A&A 492, 685-693 (2008)
DOI: 10.1051/0004-6361:200810251
R. de Grijs1,2 - S. P. Goodwin1 - M. B. N. Kouwenhoven1 - P. Kroupa3
1 - Department of Physics & Astronomy, The University of
Sheffield, Hicks Building, Hounsfield Road, Sheffield S3 7RH, UK
2 -
National Astronomical Observatories, Chinese Academy of Sciences,
20A Datun Road, Chaoyang District, Beijing 100012, PR China
3 -
Argelander-Institut für Astronomie, Universität Bonn, Auf dem
Hügel 71, 53347 Bonn, Germany
Received 23 May 2008 / Accepted 19 September 2008
Abstract
The diagnostic age versus mass-to-light ratio diagram is
often used in attempts to constrain the shape of the stellar initial
mass function (IMF) and the potential longevity of extragalactic young
to intermediate-age massive star clusters. Here, we explore its
potential for Galactic open clusters. On the basis of a small,
homogenised cluster sample, we provide useful constraints on the
presence of significant binary fractions. Using the massive young
Galactic cluster Westerlund 1 as a key example, we caution that
stochasticity in the IMF introduces significant additional
uncertainties. We conclude that, for an open cluster to survive for
any significant length of time, and in the absence of substantial
external perturbations, it is necessary but not sufficient to be
located close to or (in the presence of a significant binary
population) somewhat below the predicted photometric
evolutionary sequences for ``normal'' simple stellar populations,
although such a location may be dominated by a remaining ``bound''
cluster core and thus not adequately reflect the overall cluster
dynamics.
Key words: Galaxy: open clusters and associations: individual: Westerlund 1, NGC 1976, Hyades, Coma Berenices, Orion Nebular Cluster - stellar dynamics - methods: observational - Galaxy: open clusters and associations: general
Over the past few years, detailed studies of the stellar content and
longevity of extragalactic massive star clusters have
increasingly resorted to the use of the age versus mass-to-light
ratio (M/L) diagram as a diagnostic tool, where one usually compares
dynamically determined M/L's with those predicted by the evolution
of ``simple'' stellar populations (SSPs; e.g., Smith & Gallagher 2001; Mengel
et al. 2002; McCrady et al. 2003;
Larsen et al. 2004; McCrady et al. 2005;
Bastian et al. 2006; Goodwin & Bastian 2006;
de Grijs & Parmentier 2007;
Moll et al. 2008). Based on
high-resolution spectroscopy to obtain the objects' line-of-sight (1D)
velocity dispersions,
,
and on high spatial resolution imaging
to obtain accurate, projected half-light radii
,
,
most authors then calculate the
dynamical cluster masses,
,
using
Nevertheless, using this approach one can get at least an initial
assessment as to whether a given (unresolved) cluster may be (i)
significantly out of virial equilibrium, in particular ``super-virial'';
(ii) substantially over- or underabundant in low-mass stars; or (iii)
populated by a large fraction of binary and higher-order multiple
systems. Since the work by Bastian & Goodwin (2006) and Goodwin &
Bastian (2006), we can now also model any (super-virial) deviations
from the SSP models for the youngest ages (up to 40 Myr) if we
assume that these are predominantly due to clusters being out of
virial equilibrium after gas expulsion.
This has led a number of authors to suggest that, in the absence of
significant external perturbations, massive clusters located in the
vicinity of the SSP models and aged 108 yr may survive for a
Hubble time and eventually become old globular cluster (GC)-like
objects (e.g., Larsen et al. 2004;
Bastian et al. 2006; de Grijs & Parmentier 2007).
Encouraged by the recent progress in this area based on both observational and theoretical advances, in this paper we set out to address the following unresolved questions:
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Figure 1: Diagnostic age versus M/LV diagram, including the Galactic open clusters for which velocity dispersion measurements are available. The evolution expected for SSPs governed by both Salpeter (1955) and Kroupa (2001) IMFs is shown as the solid and short-dashed lines, respectively. The long-dashed lines represent the evolution expected for SSPs with a Kroupa (2001)-type IMF, but a range of effective star-formation efficiencies (Goodwin & Bastian 2006). The sizes of the error bars are based on the most realistic ranges of observable values used to calculate the clusters' loci in this diagram. Numbered clusters: 1, NGC 1976 (Orion Nebula Cluster); 2, NGC 2168; 3, NGC 2516; 4, NGC 2632; 5, NGC 2682; 6, NGC 3532; 7, NGC 5662; 8, NGC 6705; 9, Pleiades; 10, Coma Berenices; 11, Hyades. |
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Table 1: Observational parameters of the open cluster sample.
Table 2: Derived parameters of the open cluster sample
Table 3: References to Tables 1 and 2.
In order to test the usefulness of the diagnostic diagram of cluster
age versus M/L for Galactic open clusters (see Fig. 1,
which we will discuss in detail in Sect. 3), we rely on
published parameters. Since each of the observables has an associated
uncertainty, it is paramount that we base our results on data sets
that are as homogeneous as possible. The most crucial ingredient for
the dynamical M/L determination is the internal velocity
dispersion. We include only those Galactic open clusters for which
these velocity dispersions have been derived from the proper motions
of the individual stars (for NGC 3532 the internal velocity
dispersion used in this paper is based on individual stellar radial
velocity measurements, Gieseking 1981). Where possible we include the
core velocity dispersions, in order to match the structural
parameters we will use. We also require well-determined distances (to
obtain luminosities and linear velocity dispersions) as well as core
radii and photometric observables. The distance estimates used here
are mostly based on the recent homogenised compilations of Kharchenko
et al. (2005) and Dias et al. (2006),
supplemented with determinations
based on a number of studies focusing on individual clusters. Although
Kharchenko et al. (2005) provide values for the core radii of many of
our sample clusters, the associated uncertainties are large. In fact,
they often dominate our dynamical mass estimates, together with the
often large uncertainties in the integrated V-band magnitudes. The
latter are often difficult to obtain to any reasonable degree of
accuracy because of the crowded fields in which many of the clusters
are located and also because of uncertain stellar cluster membership
determinations (see also the discussion in Sect. 3.2).
Nevertheless, in Table 1 we have collected the
``best'' values for the core radii, ,
distances, D,
apparent V-band magnitudes, foreground extinction, E(B-V), and
velocity dispersions,
.
We also list their likely
(uncertainty) ranges for our sample clusters. We provide for both the
values and the uncertainties the references we have used, and we have
aimed to homogenise our cluster sample parameters (following a similar
procedure as Paunzen & Netopil 2006, although they used different
selection criteria). This implies that our choice of the ``best'' values
for certain parameters may depend on the values of one or more of the
other observables. We provide the full list of references used to
obtain the most likely parameter ranges. However, where we have
discarded certain values (often because they were clear statistical
outliers), the relevant respective references are bracketed.
In Table 2 we list the best ages and their uncertainty
ranges of our sample clusters using the same notation as in
Table 1, as well as the total cluster masses - based on
Eq. (1), with
- and their LV/M ratios
derived based on the parameters and their (
1
)
uncertainties from Table 1. A full list of
references to Tables 1 and 2 is
provided in Table 3.
We note that our sample selection is biased towards the nearest Galactic open clusters, for which reasonably accurate internal velocity dispersions could be obtained. However, although our sample is by no means complete in any sense, we can still use it to assess (i) the binary fractions of the clusters individually (Sect. 3); and (ii) the usefulness of the diagnostic age versus M/L diagram for cluster longevity considerations (Sect. 4).
Using the observational data from Sect. 2, we applied Eq. (1) to derive the dynamical masses for each of our sample cluster cores and calculated the relevant M/LV's. Their loci in the diagnostic diagram are shown in Fig. 1. Overplotted is the expected evolution of SSPs (Maraston 2005) for both a Salpeter (1955) and a Kroupa (2001) stellar IMF (solid and short-dashed lines, respectively).
We have also included the expected evolution of clusters formed with a
variety of effective star-formation efficiencies (eSFEs; Goodwin
& Bastian 2006). The eSFE is a measure of the extent to which a
cluster is out of equilibrium after gas expulsion, on the basis that
the virial ratio immediately before gas expulsion was
(eSFE) (where T and
are the kinetic
and potential energy of the stars, respectively, and a system in
virial equilibrium has
). The eSFE corresponds to
the true SFE if the stars and gas were initially in virial equilibrium
(see Goodwin & Bastian 2006).
Owing to the nature of our sample, only the Orion Nebula Cluster (ONC, cluster 1 in
Fig. 1) is currently young enough so as to possibly be
affected by the effects of rapid gas expulsion, as shown by the extent
(in terms of age) of the long-dashed lines in Fig. 1. The
majority of our sample clusters are old enough (
40 Myr) to have
re-virialised after gas expulsion. The dynamical state of these
objects is therefore dominated by the combined effects of (internal)
two-body relaxation, binary motions, and external perturbations.
The fact that our sample of surviving open cluster cores lie close to the SSP predictions should be expected. Clusters significantly below the SSP lines will be dynamically ``hot'' and are expected to dissolve rapidly, whilst clusters significantly above the lines will be dynamically ``cold'' and should (re-)virialise over a few crossing times to move closer to the canonical SSP lines.
We will now explore the reasons as to why most of our sample clusters (cores) are found somewhat below the SSP model curves (i.e., they seem somewhat supervirial with respect to the expectations from the SSP models), irrespective of whether or not they actually follow the SSP predictions or are characterised by roughly constant M/L's as a function of age. We expect errors in the core radii to be random, and unbiased by the mass of a cluster. However, the use of the core velocity dispersions and radii may introduce a systematic bias in the dynamical mass estimates.
There are three dynamical effects that could affect the position of the clusters relative to the SSP predictions.
First, equipartion and mass segregation could lower the core
velocity dispersion relative to the ``typical'' velocity and hence move
the clusters' positions to above the SSP predictions. The majority of
the star clusters in our sample are older than 108 yr, which
implies that they have ages greater than their half-mass relaxation
times (see, e.g., Danilov & Seleznev 1994, for the relevant
time-scales for most of our sample clusters). Therefore, these
clusters (and particularly their cores) are expected to be close to
energy equipartition, and thus are significantly mass segregated - as
observed for, e.g., NGC 2168 (Sung & Bessell 1999; Kalirai
et al. 2003), NGC 2682 (Bonatto & Bica 2003), and NGC 6705 (e.g., Sung
et al. 1999, and references therein) among our present sample.
Equipartition reduces the global velocity dispersion of
high-mass stars relative to low-mass stars, causing high-mass stars to
migrate to the cluster core. Therefore, we might expect the core
velocity dispersion of low-mass cores (as characteristic for the open
clusters discussed in this paper) to underestimate the dynamical mass
of the clusters as a whole and thus produce colder
clusters
- as seen in Fig. 1 (although we remind
the reader of the expected re-virialisation discussed above; this may
introduce an observational bias in the sense that we would not be able
to detect low-mass cluster cores that are significantly super-virial
and hence - possibly - in the process of dissolution).
Secondly, the clusters' mass functions (MFs) will have been
altered by dynamical evolution, with the preferential loss of low-mass
stars moving clusters to above the SSP predictions. This will result
in a ``top-heavy'' MF in clusters, which will in turn lead to lower
M/LV's than would be expected from the canonical SSP models. The
degree to which the MF will change depends on the two-body relaxation
time which, to first order, depends on the mass of the cluster (and
also on its size; however, we ignore this for now). Thus, old low-mass
clusters are expected to have top-heavy MFs compared to old high-mass
clusters. Therefore, we would expect low-mass clusters to lie some way
above the canonical SSP models, and high-mass clusters to lie slightly
above these lines.
Thirdly, the presence of binaries may result in an observed
velocity dispersion that is higher than the ``true'' value, moving the
clusters to below the SSP predictions. Kouwenhoven & de Grijs (2008)
pointed out that if the velocity dispersion of binary systems was
similar to the velocity dispersion of the cluster (core) as a whole,
the observationally measured velocity dispersion would
overestimate the mass of a cluster. We can explore, to first order,
whether the binary population may be a significant factor causing an
offset in Fig. 1 by using the new diagnostic proposed by
Kouwenhoven & de Grijs (2008, their Fig. 9).
In Fig. 2 we reproduce the main features of their Fig. 9, and include our open
cluster sample (using the clusters' core radii instead of their
half-mass radii; the core radii are more likely to represent the size
of the bound stellar population for these clusters; see, e.g.,
Odenkirchen et al. 1998, for arguments relating to the
open cluster in Coma Berenices). It is immediately clear from the
location of the data points that the vast majority of our sample
clusters are indeed expected to be significantly affected by binaries
(
,
cf. Fig. 2; in fact, the data points
represent upper limits to the cluster masses given that we do not know
the intrinsic masses but need to rely on dynamical tracers). This
seems to be borne out by relevant recent observations of a number of
our sample clusters, including NGC 2516 (cf. Sung et al. 2002),
NGC 2632 (M 44, also known as the Praesepe cluster: see, e.g., Bouvier
et al. 2001; Patience et al. 2002, and references therein), the Pleiades
(e.g., Martin et al. 2000, and references therein), and the Hyades
(e.g., Stefanik & Latham 1992; Patience et al. 1998).
That the cluster cores appear to lie below the SSP predictions seems to suggest that the effect of binaries outweighs mass segregation and the change in the MF in determining the position of the cluster cores in the diagram. As shown by Kouwenhoven & de Grijs (2008), this is to be expected for relatively low-mass open clusters such as the objects we are considering here. However, without a detailed investigation of each cluster and their component stars, it is impossible to reconstruct the degree to which each effect is important. However, we argue that it seems clear that the position of most clusters (cores) below the canonical SSP lines is not due to significant deviations from virial equilibirium.
Finally, we also note that we may well have overestimated the masses
by factors of a few through the universal use of
Eq. (1). For highly mass-segregated clusters containing
significant binary fractions, a range of stellar IMF representations,
and for combinations of characteristic relaxation time-scales and
cluster half-mass radii, the adoption of a single scaling factor
introduces systematic offsets, leading to lower values
of
(e.g., Fleck et al. 2006; Kouwenhoven & de Grijs 2008), and
thus to dynamical mass overestimates if
were
assumed. However, the uncertainties are too large at the present time
to reach firm conclusions regarding the dependence of our results
on
.
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Figure 2:
Diagnostic ![]() ![]() |
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The Galactic young massive star cluster, Westerlund 1 (aged 4-5 Myr;
Crowther et al. 2006), and in particular its stellar content, has been
the subject of considerable recent attention (e.g., Clark et al. 2005,
2008; Crowther et al. 2006;
Muno et al. 2006; Mengel & Tacconi-Garman
2007a,b; and references therein). It is the nearest potential
GC progenitor, and certainly the most massive young Galactic cluster
(
,
with an absolute lower limit of
;
Clark et al. 2005;
see also Mengel & Tacconi-Garman
2007a). In
order for the cluster to survive, it cannot have a stellar IMF that is
deficient in low-mass stars. Given that all observed star clusters
exhibit a range in stellar masses, we conclude that the Westerlund
1 IMF must therefore be close to ``normal'' (there is no conclusive
evidence of clusters with ``bottom-heavy'' IMFs, which could
potentially also lead to the cluster's position in
Fig. 3). In de Grijs & Parmentier (2007) we reviewed the
balance of evidence (e.g., Muno et al. 2006; Clark et al. 2008), but
the results remained inconclusive because of the difficulty of
observing the cluster's low-mass stellar population. Brandner
et al. (2008) recently completed a detailed study of the cluster's mass
function down to
,
which appears to be consistent with
a normal Kroupa or Salpeter-type IMF.
In addition, the relaxation of an idealised cluster and the
contribution of the most massive stars to the escape of stars below a
typical limiting (high) mass scales approximately in a power-law
fashion (with a power 3) with mass (e.g., Hénon 1969,
and references therein). This is key to understanding the dynamical
importance of a particular IMF. It follows that the escape rate of
low-mass stars below a certain mass is mass-independent. Moreover, in
an equilibrium system the number of stars escaped from the cluster by
an age of 4 Myr would therefore be small and hardly affect the overall
shape of the IMF. Hence, only a small modification of the IMF below
the supernova mass limit would be expected at the present age, and the
observed IMF will therefore be close to ``normal''.
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Figure 3: Westerlund 1 in the diagnostic age versus M/LV diagram. The open circle represents the cluster's locus if we were to exclude the nine brightest stars; this exemplifies the uncertainties introduced by stochastic IMF sampling and by having fortuitously caught the cluster at a time when it is dominated by a small number of very bright stars. The line coding is as in Fig. 1. |
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Using the dynamical mass estimate from Mengel & Tacconi-Garman
(2007a), combined with the integrated photometry of Piatti
et al. (1998), we reached a similar conclusion (de Grijs & Parmentier
2007), despite the significant uncertainties in the observables. Since
in the V band, on which the Piatti et al. (1998) photometry was
based, the confusion between the cluster members and the Galactic
field stellar population is substantial (in essence because of the
significant extinction along this sightline), we obtained imaging
observations at longer wavelengths, where this confusion is
significantly reduced. An I-band (peak-up) image of the cluster
(using the Ic/Iwp-ESO0845 filter), with an exposure time of 3.0 s, was
obtained with the ESO 2.2 m telescope equipped with the Wide-Field
Imager (WFI) at La Silla Observatory (Chile). The image was kindly
made available to us by Crowther. Using the photometric
zero-point offsets of Clark et al. (2005) we obtained an integrated
I-band magnitude of
mag within a radius of
108 arcsec. This includes all of the bright cluster members and
excludes bright foreground sources.
The combined integrated magnitude of the three brightest red
supergiants (objects 26, 237, and 20 of Clark et al. 2005, in order of
decreasing brightness), yellow hypergiants (objects 32, 4, and 8) and
blue supergiants (objects 243, 16, and 7) is
mag. Therefore, these nine sources alone contribute some 40% of the
cluster's total integrated I-band flux. Each of these sources is in
a rare, short-lived phase and so the luminosity of the cluster might
be expected to vary significantly on short time-scales. In addition,
we specifically discuss these nine brightest cluster members
separately, because these are the stars that make Westerlund 1 one of
the most unusual young star clusters known (e.g., Clark et al. 2005,
2008; Crowther et al. 2006). These nine sources stand out from the
overall stellar luminosity function, which appears to otherwise have
been drawn from a ``normal'' IMF. Thus, this serves as a clear caution
that stochasticity in the cluster's IMF (e.g., Brocato et al. 2000),
as well as stochasticity in the numbers of stars in unusually luminous
post-main-sequence evolutionary stages (e.g., Cerviño &
Valls-Gabaud 2003; Cerviño & Luridiana 2006) may contribute
significantly to variations in a cluster's observed M/L. On a
related note, we caution that the luminosities of all of the
clusters we analyse in this paper are subject to stochastic effects,
regardless of their age (Cerviño & Valls-Gabaud 2003;
Cerviño & Luridiana 2006).
Using the most up-to-date distance to Westerlund 1,
kpc (Kothes & Dougherty 2007), a V-band extinction of
AV = 11.6 mag (Clark et al. 2005), and the Galactic reddening law of Rieke &
Lebofsky (1985, resulting in AI = 5.6 mag), we obtain the locus in
(age versus L/M ratio) space as shown in Fig. 3.
Despite the large error bars and very young age, it appears at first
sight that Westerlund 1 is not significantly out of virial
equilibirum. Its location in Fig. 3 is consistent with
the cluster having formed with a high eSFE, and with a Kroupa or
Salpeter-like stellar IMF. Given the different filters used between
Piatti et al. (1998) and this paper, and in view of the updated
cluster mass estimate, this result confirms our earlier assertion in
de Grijs & Parmentier (2007) based on the cluster's location in the
M/LV versus age diagram (which in turn supported the conclusion of
Mengel & Tacconi-Garman 2007a, that the cluster appears to be close to
virial equilibrium). However, we note that it is not entirely clear if
Westerlund 1 would be expected to be in virial equilibrium. Bastian
et al. (2008) show that observations of massive young clusters suggest
that they may expand from initial half-mass radii of
0.1 pc to
>1 pc in the first few Myr of their lives (see, in particular,
their Fig. 4). In such a situation we might expect Westerlund 1 to lie
well below the canonical IMF lines.
If we were to exlude the nine brightest stars making up some 40% of the cluster's integrated I-band flux, its locus would shift to that of the open circle (assuming that the cluster's mass remains unchanged). We will now briefly explore whether this effect would be significantly different in the V band, as discussed in de Grijs & Parmentier (2007). Although we do not have unsaturated V-band images of Westerlund 1, we can use the results of Piatti et al. (1998, their Fig. 8) to check the above statements to first order. We base our analysis on the simplistic assumption that the innermost nine stars are the nine brightest stars.
Based on this figure, the nine innermost stars contribute a combined
mag; the full integrated cluster magnitude is
mag. We therefore conclude that these
nine innermost stars contribute
30% of the cluster's total
luminosity. Some, but not all, of these stars are clearly the very
bright stars we used in our I-band analysis, so that this estimate
provides a lower limit to the contribution of the nine brightest
stars. Given that we found that in the I band the nine brightest
stars contribute
40% of the total flux, the V and I-band
contributions of the nine brightest stars are similar, particularly in
view of the uncertainties. In fact, this could have been expected -
to first order - because the spectral energy distributions of each of
the three subgroups (blue and red supergiants and yellow hypergiants)
peak at different wavelengths. Therefore, stochasticity remains a
serious issue across these wavelengths, simply because these stars are
intrinsically so bright.
This shows the potential effects of (i) stochastic sampling of a
cluster's IMF (predominantly affecting the highest-mass end in any
cluster); and (ii) having caught the cluster at a time when it is
dominated by a few very luminous yet short-lived stars. If indeed we
are fortuitous in having observed a stochastically exceptional
situation regarding the numbers of very massive (bright) stars in the
cluster, it would indicate that Westerlund 1 may have formed with an
eSFE around the 30-40% required for clusters to survive the
gas expulsion phase (e.g., Lada et al. 1984;
Goodwin 1997a,b; Adams 2000;
Geyer & Burkert 2001;
Kroupa & Boily 2002;
Boily & Kroupa 2003a,b;
Fellhauer & Kroupa 2005;
Bastian & Goodwin 2006;
Parmentier & Gilmore 2007, their Fig. 1), although we note that
the error bars are large and will remain unchanged by the removal of
these nine brightest stars. As an aside, we note that differential
extinction towards the individual brightest cluster stars is not an
issue; extinction variations along individual sight lines are minimal,
with deviations from a mean
-band extinction of
mag (e.g., Crowther et al. 2006); this
narrow spread is due to the extinction estimates being based on
assumed intrinisic stellar colours, which in reality vary slightly.
The dynamical state of the core of the Orion Nebula Cluster (NGC 1976,
M 42; cluster 1 in Fig. 1) has been the subject of
significant observational and theoretical investigations (e.g.,
Hillenbrand & Hartmann 1998;
Kroupa et al. 1999; Kroupa 2000;
Kroupa et al. 2001;
O'Dell 2001; Scally et al. 2005,
and references therein). It is the youngest cluster in our
sample and is located (3
)
below the ``normal'' SSP
evolution in Fig. 1, even in view of the
uncertainties. This super-virial state is corroborated by current
estimates of its virial ratio, which suggest that the cluster (core)
is already unbound, but has only recently become so (e.g., Kroupa
et al. 2001; Scally et al. 2005). In fact, Hillenbrand & Hartmann (1998)
showed that in order for the ONC to be in virial equilibrium, based on
the cluster's observed velocity dispersion, the total mass within
about 2 pc of the central ``Trapezium'' configuration of massive stars
must be of order twice that of the known stellar population in the
region (and comparable to the estimated mass in molecular gas
projected onto the area). Given the youth of the cluster, and its
partially embedded nature, Hillenbrand & Hartmann (1998) argued that
if
20% of the remaining molecular gas is converted into stars,
this might result in a gravitationally bound cluster. Follow-up
N-body simulations led Scally et al. (2005) to conclude that the
size and age of the ONC imply that either the cluster is marginally
bound (or has become unbound only very recently), or else that it has
expanded quasi-statically. Kroupa et al. (2001), on the other hand,
performed binary-rich N-body models of the ONC adopting two of
the allowed initial configurations from Kroupa (2000) and showed that
it is currently expanding and was probably formed with an eSFE near 33%. In view of the uncertainties, this is roughly consistent with
its locus in the diagnostic diagram of Fig. 1.
Odenkirchen (1998) found that the open cluster in Coma Berenices
(cluster 10 in Fig. 1) has an elliptical core-halo
morphology, combined with a group of extratidal stars
either escaping stars or genuine field stars; see also Küpper et al. (2008), for a
general discussion on the distribution of escaped stars), which are
located at projected distances of 10 pc from the cluster
centre. They provide some tentative evidence of the presence of an
additional population of even lower-mass extratidal stars, and argue
that the existence of this significant population of stars beyond the
cluster's tidal radius is evidence of the cluster dissolution process
caught in the act. At the same time, they conclude that - given the
present mass and configuration of the cluster stars - the observed
(core) velocity dispersion is fully consistent with the expectations
from the SSP models.
Here we reach, in essence, the same conclusion. Based on the observational data at hand, the core of the star cluster in Coma Berenices is located very close to the expected photometric evolutionary sequences in Fig. 1, within reasonably small uncertainties. Given that there is evidence that this cluster is in the advanced stages of dissolution, this result should be considered as a strong caution. It appears that for a cluster to survive for a significant length of time, it is a necessary, but not a sufficient condition for it to be located close to the evolutionary sequences in our diagnostic diagram. We caution, however, that since our velocity dispersion measurements were weighted towards the central region of the cluster, it is possible that the cluster's locus in Fig. 1 mainly reflects its remaining bound component.
The Hyades (cluster 11 in Fig. 1) is a dynamically very
evolved, marginally bound cluster significantly depleted in low-mass
stars (e.g., Kroupa 1995; Perryman et al. 1998;
see also Portegies Zwart et al. 2001), with a stellar velocity dispersion of order
0.3-0.4 km s-1 (Makarov et al. 2000; see also
Madsen 2003). Detailed N-body simulations (e.g., Terlevich 1987;
Madsen 2003; Chumak et al. 2005) indicate that a halo
of gravitationally unbound stars can still be linked with the cluster,
and that these stars are moving along with it on similar orbits
(cf. Küpper et al. 2008), for several
yr (see Perryman
et al. 1998, for a detailed discussion). Hence, at its current age of
(Paunzen & Netopil 2006,
and references therein) it is not surprising that the Hyades moving
group is still detectable as a cluster-type object.
In Fig. 1, the Hyades occupies a locus very close to the evolutionary sequences (and with small error bars), yet the group is likely (i) unbound overall; and (ii) in the final stages of dissolution (Odenkirchen et al. 1998). Although the same caution applies to the Hyades moving group as to the cluster in Coma Berenices, we conclude again that for a cluster to survive for a significant length of time, it is a necessary, but not a sufficient condition for it to be located close to the evolutionary sequences in our diagnostic diagram.
In this paper, we have explored the usefulness of the diagnostic age versus M/L diagram in the context of Galactic open clusters. This diagram is often used in the field of extragalactic young to intermediate-age massive star clusters to constrain the shape of their stellar IMF, as well as their stability and the likelihood of their longevity.
Using a sample of Galactic open clusters for which reasonably accurate internal (core) velocity dispersions are available in the literature, we constructed a homogenised set of observational data drawn from a wide variety of publications, also including their most likely uncertainty ranges. This allowed us to derive dynamical mass estimates for our sample of open clusters, as well as their respective M/LV's and - crucially - the associated (realistic) uncertainties.
It seems clear that the effect of binaries, mass segregation, and the dynamical alteration of mass functions by two-body relaxation are important constraints that cannot be ignored.
Using the massive young Galactic cluster Westerlund 1 as a key example, we caution that stochasticity in the IMF introduces significant additional uncertainties. Therefore, the stability and long-term survival chances of Westerlund 1 remain inconclusive.
Most importantly, however, we conclude that for an open cluster to
survive for any significant length of time (in the absence of
substantial external perturbations), it is a necessary but not a
sufficient condition to be located close to the predicted photometric
evolutionary sequences for ``normal'' SSPs. This is highlighted using a
number of our sample clusters (and the parameters related to the
cluster cores) which are known to be in a late stage of dissolution,
and lie very close indeed to either of the evolutionary sequences
defined by the Salpeter (1955) or Kroupa (2001) IMFs. However, we
also note that a fair fraction of our sample clusters show the
signatures of dynamical relaxation and stability. Among our current
sample, these include NGC 2168 (M 35, Kalirai et al. 2003), NGC 2682
(M 67, Hurley et al. 2005), NGC 6705
(M 11, Mathieu 1984; McNamara & Sekiguchi 1986;
Sung et al. 1999) and the Pleiades (M 45;
McNamara & Sekiguchi 1986; Pinfield et al. 1998;
Raboud & Mermilliod 1998). Despite their relatively low masses (
)
and ages in excess of a few
yr,
this is not unexpected.
Using the vertical oscillation period, ,
around the Galactic
plane of NGC 2323 (
Myr; Clariá et al. 1998) as an
example, this cluster has only been through a few of these periods,
given its age of
(Kalirai
et al. 2003). However, at the Galactocentric distance of the Sun, a
Pleiades-like open cluster crosses the Galactic disc approximately
10-20 times before it dissolves (de la Fuente Marcos
1998a,b). The
models of Kroupa et al. (2001), which match the ONC at an age of 1 Myr
very well, as well as the Pleiades at 100 Myr, suggest that these
objects would end up below the evolutionary sequences in
Fig. 1, despite having started from the Kroupa (2001) IMF
at birth. The deviation may have been caused by the heating of the
clusters by the Galactic tidal field. In other words, it seems that
the velocity dispersion is always somewhat higher after Galactic-plane
passage, because the stars suffer from an additional acceleration.
Similarly, the age of NGC 2516 (cluster 3 in Fig. 1),
(Sung et al. 2002), is well in
excess of its period of vertical oscillations through the Galactic
plane,
yr (Dachs & Kabus
1989).
NGC 2516 is presently located some 120 pc below the Galactic plane, near
the dense molecular clouds of the Vela Sheet. These Galactic plane
passages may have contributed to rendering the cluster unstable.
Alternatively, encounters with giant molecular clouds (e.g., Gieles
et al. 2006), particularly around the time of Galactic plane passages,
may have contributed to the cluster's present dynamical state.
In addition, Bonatto & Bica (2003) show that tidal losses of stars
from NGC 2682 to the Galactic field have been effective (see also
McNamara & Sekiguchi 1986 for, e.g., NGC 2168 and NGC 2632). This
interpretation is supported by the N-body simulations of Hurley
et al. (2005). Additional (circumstantial) evidence of tidal effects
acting on NGC 2682 is present in the form of significantly elliptical
cluster isophotes (Fan et al. 1996), which might be a tidal extension
caused by the Galactic field (Bergond et al.
2001). In addition, Chupina & Vereshchagin (1998) detected several
density enhancements in the low-density extended outskirts of the
cluster. Such clumps are expected as a consequence of disc shocking
(e.g., Bergond et al. 2001). Alternatively, the cluster may have
undergone a number of encounters with giant molecular clouds, possibly
leading to a similar morphology. Similarly, Adams et al. (2001)
suggest that the flattening of the stellar mass function of the
Pleiades below
with respect to the
field-star population may have been caused by evaporation of the
lowest-mass stars into the Galactic field (see also van Leeuwen 1983),
although evidence that this may be the case remains inconclusive
(see, e.g., the simulations of de la Fuente Marcos 2000;
Moraux et al. 2004).
In a follow-up paper (Kouwenhoven et al., in prep.) we will quantitatively explore the loci in the diagnostic diagram of the Galactic open clusters, using N-body simulations.
Acknowledgements
We acknowledge research support and hospitality at the International Space Science Institute in Bern (Switzerland), as part of an International Team programme. We thank Paul Crowther for discussions and the use and basic analysis of his unpublished imaging of Westerlund 1, which was obtained as part of ESO proposal 69.D-0327(B). We are particularly grateful to Mark Gieles for very helpful comments. We also thank Vladimir Danilov for providing us with some hard-to-find data for a few of our sample clusters and acknowledge the constructive criticism in the referee report of Sverre Aarseth. R.d.G. and M.K. acknowledge financial support from STFC grant PP/D002036/1; R.d.G. also acknowledges partial support from the Royal Society in the context of a ``Frontiers of Science'' programme. This research has made use of the SIMBAD database, operated at CDS, Strasbourg (France), and of NASA's Astrophysics Data System Abstract Service.