidate.
Each entry in the catalogue also contains
various notes on the degree of crowding, close neighbours etc. based on a
visual inspection of the candidates.
Because the HST images cover limited sections of the galaxies and span a huge range in exposure times and filters, the cluster sample presented here still cannot be considered complete in any sense. The completeness is a complicated function of crowding, cluster size, underlying surface brightness, exposure time in the HST images, bandpass, seeing in the ground-based data, galaxy distance, and probably many other factors which would be next to impossible to model in a satisfactory way. It should also be emphasized that what is presented is still a list of cluster candidates, which might contain contaminants (e.g. background galaxies). A definitive list of bona-fide clusters would require spectroscopic follow-up, but such an effort is beyond the scope of this paper and left for future studies. With this in mind, it is hoped that the catalogue may still provide a useful basis for further studies. As an example, it is used in Sect. 4 to investigate trends with age and mass in the cluster sizes, degree of crowding, and shape parameters.
The search for HST archive data was concluded in October 2002 and only includes WFPC2 datasets which had been publicly released up until that time (Table 1). ACS data were available for a few galaxies, but have been excluded in order to allow a relatively simple and homogeneous reduction procedure.
Table 1:
Exposures.
indicates the offsets
in right ascension and declination between the coordinate systems of
the WFPC2 frames and the ground-based data (tied to the USNO catalogue).
The offsets are given in s/15 for right ascension and in
arcseconds for declination (i.e. the
factor has not
been applied to convert the offsets in
to true arcsecs).
Given the large volume of data, a fairly high degree of automatization had
to be incorporated in the reduction procedures. When several exposures
were available for a given field and filter, these were combined with
the CRREJ task in the STSDAS package in IRAF. In most cases, no
shifts in the image coordinate systems were required before combination
but when necessary, such shifts were applied using the IMSHIFT command
in IRAF. For each WFPC2 pointing and each filter, objects were then
detected with the DAOFIND task in DAOPHOT (Stetson 1987) running
within IRAF. Objects
with saturated pixels within a radius of 5 pixels were rejected. The
detected sources were fitted with
the ISHAPE profile-fitting algorithm (Paper II). ISHAPE models
each source assuming an analytic model for the intrinsic profile of
the source, convolved with the HST/WFPC2 point spread function (PSF).
The FWHM of the analytic model is iteratively adjusted until the best
fit to the data is obtained. The initial round of profile fitting was
carried out with a fitting radius of r=5 pixels and assuming
a model of the form
Because the main aim was to study the structure of spatially
resolved objects, only objects which were detected at >10
above
the background noise were included in the analysis. Fainter objects would
generally have too low S/N for reliable size measurements, and would have
increased the already substantial computing time required to fit the spatial
profiles. Because of the vastly different exposure times, different
bandpasses, background levels etc., the 10-
detection threshold
does not translate into a well-defined completeness limit in terms of
magnitude. A total of 82 000 sources were detected and fitted, requiring
a few days of CPU time on a 1.5 GHz Pentium PC. At this stage, many
objects appeared several times in the source list, being included in
multiple HST pointings and/or bands.
After the initial round of profile fitting, the HST object lists were
matched with the photometry data files from the ground-based surveys (all
details concerning the reduction of the ground-based data are given in Paper
II). Not only cluster candidates previously identified as such in the
ground-based surveys, but all point-like sources in each ground-based
CCD frame for which photometry was available, were matched. The matching was
done by converting pixel coordinates measured on the WFPC2 images to
celestial coordinates, using the METRIC task in the STSDAS package in IRAF.
These coordinates were then matched with the coordinates of objects measured
on the ground-based CCD images, tied to the US Naval Observatory meridian
circle catalogue as provided by the ESO SkyCat Tool (Monet et al.
1998). In many cases, there were clear systematic offsets between
the WFPC2 and USNO coordinate systems. These offsets were determined by
displaying each individual WFPC2 frame and then marking the ground-selected
cluster candidates contained within that frame. Because these would
typically be among the brightest objects in the WFPC2 frames, and the
offsets in general relatively minor (1
), identifying the
cluster candidates in the WFPC2 frames was usually unproblematic. The offsets
between the WFPC2 and USNO coordinate systems are listed in
Table 1 for each dataset. Note that large offsets
were found for the exposures belonging to snapshot programme 5446.
This is probably because these datasets were guided using the "Gyro Hold''
mode, which provides less accurate pointing and tracking than the
Fine Guidance Sensors on HST. A few WFPC2 pointings had no ground-based
cluster candidates and were excluded from further analysis. For those
pointings, no offsets are given in Table 1.
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Figure 1:
Histogram of FWHM values for objects fitted during the first
round of profile fitting. The narrow sequence of unresolved objects
with FWHM ![]() |
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Figure 2:
Illustrations of comment codes a-i. Please see
Table 2 for the full comments associated with
each code. Each panel shows a
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Figure 3: Examples of objects classified as Type 2 (uncertain). |
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Figure 4: Two examples of how the appearance of cluster candidates can change dramatically depending on bandpass. |
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The initial round of profile fitting served to separate point-like sources
(likely stars) from extended ones (cluster candidates). Figure 1
shows the distribution of intrinsic FWHM values (in pixels). There is
a very narrow peak of objects with FWHM values close to 0, although there
is no sharp division between resolved and unresolved objects. For further
analysis, a total of 3100 sources with S/N>50 (measured within the 5
pixels aperture on the WFPC2 images), FWHM >0.2 pixels (dashed line in
Fig. 1) and ground-based photometry were selected. For
comparison, the (undersampled) PSF of WFPC2 itself has a FWHM of about 1.5
pixels. At a typical distance of 5 Mpc, the size cut corresponds to a FWHM
of 0.5 pc or a core radius of about 0.25 pc. Many of these sources still
represented multiple observations of the same object.
Objects satisfying these criteria were again
fitted, but this time allowing the power-law index ()
of the EFF
profiles to vary as a free parameter. In order to better constrain
,
a fitting radius of 10 instead of 5 pixels was
chosen for this second round. The choice of a 10 pixels fitting radius
represents a compromise between reasonably accurate constraints on the fit
parameters, reducing the effects of a non-uniform background, and
computing time.
Each source was fitted three times, using
different initial guesses for
(
,
1.5 and 2.0).
The values 1.0 and 1.5 roughly bracket those typically found for real star
clusters (e.g. Elson et al. 1987 and Sect. 4.3),
while an additional more extreme value of
was included in
order to reduce any a priori bias in the measurements, in
case some clusters have steeper slopes (although tests of the
ISHAPE algorithm suggest that any dependence of the fitted parameters
on the input guesses is relatively minor; see Sect. 2.2
below).
Finally, all cluster candidates fitted in the second round
were visually inspected. In addition, all cluster candidates
identified in the ground-based surveys and with HST imaging were
inspected regardless of the S/N in the HST images. The inspection was
done by displaying a
section of all images of
each cluster candidate in an IDL widget, which also
contained a number of check-boxes corresponding to the comment
codes given in Table 2.
Table 2: Explanation of comment codes.
Each object was also assigned to one of three types: 1) likely/certain cluster, 2) uncertain classification, 3) likely/certain non-cluster. A flag was set if the ISHAPE fit was unlikely to have resulted in meaningful structural parameters, as judged from the degree of crowding etc. The inspection process was repeated three times, each time displaying the cluster candidates in a different (random) order. At the end, the final object type and ISHAPE "Fit-OK'' flag were determined as the most pessimistic of the three individual estimates. A comment code was included in the final list if selected in at least 2 out of the 3 inspection rounds.
Clearly, by its very nature any visual inspection involves an element
of subjectivity. Very often, more than one comment applied to
a given object. In such cases, all applicable comments were selected and
included in the catalogue. Furthermore, resolution (and thus distance)
effects may change the visual appearance of certain objects and
thereby e.g. cause an object classified as "double'' in a nearby galaxy
to be labelled "elongated'' in a slightly more distant one. In order
to illustrate roughly what the individual comment codes in
Table 2 represent, Fig. 2 shows
sections around some clusters that are
typical for each comment code. Figure 3 shows a few
examples of objects classified as "uncertain''.
Table 3: All clusters with S/N>50 on HST images. Photometry is from ground-based data while morphological data are from HST/WFPC2 images. The first 10 rows of the table are reproduced here; the full table (1358 rows) is only available in electronic form at the CDS.
Another difficulty is that objects can appear quite different in different bandpasses. As an example, Fig. 4 shows two objects observed in F336W, F439W, F555W and F814W. The F336W and F814W images of n3621-620 are hardly recognisable as the same object - the F336W image just shows what might be a loose association of stars, while the F814W image shows two fairly well-defined, compact sources. Conversely, n3621-513 looks compact and symmetric in F336W, while two nearby neighbours (perhaps red giants or supergiant stars) appear in F814W. If multi-wavelength data were available for all objects, one might adopt a consistent strategy for dealing with this problem, but in many cases data were only available in one band. It should therefore be borne in mind that the visual comment codes (and even the profile fits) depend on the bandpass in a way that is not easily predictable.
The final list of cluster candidates includes 1358 objects
(Table 3).
Columns (1)-(8) contain information from the
ground-based data: coordinates, UBVI photometry and projected
galactocentric distance (in arcmin). The photometry has been corrected
for Galactic foreground extinction using the Schlegel et al.
(1998) values and the extinction law in Cardelli et al.
(1989). Note that the reddening corrections differ from those in
Papers I-III, where Burstein & Heiles (1984) values were used.
In Papers I-III we used a relatively large aperture radius of 8 pixels
(
and
for the NOT and Danish 1.54 m data, respectively)
to avoid possible systematic effects in the integrated magnitudes due to the
extendedness of the objects. However, it is now clear that most clusters are
compact enough that this is not a major source of concern in ground-based
imaging, and in the present paper I therefore use a smaller aperture radius
of 4 pixels, the same as for the colours, for the ground-based magnitudes.
Information derived from HST images is listed in Cols. (9)-(16): the
number of individual detections of each cluster
(where one "detection'' is defined as the presence of the cluster in
an image taken through a given filter under a given programme),
N, is in Col. (9),
followed by the full-width at half maximum (FWHM) of the cluster profile
derived from the ISHAPE fits, the exponent
,
the effective
(half-light) radius
,
the minor/major axis ratio (x/y), the object
type (Col. 14), Fit-OK flag (Col. 15) and comments.
Unlike the classical King profiles (King 1962,1966), the
EFF models have no finite radius, and for
the volume
contained under the profile is infinite. For
only slightly larger
than unity, the total volume converges very slowly, resulting in
unrealistically large
values. To cope with these difficulties, the
values in Table 3 are computed for a finite outer radius
of 50 pc, beyond which the luminosity profiles of young clusters become
difficult to trace even in nearby galaxies such as the LMC
(e.g. Elson et al. 1987). However, it is important to note
that the estimates of
are generally based on an extrapolation of the
luminosity profiles beyond the fitting radius, and carry significant
uncertainties especially when
.
Instead of listing the FWHM, I
could have given the core radius
,
since both are always defined. The
reason for listing FWHM is that there is some ambiguity in the definition
of the core radius. Some authors define it as FWHM/2, but it may also be
defined e.g. as the scale radius
in Eq. (1).
In order to avoid confusion, I will simply use the term FWHM rather than
"core radius'' throughout the remainder of this paper. When more than
one exposure was available for a cluster candidate, the shape parameters
in Table 3 were obtained by weighting the measurements on
each exposure by its S/N.
As discussed above, the morphology of clusters can be quite
wavelength-dependent, but cases where the determination of shape
parameters is particularly uncertain can generally be recognized by
the "Fit-OK'' flag in Col. (15) of Table 3.
Errors were estimated as follows: for each exposure, the error on the shape parameters derived from that exposure were estimated as the standard deviation of the three individual fits. If only one exposure was available, this is the error listed in Table 3. When several exposures were available, the errors in Table 3 are the estimated standard errors on the mean of the weighted average. Some additional comments for a few objects, which did not fit into the codes in Table 2, are listed in Table 4.
HST photometry for each cluster candidate is given in Table 5.
Table 4: Additional comments (only 5 sample entries given).
Table 5:
HST photometry for the same clusters listed in Table 3.
All magnitudes are in the STMAG system, measured in a
aperture and applying an aperture correction of -0.2 mag. No
corrections for foreground reddening have been applied.
Only the first 10 rows are reproduced here.
In previous papers it has been documented that no systematic differences
seem to be present between ground-based and HST-based colours for the
cluster candidates (e.g. Larsen 2002), although a random scatter of
0.1-0.2 mag exists. For integrated magnitudes, on the other hand, an
offset of a few
mag has been found between ground-based and HST
magnitudes, in the sense that ground-based data tends to give brighter
magnitudes. With the larger sample of clusters available here, this
comparison can now be carried out in more detail. Of the clusters listed
in Table 5, 1245 have data in at least one of the
filters F547M, F555W and F606W, all of which are reasonable approximations
to the Johnson V-band. The mean difference
between ground-based
and WFPC2 photometry, including all objects with F547M, F555W or F606WHST data is -0.50 mag with a large scatter of
mag.
This scatter is partly due to the fact that the ground-based
photometry of some of the fainter clusters has large errors, but
decreases only slightly (to
mag) if
clusters with formal errors larger than 0.2 mag on the ground-based Vmagnitudes are excluded. Thus, most of the errors are clearly of a
systematic nature.
If the samples observed with the Danish 1.54 m telescope and the NOT are
compared with the HST photometry separately, interesting differences emerge.
The image quality of the data taken with the two telescopes differ
significantly, with typical FWHM seeing values of
and
,
respectively (Larsen 1999). For the galaxies observed with the
Danish 1.54 m, the mean difference between ground-based and HST
photometry is
mag (where the 0.61 mag refer to
the standard deviation around the mean, not the error on the mean value).
For the NOT sample, the corresponding numbers are
mag.
Thus, while the scatter remains large even in the NOT data, the systematic
difference relative to the HST photometry is clearly smaller than for the
Danish 1.54 m data.
A few individual, relatively isolated clusters observed with both the NOT
and Danish 1.54 m telescope have been analysed in detail
(Larsen et al. 2001; Larsen & Richtler 2004, in preparation) and for
these clusters there is good agreement between ground-based and HST
magnitudes (within
0.1 mag). The differences between the
mean magnitudes of the HST and ground-based samples are therefore not due
to trivial zero-point errors in the photometric calibrations.
Clusters for which any of the comment flags in Table 2
are set might be expected to show poorer agreement with the ground-based
data. Indeed, if such clusters are rejected then
mag
for all clusters, and
mag and
mag for the Danish 1.54 m and NOT samples, respectively. Thus, the systematic
difference between HST and ground-based photometry clearly decreases,
albeit still with significant scatter. The remaining offsets can probably be
attributed to contamination within the ground-based apertures which did
not trigger any comment flags. In fact, if the HST photometry is instead
carried out using an
aperture radius, assuming that such an
aperture encircles all flux from the objects (i.e. applying no
aperture corrections) then the mean offset with respect to the NOT data
is only
mag. Excluding clusters with comment codes,
the difference decreases even further to
mag.
A smaller subset of the clusters have observations in HST bandpasses
that allow a comparison with the ground-based colours. For example,
Holtzman et al. (1995) give transformations to Johnson
colours for the F439W and F555W bandpasses, which are available for
190 of the clusters in Table 5. The mean offset
between ground-based and HST
colours is
mag
with a scatter of 0.17 mag. For the Danish 1.54 m and NOT samples, the
differences are
mag and
mag,
respectively. This confirms that the ground-based colours are more accurate
than magnitudes, presumably because the objects that contaminate the
ground-based apertures tend to have similar ages and colours to the
clusters themselves.
Tests of the ISHAPE code have been carried out in several previous
papers, in particular in Paper II. The reader is
referred to that paper and to the documentation included with the
code for further details. However, because the ability to fit the
EFF
parameter is a more recent addition, some tests of
this particular feature are presented in the following.
First, artificially generated sources with known profiles were added
to images of NGC 7793 and NGC 5194. The frames used for these tests
were the WF4 chips from proposal 8591 (NGC 7793, F547M, 4
400 s) and
proposal 7375 (NGC 5194, F555W, 2
600 s). The artificial sources
were generated by convolving EFF profiles with the HST PSF, generated
by TinyTim, and then added to the science images with the MKSYNTH program
(Paper II). For each galaxy, artificial sources with FWHM of 1 and 2 pc,
= 1.0 and 1.5, and magnitudes of V=20, 21 and 22 were added. For
each combination of these input parameters, 25 objects were added at random
positions and then fitted with ISHAPE using three different values
(
= 1.0, 1.5 and 2.0) for the initial guess of
.
The test results are summarised in Table 6.
Table 6:
Tests of the ISHAPE profile-fitting algorithm. For each
combination of input parameters (
V=20/21/22, FWHM =1 pc/2 pc)
the output fitted FWHM and
parameters are
shown for three initial guesses of
(
).
Numbers in parentheses denote the object-to-object rms deviation
around the mean values, excluding the two most deviating points at
each extreme.
Neither the FWHM nor the
values returned by ISHAPE show any
systematic dependence on
,
and the mean fitted
FWHM and
are generally quite close to the input values. For
,
the rms scatter of the FWHM and
measurements are typically
0.15-0.20
pixels and 0.10-0.15 (dimensionless units), respectively. The mean standard
deviation (
)
of the three individual fits with different
is much smaller than the object-to-object rms scatter,
indicating that
the uncertainty due to a particular initial guess is smaller
than the random measurement error. For typical cluster-like objects
with
,
these tests suggest that the FWHM and
parameter can be measured with an accuracy
20% and
10-15%, respectively on a single image.
Note, however, that the error on the effective radius
derived from the FWHM and
can be very much larger, especially if
.
It is therefore important that any error selection be
carried out on FWHM and
,
and not on the derived quantity
.
As a final caution, the tests carried out here do not take into account any uncertainty on the PSF itself. In practice, this can be important even for HST images, since the so-called "diffusion'' kernel makes an important contribution to the scattering of light between neighbouring pixels. While this effect is included in the modelling done by ISHAPE, the diffusion kernel has so far only been properly characterised for the F555W filter. It is probably appropriate also for the F547M and F606W filters, but for bluer or redder passbands the true diffusion kernel could be significantly different.
One of the motivations for this study was to test the reliability of the ground-based cluster identifications and quantify how much any mis-identifications might affect the specific luminosities for the cluster systems derived in Paper III. To this aim, Table 7 lists photometry and object classifications for all cluster candidates originally detected in the ground-based surveys, which are also included in the HST datasets.
Table 7:
Clusters from ground-based survey reidentified in HST images.
Photometry is from ground-based data. FWHM is the
full-width-at-half-maximum in pc derived from EFF
fits. The
first 10 rows of the table are reproduced here; the full table (330 rows)
is only available in electronic form at the CDS.
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Figure 5: Size distribution for cluster candidates identified in ground-based surveys. |
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Figure 6:
Magnitude distributions for ground-selected cluster
candidates with FWHM <0.5 pc ( top) and FWHM >0.5 pc ( bottom).
Sizes are from EFF fits with fixed
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Figure 7: Examples of objects identified as clusters from the ground, but which are unresolved according to the ISHAPE fits. |
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The distribution of FWHM values for objects listed in Table 7
is shown in Fig. 5. Only 6 objects were classified as
"likely/certain non-clusters'' during the visual inspection, but
69 objects (or about 21% of the sample) have FWHM <0.5 pc and are
thus essentially unresolved even on the HST images.
Unresolved objects are not necessarily individual, isolated stars, but
can also be loose groupings of stars (OB associations etc.) where
ISHAPE simply picks one star and fits it.
This is the case for about 25 out of the 69 unresolved objects,
or about 1/3. A few examples are shown in Fig. 7.
The size distribution in Fig. 5 does not appear
strongly peaked at a particular value (except for the unresolved
sources near FWHM =0), but spans a range from the resolution limit
up to 8-10 pc. If objects with FWHM <0.5 pixels are excluded, the
formal estimate of the mean FWHM is 3.8 pc.
For an EFF model with
,
this corresponds to a half-light radius of 4.3 pc (or
4.0 pc if the profile is truncated at 50 pc), consistent
with the typical sizes of young star clusters in the Milky Way and elsewhere.
The mean FWHM changes only slightly (to 3.9 pc) if objects with
"Fit OK? = N'' are excluded. The issue of cluster sizes will be discussed
in more detail below for the full sample.
Figure 6 shows the distributions of absolute MV
magnitudes for unresolved and resolved objects. While the magnitude
distribution for resolved objects (bottom panel) does extend to brighter
magnitudes than for unresolved ones, there are several unresolved
objects brighter than MV=-10. Such bright objects are unlikely to be
individual stars, but most of them have the "c'' comment set (companions
within
), indicating that the ground-based magnitudes are
likely contaminated by
nearby objects. Another possibility is that some of them are very compact
star clusters. At magnitudes fainter than MV=-9,
Fig. 6 shows a
clear excess of unresolved objects
(26 out of 66, or 39%, compared to 21% for the whole sample),
many of which may indeed be individual stars.
Another way of estimating the contamination fraction is to use the object types from the visual inspection. Of the 330 objects, 61 are labelled as type "2'' (uncertain) or "3'' (very likely non-cluster). Although classification as type "2'' does not necessarily mean that the object in question is a non-cluster, these numbers again suggest a contamination rate of order 20% or less.
![]() |
Figure 8:
Specific U-band luminosity (TL(U)) vs. area-normalised
star formation rate (
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How does contamination of the cluster samples affect the
relations between specific luminosity (TL(U)) and host galaxy
properties derived in Paper III? If the contamination rate were the
same in different galaxies, all specific luminosities would just be
scaled down by a constant factor, leaving any relations involving
TL(U) intact. However, it is possible that the ratio of clusters
to potential contaminants varies from galaxy to galaxy.
Table 8
lists the total number of cluster candidates identified from the ground
(
)
in each galaxy, followed by the number of ground-identified
candidates contained within the WFPC2 images covering that galaxy
(
), and the number of objects among these that are unresolved
(
).
Table 8:
Contamination of ground-based sample for individual galaxies.
is the total number of cluster candidates identified in the
ground-based surveys,
is the subset of those candidates covered
by HST images and
is the number of unresolved
sources. The adopted distance moduli are also listed (see Paper I for
references).
An updated version of the TL(U) vs.
plot
from Paper III is shown in Fig. 8. The original data
are shown with error bars, while new updated TL(U) values,
corrected according to Table 8, are shown with filled
diamonds. The basic trend for TL(U) to increase as a function
of
is clearly preserved.
The scatter increases somewhat after the exclusion of unresolved objects.
However, with only a couple of clusters left in galaxies
like NGC 2835 and NGC 5204 after exclusion of unresolved objects,
TL(U) as defined in Paper III is probably no longer a good approximation
to the true relative luminosity of the cluster system. To obtain a more
useful number, one would likely have to probe to fainter magnitudes and
thereby sample the cluster population more completely.
The broad-band colours of simple stellar populations (such as star
clusters) are functions of both age and metallicity, with additional
complications arising from unknown reddenings and stochastic
effects due to the finite number of stars in a cluster (e.g. Girardi et al. 1995). However, for
clusters younger than 109 years it is still possible
to obtain reasonably accurate photometric age estimates, especially if
U-band data are included. For such clusters, metallicity effects
are weak, except for a brief period around 107 years when the
integrated light is dominated by red supergiants.
Here, cluster ages were obtained by fitting Bruzual & Charlot
(2000; priv. comm.) model
colours to the observed UBVI cluster colours. The SSP model fits were done
by minimizing the rms deviation between model- and observed colours (weighted by their errors) as a function of age and reddening. In order to reduce the
uncertainty on the age determinations, age estimates were made only for
clusters with
mag,
mag and
mag.
The ages of individual clusters derived from broad-band colours should
only be regarded as approximate. The ground-based apertures may be
contaminated by objects other than the cluster candidate itself, and model
uncertainties also make the absolute ages uncertain. However, the relative age
ranking of clusters should still be reasonably reliable. Line emission can
also affect the broad-band colours of very young objects, and must be taken into
account if accurate age estimates for objects younger than 107years are required (e.g. Anders & Fritze-v. Alvensleben 2003).
Table 9 lists the mean FWHM and
for clusters in
each galaxy.
Table 9: Mean FWHM and effective radii for clusters in galaxies.
Only clusters classified as "Type 1'' and with "Fit OK = YES'' were included. The number of clusters in each galaxy satisfying these criteria is given in the second column. Because of the poorly defined
The mean effective radius is
pc,
perhaps slightly larger compared to those for the ground-selected sample
(Sect. 3) and other young and old star clusters. This may be due
to the fact that the clusters are assumed to follow a single power-law out to
a total radius of 50 pc, while in reality the behaviour at large radii is
poorly constrained. If the cluster profiles decline more rapidly at large
radii the effective radii would decrease, especially for objects with
-values close to 1. Also, the size cut imposed in order to
exclude point sources may introduce a bias against the most compact clusters.
In the last column of Table 9, the
requirement is abandoned. Clearly, this leads to an increase in the
mean effective radii, but it is stressed that the
values
for these objects are very uncertain and depend strongly on the
adopted outer radii.
![]() |
Figure 9: Mean FWHM ( top) and effective radii ( bottom) for stellar clusters versus parent galaxy distance modulus. |
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Figure 10:
Mean FWHM ( top) and effective radii ( bottom) for
stellar clusters versus parent galaxy
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The mean FWHM and
from Table 9 are plotted versus
parent galaxy distance modulus in Fig. 9.
The sizes do show some correlation with parent galaxy distance,
possibly due to a less than perfect correction for the PSF. It
is also possible that a larger fraction of the objects detected
in more nearby galaxies are individual stars, rather than clusters, which
made it into the list of cluster candidates despite the size cut.
Contributing to this effect, the number of individual stars bright enough
to be detected in the ground-based photometry would increase at
smaller distances.
Furthermore, the fixed size cut at FWHM =0.2 pixels corresponds to
a different physical cluster size in different galaxies, ranging from 0.3 pc
(core radius
pc) in NGC 7793 to 1.2 pc (core radius
pc) in
NGC 7424.
Figure 10 shows the mean cluster sizes as a function of
the area-normalised host galaxy star formation rate,
(see Paper III).
The scatter in the FWHM plot is somewhat larger than for the
values,
but neither
FWHM nor
shows any obvious correlation with
.
Below, the
data for all galaxies is combined and analysed collectively in order to
improve statistics, but in order to reduce possible systematic effects
due to different distances, clusters in the closest (NGC 7793) and two most
distant galaxies (NGC 2997 and NGC 7424) are excluded for the analysis of
structural parameters. Thus, in summary the selection parameters applied to
the cluster candidates in Table 3 for the following analysis are:
![]() |
Figure 11:
Distributions of FWHM and effective radii for the combined
cluster samples. In the lower panel, the hatched
and outlined histograms are for clusters with
![]() ![]() |
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Figure 12: Distribution of envelope slopes in four age bins. |
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Figure 12 shows the distribution of
values for
clusters in four different age bins. For reference, typical errors of 0.15
on
are indicated (cf. Sect. 2.2). The distribution
of
values is peaked around 1.2-1.3 in all four bins, but there may
be a trend of mean
increasing with age,
at least in the sense that the youngest age bin appears to contain a
higher fraction of clusters with
.
One important caveat in this comparison is that the mean mass is likely to
increase from the youngest to the older age bins, due to the increase in
M/L ratio with age.
Thus, in principle the difference between the
-distributions in
Fig. 12 might be due to the different mass ranges sampled
in each bin, rather than being an evolutionary effect.
To test whether this might be the case, Fig. 13
displays the
-distributions in two age intervals,
but now also divided into different mass intervals. The number of
clusters in each panel is quite small, but the excess of clusters with
seems to be present for log(age) <7.0 in all three mass bins.
This suggests that the difference between the
-distributions in
Fig. 12 is not just an effect of the different mass
intervals covered at different ages.
![]() |
Figure 13: Envelope slopes in different mass bins for two different ages. |
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Other authors have previously noted that very young clusters tend to be surrounded by relatively extended envelopes with more light at large radii than a King profile. Elson et al. (1987) estimated that as much as 50% of the mass in young LMC clusters may reside in unbound halos. In the Milky Way, a large fraction of the youngest open clusters also have very large radii and may be unbound and in the process of dispersing away (Janes et al. 1988). Whitmore et al. (1999) showed radial profiles for three clusters in the Antennae, illustrating a gradual transition from extended envelopes with no well-defined outer radius (for the highly luminous "Knot S'', only a few Myr old) to older clusters where a tidal cut-off becomes apparent. A similar extended halo was observed for a very luminous, 15 Myr old cluster in NGC 6946 (Larsen et al. 2001). The structure of very young clusters may hold important clues to the structure of the progenitor clouds out of which the clusters formed, although it may prove challenging to disentangle this from the effects of early dynamical evolution.
Figures 14 and 15 show FWHM
and
versus age and mass for the combined cluster samples. Clusters
older than 1 Gyr are excluded from
the plots because of large uncertainties on the ages (and therefore also on
the masses derived from integrated photometry). The masses were estimated
using M/L ratios from the Bruzual & Charlot models, assuming a Salpeter IMF
from 0.1
-100
.
Absolute masses are sensitive to the
shape and lower-mass cut-off of the IMF,
and would be
30% lower for a Kroupa (2002) IMF,
but this is not important for
the relative comparison attempted here.
In Fig. 15, clusters with
and
are shown with different symbols (diamonds and plus markers).
The dashed lines superimposed on the lower panels (size vs. mass) of each
figure represent the size
relation corresponding to
constant cluster density.
Typical error bars are also shown in the top right corner in each of
the lower plots.
![]() |
Figure 14:
Cluster FWHM versus age ( top) and mass ( bottom). The dashed curve
in the lower panel represents a relation corresponding to FWHM
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![]() |
Figure 15:
Effective radius (
![]() ![]() ![]() ![]() ![]() ![]() |
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Neither Fig. 14 nor Fig. 15
shows any strong evidence for a general correlation between
cluster size and age. Since FWHM is roughly equivalent to
,
it is interesting to note that Mackey & Gilmore
(2003) found a strong correlation between core radius and age for
clusters in the LMC (confirming earlier results by
Elson et al. 1989). For ages < 108 years, there are essentially
no clusters in the LMC with
pc, while clusters with ages
109 years show a full range of core radii from less than 1 to 8 pc.
While it
cannot be ruled out that the most extended clusters are missing from the
sample, Fig. 14 shows no strong
-age relation
similar to that in the LMC.
Likewise, the size-mass plots show no strong correlations. The solid
line in the bottom panel of Fig. 14 is a least-squares
power-law fit of the form
![]() |
Figure 16:
Histograms of the
![]() ![]() ![]() ![]() |
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Using the HST photometry in F547M, F555W or F606W in Table 5
instead of the ground-based data to estimate the V-band luminosities
of the clusters has no significant effect on the FWHM vs. M and
vs. M relations. For
vs. M, the slope B decreases by only 0.003,
while for FWHM vs. M it decreases by 0.014, in both cases well within
the formal uncertainties on the fits. Note, however, that the ages still
have to be estimated from the ground-based photometry.
The high degree of crowding in Fig. 15 makes it
difficult to visually assess to what extent the two curves agree with the
data. A somewhat different representation of the same data is shown in
Fig. 16, where the
distributions are shown for
five separate mass bins. The sample with
is shown with
shaded histograms, while the sample including clusters with
is shown with outlined histograms. In each panel, the range of sizes
corresponding to an
scaling is also indicated
(using the same arbitrary normalisation as in Fig. 15,
not taking any scatter into account). The paucity of objects in the
0-1 pc bin is due to the size cut, but even if one accounts for the
fact that the most compact clusters may have been systematically excluded, it
seems difficult to reconcile the overall
distributions in
Fig. 16 with a constant-density relation.
Although completeness effects are difficult to quantify, they would most likely tend to work against the detection of faint, extended clusters, thereby strengthening rather than weakening any existing size-mass trend. Thus, while the size-mass trends suggested by Figs. 14 and 15 remain quantitatively uncertain, the current data seem to imply that any size-mass trend, if it exists, is substantially shallower than for a constant-density relation.
While more subjective than the measurements of FWHM,
and
,
the comment flags in Table 3 hold useful information
about the surroundings of each cluster.
It is reiterated that the comment flags are based on visual inspection
of a fairly heterogeneous dataset, and that the morphology of cluster
candidates can be quite wavelength-dependent. Note, however, that the
comment codes are based on visual inspection of all available images of
each cluster, and the majority of the clusters have imaging in a
B-, V- or I-band equivalent filter.
Table 10 lists the fraction of cluster candidates with
comment flags a/c/d (all of which are likely indicators of multiplicity) and
e/f/g/h (more general crowding indicators).
Table 10: Statistics on comment flags in different age intervals.
These statistics are given both for the best cluster candidates with structural parameters determinations (i.e. Fit-OK = YES and Type 1), as well as for all potential cluster candidates including those without reliable fits (Fit-OK = YES/NO and Type 1/2). Data for all galaxies are included in Table 10, but the numbers remain unchanged within 1-2% even if NGC 2997, NGC 7793 and NGC 7424 are omitted as in the previous sections.Very few clusters with Fit-OK = YES have any of the /e/f/g/h flags set. This is no coincidence, because these flags indicate exactly those conditions which would make profile fits uncertain. The table indicates a strong evolution in the environment as a function of age. Most of the cluster candidates in crowded environments (flags e/f/g/h) are younger than 107 years. The tendency for the crowding to decrease with age is probably a consequence of fading and dispersion of the surrounding stellar population. Assuming typical velocity dispersions of a few km s-1 within (unbound) star forming regions, the expansion will amount to a few tens of pc in 107 years.
Of the objects in the youngest age bin, 42% were classified as "uncertain'' (type 2). This underscores the fundamental difficulty of identifying the youngest clusters. The problem of defining an appropriate selection criterion for bona-fide clusters is far from trivial. Sometimes the main problem is simply that an object is only barely resolved. In such cases, better angular resolution would help confirm or rule out the cluster nature. For objects with complex morphology it can be difficult to determine whether or not an object is a true star cluster, even if resolution would otherwise not be a problem. Examples can be seen in Fig. 3 and in panels (e), (f) and (g) of Fig. 2. In these cases, it is difficult to determine whether a well-defined stellar cluster is present. For low-mass clusters of low age, an additional problem is that the integrated light can be dominated by a few luminous stars, making it difficult to distinguish such objects from random superpositions of individual stars along the line-of-sight.
As pointed out in the introduction, most stars probably form in clusters, but only a small fraction of young embedded clusters survive as bound objects. It may also happen that only a fraction of the stars in a cluster remain bound, while the rest disperse away (Kroupa 2001). Thus, a few Myr-old concentration of stars may be a bound star cluster, a bound core surrounded by an expanding envelope, or an entirely unbound association which will soon disperse away completely. Other star formation may also be taking place nearby, perhaps triggered by the young cluster. So it is not surprising that a large fraction of the youngest objects have a messy morphology.
The age distribution of double or multiple objects does not appear to be as strongly peaked at young ages, with some objects flagged "a/c/d'' even in the oldest age bin. It is possible that at least a fraction of these objects are genuine double clusters, similar to those found in the Large Magellanic Cloud. The LMC binary clusters tend to be predominantly young, though a few pairs as old as several times 108 years exist (Dieball et al. 2002). This seems to be consistent with the decreasing fraction of multiple objects at high ages in Table 10. However, because the main source of photometry in this paper is ground-based imaging, no information is available about possible colour/age differences between the components in such pairs. Multi-colour HST imaging, especially including U-band data, would allow a more thorough investigation of double clusters and make a comparison with the LMC sample possible.
Using a combination of HST/WFPC2 imaging and ground-based UBVI photometry,
a catalogue of cluster candidates in 18 nearby spiral galaxies has been
compiled. Only objects with a S/N>50 on the HST images (within an r=5
pixels aperture) were included, allowing for a detailed analysis of the
structure of individual clusters. Analytic profile fits of the form
were carried out, including a proper
modelling of the HST/WFPC2 PSF, and allowing both the core radius
and
envelope slope parameter
to vary. Each cluster candidate has
also been visually inspected and comment flags relating to crowding and
multiplicity are given. These comment flags, combined with
the photometric data and structural parameters, may be helpful when using
the list of cluster candidates to select targets e.g. for spectroscopic
studies.
The HST imaging indicates a mean contamination rate of 20% or less
for the ground-based cluster surveys in Papers I-III,
although the contamination rates in some individual galaxies (most notably
NGC 5204) are higher. However, the relation between specific U-band
luminosity of the cluster systems TL(U) and
(Paper III) remains
valid after correction for contamination. Because very few clusters
remain after the contamination correction in some galaxies, sampling the
cluster populations to fainter magnitudes than the limits defined in
Papers I-III would probably reduce the scatter in the
TL(U) vs.
relation.
The cluster catalogue has been used to investigate trends and relations
between various cluster properties, although the analysis is complicated
by the fact that selection effects are difficult or impossible to control
with a dataset as heterogeneous as the one used here.
Most clusters have FWHM less than about 8 pc with a formal mean
of about 2.7 pc, corresponding to a mean core radius
1.3 pc, but very compact clusters may be
missing because of the size cut imposed in order to exclude point sources
(individual stars).
For the subset of the clusters which have
and therefore
reasonably reliable measurements of the effective radius
,
the mean
value is
pc, but this
mean value may again be affected by a selection bias against the most
compact clusters. The effective radii are also
sensitive to the poorly constrained behaviour of the luminosity
profiles at large radii. Here, the profiles are truncated at 50 pc.
The structural parameters show little or no variation from
galaxy to galaxy, especially when considering that the distances are not
always known very accurately. In particular, the effective radii are
uncorrelated with the host galaxy area-normalised star formation rate, and
are also very similar to those of open and globular clusters in the
Milky Way, globular clusters in early-type galaxies, and young clusters
in merger galaxies and starbursts. It is quite remarkable that the
sizes of stellar clusters are largely invariant with
respect to the properties of the parent galaxy. Physical parameters such
as gas density and -pressure probably play a major role in regulating the
star formation rate (Kennicutt 1998), but while these factors
may affect the formation efficiency of bound clusters (Paper III) they do
not appear to have a strong impact on the structure of the clusters
themselves, once formed. Exceptions to this rule do exist, including
the "faint fuzzy'' star clusters observed in some lenticular galaxies
(Brodie & Larsen 2002), and there is a general trend for the
sizes of globular clusters to increase as a function of galactocentric
distance (van den Bergh et al. 1991)
While both the FWHM and
are found to correlate with cluster mass,
least-squares power-law fits yield slopes that are substantially shallower
than for a constant-density relation, implying an increase in cluster density
as a function of mass.
Qualitatively, these results are in agreement with previous data for
young star clusters as well as old globular clusters.
Quantitatively, the relations show a large scatter and remain uncertain.
Ashman & Zepf (2001) have argued that an increasing star formation
efficiency as a function of cluster mass may explain (at least partially)
the lack of a strong size-mass relation.
Many of the youngest clusters have extended, shallow outer envelopes. This tendency seems to be a general one, noted previously for a few isolated cases in the Antennae, NGC 6946, and for LMC clusters. The structure of these young objects may hold important clues to the early dynamical evolution of clusters and the density distribution of the parent proto-cluster clouds. Older clusters gradually evolve towards King-type profiles with a finite tidal radius.
Finally, a strong correlation between cluster age and crowding is found, with most of the strongly crowded clusters having ages <107 years. About 1/3-1/2 of these young objects are double or multiple sources, but the identification as bona-fide clusters is often uncertain even on WFPC2 images in these fairly nearby galaxies. Future multi-colour imaging with the Advanced Camera for Surveys on HST should help resolve many of the difficulties encountered in this study.
Acknowledgements
This work was partially supported by National Science Foundation grant AST 02-06139 and by HST grant AR-09523. I am grateful to T. Richtler and S. M. Fall for several helpful comments, and to the referee, R. de Grijs, for a very detailed report which helped improve the paper.