ApJS, in press; version with embedded figures can be obtained
at http://spider.ipac.caltech.edu/staff/stauffer/
Near and Mid-IR Photometry of the Pleiades, and a New List of
Substellar Candidate Members1,2
John R. Stauffer
Spitzer Science Center, Caltech 314-6, Pasadena, CA 91125
stauffer@ipac.caltech.edu
Lee W. Hartmann
Astronomy Department, University of Michigan
Giovanni G. Fazio, Lori E. Allen, Brian M. Patten
Harvard-Smithsonian Center for Astrophysics, 60 Garden St., Cambridge, MA 02138
Patrick J. Lowrance, Robert L. Hurt, Luisa M. Rebull
Spitzer Science Center, Caltech , Pasadena, CA 91125
Roc M. Cutri, Solange V. Ramirez
Infrared Processing and Analysis Center, Caltech 220-6, Pasadena, CA 91125
Erick T. Young, George H. Rieke, Nadya I. Gorlova3, James C. Muzerolle
Steward Observatory, University of Arizona, Tucson, AZ 85726
Cathy L. Slesnick
Astronomy Department, Caltech, Pasadena, CA 91125
Michael F. Skrutskie
Astronomy Department, University of Virginia, Charlottesville, VA 22903
1This work is based (in part) on observations made with the Spitzer Space Telescope, which is operated
by the Jet Propulsion Laboratory, California Institute of Technology, under NASA contract 1407.
2This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project
of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute
of Technology, funded by the National Aeronautics and Space Administration and the National Science
Foundation.
3Current address: University of Florida, 211 Bryant Space Center, Gainesville, FL 32611
http://arxiv.org/abs/0704.1832v1
http://spider.ipac.caltech.edu/staff/stauffer/
– 2 –
ABSTRACT
We make use of new near and mid-IR photometry of the Pleiades cluster in
order to help identify proposed cluster members. We also use the new photometry
with previously published photometry to define the single-star main sequence
locus at the age of the Pleiades in a variety of color-magnitude planes.
The new near and mid-IR photometry extend effectively two magnitudes
deeper than the 2MASS All-Sky Point Source catalog, and hence allow us to
select a new set of candidate very low mass and sub-stellar mass members of the
Pleiades in the central square degree of the cluster. We identify 42 new candi-
date members fainter than Ks =14 (corresponding to 0.1 Mo). These candidate
members should eventually allow a better estimate of the cluster mass function
to be made down to of order 0.04 solar masses.
We also use new IRAC data, in particular the images obtained at 8 um,
in order to comment briefly on interstellar dust in and near the Pleiades. We
confirm, as expected, that – with one exception – a sample of low mass stars
recently identified as having 24 um excesses due to debris disks do not have
significant excesses at IRAC wavelengths. However, evidence is also presented
that several of the Pleiades high mass stars are found to be impacting with local
condensations of the molecular cloud that is passing through the Pleiades at the
current epoch.
Subject headings: stars: low mass — young; open clusters — associations:
individual (Pleiades)
1. Introduction
Because of its proximity, youth, richness, and location in the northern hemisphere, the
Pleiades has long been a favorite target of observers. The Pleiades was one of the first open
clusters to have members identified via their common proper motion (Trumpler 1921), and
the cluster has since then been the subject of more than a dozen proper motion studies.
Some of the earliest photoelectric photometry was for members of the Pleiades (Cummings
1921), and the cluster has been the subject of dozens of papers providing additional optical
photometry of its members. The youth and nearness of the Pleiades make it a particularly
attractive target for identifying its substellar population, and it was the first open cluster
studied for those purposes (Jameson & Skillen 1989; Stauffer et al. 1989). More than 20 pa-
pers have been subsequently published, identifying additional substellar candidate members
of the Pleiades or studying their properties.
– 3 –
We have three primary goals for this paper. First, while extensive optical photometry
for Pleiades members is available in the literature, photometry in the near and mid-IR is
relatively spotty. We will remedy this situation by using new 2MASS JHKs and Spitzer
IRAC photometry for a large number of Pleiades members. We will use these data to help
identify cluster non-members and to define the single-star locus in color-magnitude diagrams
for stars of 100 Myr age. Second, we will use our new IR imaging photometry of the center of
the Pleiades to identify a new set of candidate substellar members of the cluster, extending
down to stars expected to have masses of order 0.04 M⊙. Third, we will use the IRAC data
to briefly comment on the presence of circumstellar debris disks in the Pleiades and the
interaction of the Pleiades stars with the molecular cloud that is currently passing through
the cluster.
In order to make best use of the IR imaging data, we will begin with a necessary
digression. As noted above, more than a dozen proper motion surveys of the Pleiades have
been made in order to identify cluster members. However, no single catalog of the cluster
has been published which attempts to collect all of those candidate members in a single
table and cross-identify those stars. Another problem is that while there have been many
papers devoted to providing optical photometry of cluster members, that photometry has
been bewilderingly inhomogeneous in terms of the number of photometric systems used. In
Sec. 3 and in the Appendix, we describe our efforts to create a reasonably complete catalog
of candidate Pleiades members and to provide optical photometry transformed to the best
of our ability onto a single system.
2. New Observational Data
2.1. 2MASS “6x” Imaging of the Pleiades
During the final months of Two Micron All Sky Survey (2MASS; Skrutskie et al. (2006))
operations, a series of special observations were carried out that employed exposures six times
longer than used for the the primary survey. These so-called “6x” observations targeted 30
regions of scientific interest including a 3 deg x 2 deg area centered on the Pleiades cluster.
The 2MASS 6x data were reduced using an automated processing pipeline similar to that
used for the main survey data, and a calibrated 6x Image Atlas and extracted 6x Point and
Extended Source Catalogs (6x-PSC and 6x-XSC) analogous to the 2MASS All-Sky Atlas,
PSC and XSC have been released as part of the 2MASS Extended Mission. A description
of the content and formats of the 6x image and catalog products, and details about the 6x
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observations and data reduction are given by Cutri et al. (2006; section A3). 1 The 2MASS
6x Atlas and Catalogs may be accessed via the on-line services of the NASA/IPAC Infrared
Science Archive (http://irsa.ipac.caltech.edu).
Figure 1 shows the area on the sky imaged by the 2MASS 6x observations in the Pleiades
field. The region was covered by two rows of scans, each scan being one degree long (in
declination) and 8.5’ wide in right ascension. Within each row, the scans overlap by approx-
imately one arcminute in right ascension. There are small gaps in coverage in the declination
boundary between the rows, and one complete scan in the southern row is missing because
the data in that scan did not meet the minimum required photometric quality. The total
area covered by the 6x Pleiades observations is approximately 5.3 sq. degrees.
There are approximately 43,000 sources extracted from the 6x Pleiades observations
in the 2MASS 6x-PSC, and nearly 1,500 in the 6x-XSC. Because there are at most about
1000 Pleiades members expected in this region, only ∼2% of the 6x-PSC sources are cluster
members, and the rest are field stars and background galaxies. The 6x-XSC objects are
virtually all resolved background galaxies. Near infrared color-magnitude and color-color
diagrams of the unresolved sources from the 2MASS 6x-PSC and all sources in the 6x-XSC
sources from the Pleiades region are shown in Figures 2 and 3, respectively. The extragalactic
sources tend to be redder than most stars, and the galaxies become relatively more numerous
towards fainter magnitudes. Unresolved galaxies dominate the point sources that are fainter
than Ks > 15.5 and redder than J −Ks > 1.2 mag.
The 2MASS 6x observations were conducted using the same freeze-frame scanning tech-
nique used for the primary survey (Skrutskie et al. 2006). The longer exposure times were
achieved by increasing the “READ2-READ1” integration to 7.8 sec from the 1.3 sec used for
primary survey. However, the 51 ms “READ1” exposure time was not changed for the 6x
observations. As a result, there is an effective “sensitivity gap” in the 8-11 mag region where
objects may be saturated in the 7.8 sec READ2-READ1 6x exposures, but too faint to be
detected in the 51 ms READ1 exposures. Because the sensitivity gap can result in incom-
pleteness and/or flux bias in the photometric overlap regime, the near infrared photometry
for sources brighter than J=11 mag in the 6x-PSC was taken from the 2MASS All-Sky PSC
during compilation of the catalog of Pleiades candidate members presented in Table 2 (c.f.
Section 3).
1http://www.ipac.caltech.edu/2mass/releases/allsky/doc/explsup.html
http://irsa.ipac.caltech.edu
– 5 –
2.2. Shallow IRAC Imaging
Imaging of the Pleiades with Spitzer was obtained in April 2004 as part of a joint GTO
program conducted by the IRAC instrument team and the MIPS instrument team. Initial
results of the MIPS survey of the Pleiades have already been reported in Gorlova et al. (2006).
The IRAC observations were obtained as two astronomical observing requests (AORs). One
of them was centered near the cluster center, at RA=03h47m00.0s and Dec=24d07m (2000),
and consisted of a 12 row by 12 column map, with “frametimes” of 0.6 and 12.0 seconds
and two dithers at each map position. The map steps were 290′′ in both the column and
row direction. The resultant map covers a region of approximately one square degree, and
a total integration time per position of 24 sec over most of the map. The second AOR used
the same basic mapping parameters, except it was smaller (9 rows by 9 columns) and was
instead centered northwest from the cluster center at RA=03h44m36.0s and Dec=25d24m.
A two-band color image of the AOR covering the center of the Pleiades is shown in Figure 4.
A pictorial guide to the IRAC image providing Greek names for a few of the brightest stars,
and Hertzsprung (1947) numbers for several stars mentioned in Section 6 is provided in
Figure 5.
We began our analysis with the basic calibrated data (BCDs) from the Spitzer pipeline,
using the S13 version of the Spitzer Science Center pipeline software. Artifact mitigation
and masking was done using the IDL tools provided on the Spitzer contributed software
website. For each AOR, the artifact-corrected BCDs were combined into single mosaics
for each channel using the post-BCD “MOPEX” package (Makovoz & Marleau 2005). The
mosaic images were constructed with 1.22×1.22 arcsecond pixels (i.e., approximately the
same pixel size as the native IRAC arrays).
We derived aperture photometry for stars present in these IRAC mosaics using both
APEX (a component of the MOPEX package) and the “phot” routine in DAOPHOT. In
both cases, we used a 3 pixel radius aperture and a sky annulus from 3 to 7 pixels (except
that for Channel 4, for the phot package we used a 2 pixel radius aperture and a 2 to 6 pixel
annulus because that provided more reliable fluxes at low flux levels). We used the flux for
zero magnitude calibrations provided in the IRAC data handbook (280.9, 179.7, 115.0 and
64.1 Jy for Ch 1 through Ch 4, respectively), and the aperture corrections provided in the
same handbook (multiplicative flux correction factors of 1.124, 1.127, 1.143 and 1.584 for Ch
1-4, inclusive. The Ch4 correction factor is much bigger because it is for an aperture radius
of 2 rather than 3 pixels.).
Figure 6 and Figure 7 provide two means to assess the accuracy of the IRAC photometry.
The first figure compares the aperture photometry from APEX to that from phot, and shows
that the two packages yield very similar results when used in the same way. For this reason,
– 6 –
we have simply averaged the fluxes from the two packages to obtain our final reported value.
The second figure shows the difference between the derived 3.6 and 4.5 µm magnitudes
for Pleiades members. Based on previous studies (e.g. Allen et al. (2004)), we expected
this difference to be essentially zero for most stars, and the Pleiades data corroborate that
expectation. For [3.6]<10.5, the RMS dispersion of the magnitude difference between the two
channels is 0.024 mag. Assuming that each channel has similar uncertainties, this indicates
an internal 1-σ accuracy of order 0.017 mag. The absolute calibration uncertainty for the
IRAC fluxes is currently estimated at of order 0.02 mag. Figure 7 also shows that fainter than
[3.6]=10.5 (spectral type later than about M0), the [3.6]−[4.5] color for M dwarfs departs
slightly from zero, becoming increasingly redder to the limit of the data (about M6).
3. A Catalog of Pleiades Candidate Members
If one limits oneself to only stars visible with the naked eye, it is easy to identify which
stars are members of the Pleiades – all of the stars within a degree of the cluster center that
have V < 6 are indeed members. However, if one were to try to identify the M dwarf stellar
members of the cluster (roughly 14 < V < 23), only of order 1% of the stars towards the
cluster center are likely to be members, and it is much harder to construct an uncontaminated
catalog. The problem is exacerbated by the fact that the Pleiades is old enough that mass
segregation through dynamical processes has occurred, and therefore one has to survey a
much larger region of the sky in order to include all of the M dwarf members.
The other primary difficulty in constructing a comprehensive member catalog for the
Pleiades is that the pedigree of the candidates varies greatly. For the best studied stars,
astrometric positions can be measured over temporal baselines ranging up to a century or
more, and the separation of cluster members from field stars in a vector point diagram
(VPD) can be extremely good. In addition, accurate radial velocities and other spectral
indicators are available for essentially all of the bright cluster members, and these further
allow membership assessment to be essentially definitive. Conversely, at the faint end (for
stars near the hydrogen burning mass limit in the Pleiades), members are near the detection
limit of the existing wide-field photographic plates, and the errors on the proper motions
become correspondingly large, causing the separation of cluster members from field stars
in the VPD to become poor. These stars are also sufficiently faint that spectra capable of
discriminating members from field dwarfs can only be obtained with 8m class telescopes, and
only a very small fraction of the faint candidates have had such spectra obtained. Therefore,
any comprehensive catalog created for the Pleiades will necessarily have stars ranging from
certain members to candidates for which very little is known, and where the fraction of
– 7 –
spurious candidate members increases to lower masses.
In order to address the membership uncertainties and biases, we have chosen a sliding
scale for inclusion in our catalog. For all stars, we require that the available photometry yields
location in color-color and color-magnitude diagrams consistent with cluster membership.
For the stars with well-calibrated photoelectric photometry, this means the star should not
fall below the Pleiades single-star locus by more than about 0.2 mag or above that locus
by more than about 1.0 mag (the expected displacement for a hierarchical triple with three
nearly equal mass components). For stars with only photographic optical photometry, where
the 1-σ uncertainties are of order 0.1 to 0.2 mag, we still require the star’s photometry to
be consistent with membership, but the allowed displacements from the single star locus are
considerably larger. Where accurate radial velocities are known, we require that the star
be considered a radial velocity member based on the paper where the radial velocities were
presented. Where stars have been previously identified as non-members based on photometric
or spectroscopic indices, we adopt those conclusions.
Two other relevant pieces of information are sometimes available. In some cases, in-
dividual proper motion membership probabilities are provided by the various membership
surveys. If no other information is available, and if the membership probability for a given
candidate is less than 0.1, we exclude that star from our final catalog. However, often a star
appears in several catalogs; if it appears in two or more proper motion membership lists we
include it in the final catalog even if P < 0.1 in one of those catalogs. Second, an entirely
different means to identify candidate Pleiades members is via flare star surveys towards the
cluster (Haro et al. 1982; Jones 1981). A star with a formally low membership probability
in one catalog but whose photometry is consistent with membership and that was identified
as a flare star is retained in our catalog.
Further details of the catalog construction are provided in the appendix, as are details of
the means by which the B, V , and I photometry have been homogenized. A full discussion
and listing of all of the papers from which we have extracted astrometric and photometric
information is also provided in the appendix. Here we simply provide a very brief description
of the inputs to the catalog.
We include candidate cluster members from the following proper motion surveys: Trumpler
(1921), Hertzsprung (1947), Jones (1981), Pels and Lub – as reported in van Leeuwen, Alphenaar, & Brand
(1986), Stauffer et al. (1991), Artyukhina (1969), Hambly et al. (1993), Pinfield et al. (2000),
Adams et al. (2001) and Deacon & Hambly (2004). Another important compilation which
provides the initial identification of a significant number of low mass cluster members is the
flare star catalog of Haro et al. (1982). Table 1 provides a brief synopsis of the characteristics
of the candidate member catalogs from these papers. The Trumpler paper is listed twice
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in Table 1 because there are two membership surveys included in that paper, with differing
spatial coverages and different limiting magnitudes.
In our final catalog, we have attempted to follow the standard naming convention
whereby the primary name is derived from the paper where it was first identified as a cluster
member. An exception to this arises for stars with both Trumpler (1921) and Hertzsprung
(1947) names, where we use the Hertzsprung numbers as the standard name because that
is the most commonly used designation for these stars in the literature. The failure for
the Trumpler numbers to be given precedence in the literature perhaps stems from the fact
that the Trumpler catalog was published in the Lick Observatory Bulletins as opposed to a
refereed journal. In addition to providing a primary name for each star, we provide cross-
identifications to some of the other catalogs, particularly where there is existing photometry
or spectroscopy of that star using the alternate names. For the brightest cluster members, we
provide additional cross-references (e.g., Greek names, Flamsteed numbers, HD numbers).
For each star, we attempt to include an estimate for Johnson B and V , and for Cousins
I (IC). Only a very small fraction of the cluster members have photoelectric photometry in
these systems, unfortunately. Photometry for many of the stars has often been obtained in
other systems, including Walraven, Geneva, Kron, and Johnson. We have used previously
published transformations from the appropriate indices in those systems to Johnson BV or
Cousins I. In other cases, photometry is available in a natural I band system, primarily
for some of the relatively faint cluster members. We have attempted to transform those I
band data to IC by deriving our own conversion using stars for which we already have a IC
estimate as well as the natural I measurement. Details of these issues are provided in the
Appendix.
Finally, we have cross-correlated the cluster candidates catalog with the 2MASS All-Sky
PSC and also with the 6x-PSC for the Pleiades. For every star in the catalog, we obtain
JHKs photometry and 2MASS positions. Where we have both main survey 2MASS data
and data from the 6x catalog, we adopt the 6x data for stars with J >11, and data from
the standard 2MASS catalog otherwise. We verified that the two catalogs do not have any
obvious photometric or astrometric offsets relative to each other. The coordinates we list in
our catalog are entirely from these 2MASS sources, and hence they inherit the very good
and homogeneous 2MASS positional accuracies of order 0.1 arcseconds RMS.
We have then plotted the candidate Pleiades members in a variety of color-magnitude
diagrams and color-color diagrams, and required that a star must have photometry that is
consistent with cluster membership. Figure 8 illustrates this process, and indicates why (for
example) we have excluded HII 1695 from our final catalog.
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Table 2 provides the collected data for the 1417 stars we have retained as candidate
Pleiades members. The first two columns are the J2000 RA and Dec from 2MASS; the next
are the 2MASS JHKs photometry and their uncertainties, and the 2MASS photometric
quality flag (“ph-qual”). If the number following the 2MASS quality flag is a 1, the 2MASS
data come from the 2MASS All-Sky PSC; if it is a 2, the data come from the 6x-PSC. The
next three columns provide the B, V and IC photometry, followed by a flag which indicates
the provenance of that photometry. The last column provides the most commonly used
names for these stars. The hydrogen burning mass limit for the Pleiades occurs at about
V=22, I=18, Ks=14.4. Fifty-three of the candidate members in the catalog are fainter than
this limit, and hence should be sub-stellar if they are indeed Pleiades members.
Table 3 provides the IRAC [3.6], [4.5], [5.8] and [8.0] photometry we have derived for
Pleiades candidate members included within the region covered by the IRAC shallow survey
of the Pleiades (see section 2). The brightest stars are saturated even in our short integration
frame data, particularly for the more sensitive 3.6 and 4.5 µm channels. At the faint end,
we provide photometry only for 3.6 and 4.5 µm because the objects are undetected in the
two longer wavelength channels. At the “top” and “bottom” of the survey region, we have
incomplete wavelength coverage for a band of width about 5′, and for stars in those areas
we report only photometry in either the 3.6 and 5.8 bands or in 4.5 and 8.0 bands.
Because Table 2 is an amalgam of many previous catalogs, each of which have different
spatial coverage, magnitude limits and other idiosyncrasies, it is necessarily incomplete and
inhomogeneous. It also certainly includes some non-members. For V < 12, we expect very
few non-members because of the extensive spectroscopic data available for those stars; the
fraction of non-members will likely increase to fainter magnitudes, particularly for stars
located far from the cluster center. The catalog is simply an attempt to collect all of the
available data, identify some of the non-members and eliminate duplications. We hope that
it will also serve as a starting point for future efforts to produce a “cleaner” catalog.
Figure 9 shows the distribution on the sky of the stars in Table 2. The complete
spatial distribution of all members of the Pleiades may differ slightly from what is shown
due to the inhomogeneous properties of the proper motion surveys. However, we believe
that those effects are relatively small and the distribution shown is mostly representative
of the parent population. One thing that is evident in Figure 9 is mass segregation – the
highest mass cluster members are much more centrally located than the lowest mass cluster
members. This fact is reinforced by calculating the cumulative number of stars as a function
of distance from the cluster center for different absolute magnitude bins. Figure 10 illustrates
this fact. Another property of the Pleiades illustrated by Figure 9 is that the cluster appears
to be elongated parallel to the galactic plane, as expected from n-body simulations of galactic
– 10 –
clusters (Terlevich 1987). Similar plots showing the flattening of the cluster and evidence for
mass segregation for the V < 12 cluster members were provided by (Raboud & Mermilliod
1998).
4. Empirical Pleiades Isochrones and Comparison to Model Isochrones
Young, nearby, rich open clusters like the Pleiades can and should be used to provide
template data which can help interpret observations of more distant clusters or to test
theoretical models. The identification of candidate members of distant open clusters is often
based on plots of stars in a color-magnitude diagram, overlaid upon which is a line meant
to define the single-star locus at the distance of the cluster. The stars lying near or slightly
above the locus are chosen as possible or probable cluster members. The data we have
collected for the Pleiades provide a means to define the single-star locus for 100 Myr, solar
metallicity stars in a variety of widely used color systems down to and slightly below the
hydrogen burning mass limit. Figure 11 and Figure 12 illustrate the appearance of the
Pleiades stars in two of these diagrams, and the single-star locus we have defined. The curve
defining the single-star locus was drawn entirely “by eye.” It is displaced slightly above the
lower envelope to the locus of stars to account for photometric uncertainties (which increase
to fainter magnitudes). We attempted to use all of the information available to us, however.
That is, there should also be an upper envelope to the Pleiades locus in these diagrams, since
equal mass binaries should be displaced above the single star sequence by 0.7 magnitudes
(and one expects very few systems of higher multiplicity). Therefore, the single star locus
was defined with that upper envelope in mind. Table 4 provides the single-star loci for the
Pleiades for BV IcJKs plus the four IRAC channels. We have dereddened the empirical loci
by the canonical mean extinction to the Pleiades of AV = 0.12 (and, correspondingly, AB
= 0.16, AI = 0.07, AJ = 0.03, AK = 0.01, as per the reddening law of Rieke & Lebofsky
(1985)).
The other benefit to constructing the new catalog is that it can provide an improved
comparison dataset to test theoretical isochrones. The new catalog provides homogeneous
photometry in many photometric bands for stars ranging from several solar masses down
to below 0.1 M⊙. We take the distance to the Pleiades as 133 pc, and refer the reader to
Soderblom et al. (2005) for a discussion and a listing of the most recent determinations. The
age of the Pleiades is not as well-defined, but is probably somewhere between 100 and 125
Myr (Meynet, Mermilliod, & Maeder 1993; Stauffer et al. 1998). We adopt 100 Myr for the
purposes of this discussion; our conclusions relative to the theoretical isochrones would not be
affected significantly if we instead chose 125 Myr. As noted above, we adopt AV=0.12 as the
– 11 –
mean Pleiades extinction, and apply that value to the theoretical isochrones. A small number
of Pleiades members have significantly larger extinctions (Breger 1986; Stauffer & Hartmann
1987), and we have dereddened those stars individually to the mean cluster reddening.
Figures 13 and 14 compare theoretical 100 Myr isochrones from Siess et al. (2000) and
Baraffe et al. (1998) to the Pleiades member photometry from Table 2 for stars for which
we have photoelectric photometry. Neither set of isochrones are a good fit to the V − I
based color-magnitude diagram. For Baraffe et al. (1998) this is not a surprise because they
illustrated that their isochrones are too blue in V−I for cool stars in their paper, and ascribed
the problem as likely the result of an incomplete line list, resulting in too little absorption in
the V band. For Siess et al. (2000), the poor fit in the V − I CMD is somewhat unexpected
in that they transform from the theoretical to the observational plane using empirical color-
temperature relations. In any event, it is clear that neither model isochrones match the
shape of the Pleiades locus in the V vs. V − I plane, and therefore use of these V − I based
isochrones for younger clusters is not likely to yield accurate results (unless the color-Teff
relation is recalibrated, as described for example in Jeffries & Oliveira (2005)). On the other
hand, the Baraffe et al. (1998) model provides a quite good fit to the Pleiades single star
locus for an age of 100 Myr in the K vs. I − K plane.2. This perhaps lends support to
the hypothesis that the misfit in the V vs. V − I plane is due to missing opacity in their
V band atmospheres for low mass stars (see also Chabrier et al. (2000) for further evidence
in support of this idea). The Siess et al. (2000) isochrones do not fit the Pleiades locus in
the K vs. I − K plane particularly well, being too faint near I − K=2 and too bright for
I −K > 2.5.
5. Identification of New Very Low Mass Candidate Members
The highest spatial density for Pleiades members of any mass should be at the cluster
center. However, searches for substellar members of the Pleiades have generally avoided
the cluster center because of the deleterious effects of scattered light from the high mass
cluster members and because of the variable background from the Pleiades reflection nebulae.
The deep 2MASS and IRAC 3.6 and 4.5 µm imaging provide accurate photometry to well
below the hydrogen burning mass limit, and are less affected by the nebular emission than
shorter wavelength images. We therefore expect that it should be possible to identify a new
2These isochrones are calculated for the standard K filter, rather than Ks. However, the difference in
location of the isochrones in these plots because of this should be very slight, and we do not believe our
conclusions are significantly affected.
– 12 –
set of candidate Pleiades substellar members by combining our new near and mid-infrared
photometry.
The substellar mass limit in the Pleiades occurs at about Ks=14.4, near the limit of
the 2MASS All-Sky PSC. As illustrated in Figure 2, the deep 2MASS survey of the Pleiades
should easily detect objects at least two magnitudes fainter than the substellar limit. The
key to actually identifying those objects and separating them from the background sources
is to find color-magnitude or color-color diagrams which separate the Pleiades members from
the other objects. As shown in Figure 15, late-type Pleiades members separate fairly well
from most field stars towards the Pleiades in a Ks vs. Ks − [3.6] color-magnitude diagram.
However, as illustrated in Figure 2, in the Ks magnitude range of interest there is also a
large population of red galaxies, and they are in fact the primary contaminants to identi-
fying Pleiades substellar objects in the Ks vs. Ks − [3.6] plane. Fortunately, most of the
contaminant galaxies are slightly resolved in the 2MASS and IRAC imaging, and we have
found that we can eliminate most of the red galaxies by their non-stellar image shape.
Figure 15 shows the first step in our process of identifying new very low mass members of
the Pleiades. The red plus symbols are the known Pleiades members from Table 2. The red
open circles are candidate Pleiades substellar members from deep imaging surveys published
in the literature, mostly of parts of the cluster exterior to the central square degree, where
the IRAC photometry is from Lowrance et al. (2007). The blue, filled circles are field M
and L dwarfs, placed at the distance of the Pleiades, using photometry from Patten et al.
(2006). Because the Pleiades is ∼100 Myr, its very low mass stellar and substellar objects
will be displaced about 0.7 mag above the locus of the field M and L dwarfs according to the
Baraffe et al. (1998) and Chabrier et al. (2000) models, in accord with the location in the
diagram of the previously identified, candidate VLM and substellar objects. The trapezoidal
shaped region outlined with a dashed line is the region in the diagram which we define as
containing candidate new VLM and substellar members of the Pleiades. We place the faint
limit of this region at Ks=16.2 in order to avoid the large apparent increase in faint, red
objects for Ks> 16.2, caused largely by increasing errors in the Ks photometry. Also, the
2MASS extended object flags cease to be useful fainter than about Ks= 16.
We took the following steps to identify a set of candidate substellar members of the
Pleiades:
• keep only objects which fall in the trapezoidal region in Figure 15.
• remove objects flagged as non-stellar by the 2MASS pipeline software;
• remove objects which appear non-stellar to the eye in the IRAC images;
– 13 –
• remove objects which do not fall in or near the locus of field M and L dwarfs in a J−H
vs. H −Ks diagram;
• remove objects which have 3.6 and 4.5 µm magnitudes that differ by more than 0.2
• remove objects which fall below the ZAMS in a J vs. J −Ks diagram.
As shown in Figure 15, all stars earlier than about mid-M have Ks − [3.6] colors bluer
than 0.4. This ensures that for most of the area of the trapezoidal region, the primary
contaminants are distant galaxies. Fortunately, the 2MASS catalog provides two types of
flags for identifying extended objects. For each filter, a chi-square flag measures the match
between the objects shape and the instrumental PSF, with values greater than 2.0 generally
indicative of a non-stellar object. In order not to be misguided by an image artifact in one
filter, we throw out the most discrepant of the three flags and average the other two. We
discard objects with mean χ2 greater than 1.9. The other indicator is the 2MASS extended
object flag, which is the synthesis of several independent tests of the objects shape, surface
brightness and color (see Jarrett, T. et al (2000) for a description of this process). If one
simply excludes the objects classified as extended in the 2MASS 6x image by either of these
techniques, the number of candidate VLM and substellar objects lying inside the trapezoidal
region decreases by nearly a half.
We have one additional means to demonstrate that many of the identified objects are
probably Pleiades members, and that is via proper motions. The mean Pleiades proper
motion is ∆RA = 20 mas yr−1 and ∆Dec = −45 mas yr−1 (Jones 1973). With an epoch
difference of only 3.5 years between the deep 2MASS and IRAC imaging, the expected motion
for a Pleiades member is only 0.07 arcseconds in RA and −0.16 arcseconds in Dec. Given
the relatively large pixel size for the two cameras, and the undersampled nature of the IRAC
3.6 and 4.5 µm images, it is not a priori obvious that one would expect to reliably detect the
Pleiades motion. However, both the 2MASS and IRAC astrometric solutions have been very
accurately calibrated. Also, for the present purpose, we only ask whether the data support a
conclusion that most of the identified substellar candidates are true Pleiades members (i.e.,
as an ensemble), rather than that each star is well enough separated in a VPD to derive a
high membership probability.
Figure 16 provides a set of plots that we believe support the conclusion that the majority
of the surviving VLM and substellar candidates are Pleiades members. The first plot shows
the measured motions between the epoch of the 2MASS and IRAC observations for all known
Pleiades members from Table 2 that lie in the central square degree region and have 11 <
Ks < 14 (i.e., just brighter than the substellar candidates). The mean offset of the Pleiades
– 14 –
stellar members from the background population is well-defined and is quantitatively of the
expected magnitude and sign (+0.07 arcsec in RA and −0.16 arcsec in Dec). The RMS
dispersion of the coordinate difference for the field population in RA and Dec is 0.076 and
0.062 arcseconds, supportive of our claim that the relative astrometry for the two cameras
is quite good. Because we expect that the background population should have essentially
no mean proper motion, the non-zero mean “motion” of the field population of about <
∆RA>=0.3 arcseconds is presumably not real. Instead, the offset is probably due to the
uncertainty in transferring the Spitzer coordinate zero-point between the warm star-tracker
and the cryogenic focal plane. Because it is simply a zero-point offset applicable to all the
objects in the IRAC catalog, it has no effect on the ability to separate Pleiades members
from the field star population.
The second panel in Figure 16 shows the proper motion of the candidate Pleiades VLM
and substellar objects. While these objects do not show as clean a distribution as the known
members, their mean motion is clearly in the same direction. After removing 2-σ deviants,
the median offsets for the substellar candidates are 0.04 and −0.11 arcseconds in RA and
Dec, respectively. The objects whose motions differ significantly from the Pleiades mean
may be non-members or they may be members with poorly determined motions (since a few
of the high probability members in the first panel also show discrepant motions).
The other two panels in Figure 16 show the proper motions of two possible control
samples. The first control sample was defined as the set of stars that fall up to 0.3 magnitudes
below the lower sloping boundary of the trapezoid in Figure 15. These objects should be late
type dwarfs that are either older or more distant than the Pleiades or red galaxies. We used
the 2MASS data to remove extended or blended objects from the sample in the same way as
for the Pleiades candidates. If the objects are nearby field stars, we expect to see large proper
motions; if galaxies, the real proper motions would be small – but relatively large apparent
proper motions due to poor centroiding or different centroids at different effective wavelengths
could be present. The second control set was defined to have −0.1 < K − [3.6] < 0.1 and
14.0 < K < 14.5, and to be stellar based on the 2MASS flags. This control sample should
therefore be relatively distant G and K dwarfs primarily. Both control samples have proper
motion distributions that differ greatly from the Pleiades samples and that make sense for,
respectively, a nearby and a distant field star sample.
Figure 17 shows the Pleiades members from Table 2 and the 55 candidate VLM and
substellar members that survived all of our culling steps. We cross-correlated this list with the
stars from Table 2 and with a list of the previously identified candidate substellar members
of the cluster from other deep imaging surveys. Fourteen of the surviving objects correspond
to previously identified Pleiades VLM and substellar candidates. We provide the new list
– 15 –
of candidate members in Table 5. The columns marked as µ(RA) and µ(DEC) are the
measured motions, in arcsec over the 3.5 year epoch difference between the 2MASS-6x and
IRAC observations. Forty-two of these objects have Ks> 14.0, and hence inferred masses
less than about 0.1 M⊙; thirty-one of them have Ks> 14.4, and hence have inferred masses
below the hydrogen burning mass limit.
Our candidate list could be contaminated by foreground late type dwarfs that happen
to lie in the line of sight to the Pleiades. How many such objects should we expect? In
order to pass our culling steps, such stars would have to be mid to late M dwarfs, or early
to mid L dwarfs. We use the known M dwarfs within 8 pc to estimate how many field M
dwarfs should lie in a one square degree region and at distance between 70 and 100 parsecs
(so they would be coincident in a CMD with the 100 Myr Pleiades members). The result is
∼3 such field M dwarf contaminants. Cruz et al. (2006) estimate that the volume density of
L dwarfs is comparable to that for late-M dwarfs, and therefore a very conservative estimate
is that there might also be 3 field L dwarfs contaminating our sample. We regard this (6
contaminating field dwarfs) as an upper limit because our various selection criteria would
exclude early M dwarfs and late L dwarfs. Bihain et al. (2006) made an estimate of the
number of contaminating field dwarfs in their Pleiades survey of 1.8 square degrees; for the
spectral type range of our objects, their algorithm would have predicted just one or two
contaminating field dwarfs for our survey.
How many substellar Pleiades members should there be in the region we have surveyed?
That is, of course, part of the question we are trying to answer. However, previous studies
have estimated that the Pleiades stellar mass function for M < 0.5 M⊙ can be approximated
as a power-law with an exponent of -1 (dN/dM ∝ M−1). Using the known Pleiades members
from Table 2 that lie within the region of the IRAC survey and that have masses of 0.2 <
M/M⊙< 0.5 (as estimated from the Baraffe et al. (1998) 100 Myr isochrone) to normalize
the relation, the M−1 mass function predicts about 48 members in our search region and
with 14 < K < 16.2 (corresponding to 0.1 < M/M⊙< 0.035). Other studies have suggested
that the mass function in the Pleiades becomes shallower below 0.1 M⊙, dN/dM ∝ M
−0.6.
Using the same normalization as above, this functional form for the Pleiades mass function
for M < 0.1 M⊙ yields a prediction of 20 VLM and substellar members in our survey. The
number of candidates we have found falls between these two estimates. Better proper motions
and low-resolution spectroscopy will almost certaintly eliminate some of these candidates as
non-members.
– 16 –
6. Mid-IR Observations of Dust and PAHS in the Pleiades
Since the earliest days of astrophotography, it has been clear that the Pleiades stars
are in relatively close proximity to interstellar matter whose optical manifestation is the
spider-web like network of filaments seen particularly strongly towards several of the B stars
in the cluster. High resolution spectra of the brightest Pleiades stars as well as CO maps
towards the cluster show that there is gas as well as dust present, and that the (primary)
interstellar cloud has a significant radial velocity offset relative to the Pleiades (White 2003;
Federman & Willson 1984). The gas and dust, therefore, are not a remnant from the forma-
tion of the cluster but are simply evidence of a a transitory event as this small cloud passes by
the cluster in our line of sight (see also Breger (1986)). There are at least two claimed mor-
phological signatures of a direct interaction of the Pleiades with the cloud. White & Bally
(1993) provided evidence that the IRAS 60 and 100 µm image of the vicinity of the Pleiades
showed a dark channel immediately to the east of the Pleiades, which they interpreted as
the “wake” of the Pleiades as it plowed through the cloud from the east. Herbig & Simon
(2001) provided a detailed analysis of the optically brightest nebular feature in the Pleiades
– IC 349 (Barnard’s Merope nebula) – and concluded that the shape and structure of that
nebula could best be understood if the cloud was running into the Pleiades from the south-
east. Herbig & Simon (2001) concluded that the IC 349 cloudlet, and by extension the rest
of the gas and dust enveloping the Pleiades, are relatively distant outliers of the Taurus
molecular clouds (see also Eggen (1950) for a much earlier discussion ascribing the Merope
nebulae as outliers of the Taurus clouds). White (2003) has more recently proposed a hybrid
model, where there are two separate interstellar cloud complexes with very different space
motions, both of which are colliding simultaneously with the Pleiades and with each other.
Breger (1986) provided polarization measurements for a sample of member and back-
ground stars towards the Pleiades, and argued that the variation in polarization signatures
across the face of the cluster was evidence that some of the gas and dust was within the clus-
ter. In particular, Figure 6 of that paper showed a fairly distinct interface region, with little
residual polarization to the NE portion of the cluster and an L-shaped boundary running
EW along the southern edge of the cluster and then north-south along the western edge of
the cluster. Stars to the south and west of that boundary show relatively large polarizations
and consistent angles (see also our Figure 5 where we provide a few polarization vectors from
Breger (1986) to illustrate the location of the interface region and the fact that the position
angle of the polarization correlates well with the location in the interface).
There is a general correspondence between the polarization map and what is seen with
IRAC, in the sense that the B stars in the NE portion of the cluster (Atlas and Alcyone)
have little nebular emission in their vicinity, whereas those in the western part of the cluster
– 17 –
(Maia, Electra and Asterope) have prominent, filamentary dust emission in their vicinity.
The L-shaped boundary is in fact visible in Figure 4 as enhanced nebular emission running
between and below a line roughly joining Merope and Electra, and then making a right
angle and running roughly parallel to a line running from Electra to Maia to HII1234 (see
Figure 5).
6.1. Pleiades Dust-Star Encounters Imaged with IRAC
The Pleiades dust filaments are most strongly evident in IRAC’s 8 µm channel, as
evidenced by the distinct red color of the nebular features in Figure 4. The dominance at 8
µm is an expected feature of reflection nebulae, as exemplified by NGC 7023 (Werner et al.
2004), where most of the mid-infrared emission arises from polycyclic aromatic hydrocarbons
(PAHs) whose strongest bands in the 3 to 10 µm region fall at 7.7 and 8.6 µm. One might
expect that if portions of the passing cloud were particularly near to one of the Pleiades
members, it might be possible to identify such interactions by searching for stars with 8.0
µm excesses or for stars with extended emission at 8 µm. Figure 18 provides two such plots.
Four stars stand out as having significant extended 8 µm emission, with two of those stars
also having an 8 µm excess based on their [3.6]−[8.0] color. All of these stars, plus IC 349,
are located approximately along the interface region identified by Breger (1986).
We have subtracted a PSF from the 8 µm images for the stars with extended emission,
and those PSF-subtracted images are provided in Figure 19. The image for HII 1234 has
the appearance of a bow-shock. The shape is reminiscent of predictions for what one should
expect from a collision between a large cloud or a sheet of gas and an A star as described
in Artymowicz & Clampin (1997). The Artymowicz & Clampin (1997) model posits that
A stars encountering a cloud will carve a paraboloidal shaped cavity in the cloud via radi-
ation pressure. The exact size and shape of the cavity depend on the relative velocity of
the encounter, the star’s mass and luminosity and properties of the ISM grains. For typical
parameters, the predicted characteristic size of the cavity is of order 1000 AU, quite compa-
rable to the size of the structures around HII 652 and HII 1234. The observed appearance
of the cavity depends on the view angle to the observer. However, in any case, the direction
from which the gas is moving relative to the star can be inferred from the location of the
star relative to the curved rim of the cavity; the “wind” originates approximately from the
direction connecting the star and the apex of the rim. For HII 1234, this indicates the cloud
which it is encountering has a motion relative to HII 1234 from the SSE, in accord with a
Taurus origin and not in accord for where a cloud is impacting the Pleiades from the west
as posited in White (2003). The nebular emission for HII 652 is less strongly bow-shaped,
– 18 –
but the peak of the excess emission is displaced roughly southward from the star, consistent
with the Taurus model and inconsistent with gas flowing from the west.
Despite being the brightest part of the Pleiades nebulae in the optical, IC 349 appears
to be undetected in the 8 µm image. This is not because the 8 µm image is insensitive to
the nebular emission - there is generally good agreement between the structures seen in the
optical and at 8 µm, and most of the filaments present in optical images of the Pleiades are
also visible on the 8 µm image (see Figures 4 and 19) and even the psf-subtracted image of
Merope shows well-defined nebular filaments. The lack of enhanced 8 µm emission from the
region of IC 349 is probably because all of the small particles have been scoured away from
this cloudlet, consistent with Herbig’s model to explain the HST surface photometry and
colors. There is no PAH emission from IC 349 because there are none of the small molecules
that are the postulated source of the PAH emission.
IC349 is very bright in the optical, and undetected to a good sensitivity limit at 8 µm; it
must be detectable via imaging at some wavelength between 5000 Å and 8 µm. We checked
our 3.6 µm data for this purpose. In the standard BCD mosaic image, we were unable to
discern an excess at the location of IC349 either simply by displaying the image with various
stretches or by doing cuts through the image. We performed a PSF subtraction of Merope
from the image in order to attempt to improve our ability to detect faint, extended emission
30” from Merope - unfortunately, bright stars have ghost images in IRAC Ch. 1, and in this
case the ghost image falls almost exactly at the location of IC349. IC349 is also not detected
in visual inspection of our 2MASS 6x images.
6.2. Circumstellar Disks and IRAC
As part of the Spitzer FEPS (Formation and Evolution of Planetary Systems) Legacy
program, using pointed MIPS photometry, Stauffer et al. (2005) identified three G dwarfs
in the Pleiades as having 24 µm excesses probably indicative of circumstellar dust disks.
Gorlova et al. (2006) reported results of a MIPS GTO survey of the Pleiades, and identified
nine cluster members that appear to have 24 µm excesses due to circumstellar disks. However,
it is possible that in a few cases these apparent excesses could be due instead to a knot of
the passing interstellar dust impacting the cluster member, or that the 24 µm excess could
be flux from a background galaxy projected onto the line of sight to the Pleiades member.
Careful analysis of the IRAC images of these cluster members may help confirm that the
MIPS excesses are evidence for debris disks rather than the other possible explanations.
Six of the Pleiades members with probable 24 µm excesses are included in the region
– 19 –
mapped with IRAC. However, only four of them have data at 8 µm – the other two fall
near the edge of the mapped region and only have data at 3.6 and 5.8 µm. None of the
six stars appear to have significant local nebular dust from visual inspection of the IRAC
mosaic images. Also, none of them appear problematic in Figure 18. For a slightly more
quantitative analysis of possible nebular contamination, we also constructed aperture growth
curves for the six stars, and compared them to other Pleiades members. All but one of the
six show aperture growth curves that are normal and consistent with the expected IRAC
PSF. The one exception is HII 489, which has a slight excess at large aperture sizes as is
illustrated in Figure 20. Because HII 489 only has a small 24 µm excess, it is possible that
the 24 µm excess is due to a local knot of the interstellar cloud material and is not due to a
debris disk. For the other five 24 µm excess stars we find no such problem, and we conclude
that their 24 µm excesses are indeed best explained as due to debris disks.
7. Summary and Conclusions
We have collated the primary membership catalogs for the Pleiades to produce the first
catalog of the cluster extending from its highest mass members to the substellar limit. At the
bright end, we expect this catalog to be essentially complete and with few or no non-member
contaminants. At the faint end, the data establishing membership are much sparser, and we
expect a significant number of objects will be non-members. We hope that the creation of
this catalog will spur efforts to obtain accurate radial velocities and proper motions for the
faint candidate members in order to eventually provide a well-vetted membership catalog
for the stellar members of the Pleiades. Towards that end, it would be useful to update
the current catalog with other data – such as radial velocities, lithium equivalent widths,
x-ray fluxes, Hα equivalent widths, etc. – which could be used to help accurately establish
membership for the low mass cluster candidates. It is also possible to make more use of
“negative information” present in the proper motion catalogs. That is, if a member from
one catalog is not included in another study but does fall within its areal and luminosity
coverage, that suggests that it likely failed the membership criteria of the second study. For
a few individual stars, we have done this type of comparison, but a systematic analysis of
the proper motion catalogs should be conducted. We intend to undertake these tasks, and
plan to establish a website where these data would be hosted.
We have used the new Pleiades member catalog to define the single-star locus at 100
Myr for BV IcKs and the four IRAC bands. These curves can be used as empirical calibration
curves when attempting to identify members of less well-studied, more distant clusters of
similar age. We compared the Pleiades photometry to theoretical isochrones from Siess et al.
– 20 –
(2000) and Baraffe et al. (1998). The Siess et al. (2000) isochrones are not, in detail, a good
fit to the Pleiades photometry, particularly for low mass stars. The Baraffe et al. (1998) 100
Myr isochrone does fit the Pleiades photometry very well in the I vs. I −K plane.
We have identified 31 new substellar candidate members of the Pleiades using our com-
bined seven-band infrared photometry, and have shown that the majority of these objects
appear to share the Pleiades proper motion. We believe that most of the objects that may
be contaminating our list of candidate brown dwarfs are likely to be unresolved galaxies, and
therefore low resolution spectroscopy should be able to provide a good criterion for culling
our list of non-members.
The IRAC images, particularly the 8 µm mosaic, provide vivid evidence of the strong in-
teraction of the Pleiades stars and the interstellar cloud that is passing through the Pleiades.
Our data are supportive of the model proposed by Herbig & Simon (2001) whereby the pass-
ing cloud is part of the Taurus cloud complex and hence is encountering the Pleiades from
the SSE direction. White & Bally (1993) had proposed a model whereby the cloud was
encountering the Pleiades from the west and used this to explain features in the IRAS 60
and 100 µm images of the region as the wake of the Pleiades moving through the cloud.
Our data appear to not be supportive of that hypothesis, and therefore leaves the apparent
structure in the IRAS maps as unexplained.
Most of the support for this work was provided by the Jet Propulsion Laboratory, Cal-
ifornia Institute of Technology, under NASA contract 1407. This research has made use of
NASA’s Astrophysics Data System (ADS) Abstract Service, and of the SIMBAD database,
operated at CDS, Strasbourg, France. This research has made use of data products from
the Two Micron All-Sky Survey (2MASS), which is a joint project of the University of Mas-
sachusetts and the Infrared Processing and Analysis Center, funded by the National Aero-
nautics and Space Administration and the National Science Foundation. These data were
served by the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion
Laboratory, California Institute of Technology, under contract with the National Aeronautics
and Space Administration. The research described in this paper was partially carried out at
the Jet Propulsion Laboratory, California Institute of Technology, under contract with the
National Aeronautics and Space Administration.
– 21 –
A. APPENDIX
A.1. Membership Catalogs
Membership lists of the Pleiades date back to antiquity if one includes historical and
literary references to the Seven Sisters (Alcyone, Maia, Merope, Electra, Taygeta, Asterope
and Celeno) and their parents (Atlas and Pleione). The first paper discussing relative proper
motions of a large sample of stars in the Pleiades (based on visual observations) was published
by Pritchard (1884). The best of the early proper motion surveys of the Pleiades derived
from photographic plate astrometry was that by Trumpler (1921), based on plates obtained
at Yerkes and Lick observatories. The candidate members from that survey were presented
in two tables, with the first being devoted to candidate members within about one degree
from the cluster center (operationally, within one degree from Alcyone) and the second table
being devoted to candidates further than one degree from the cluster center. Most of the
latter stars were denoted by Trumpler by an S or R, followed by an identification number.
We use Tr to designate the Trumpler stars (hence Trnnn for a star from the 1st table and
the small number of stars in the second table without an “S” or an “R”, and TrSnnn or
TrRnnn for the other stars). For the central region, Trumpler’s catalog extends to V ∼ 13,
while the outer region catalog includes stars only to about V ∼ 9.
The most heavily referenced proper motion catalog of the Pleiades is that provided by
Hertzsprung (1947). That paper makes reference to two separate catalogs: a photometric
catalog of the Pleiades published by Hertzsprung in 1923 (Hertzsprung 1923), whose members
are commonly referred to by HI numbers, and the new proper motion catalog from the
1947 paper, commonly referenced as the HII catalog. While both HI and HII numbers
have been used in subsequent observational papers, it is the HII identification numbers
that predominate. That catalog – derived from Carte du Ciel blue-sensitive plates from
14 observatories – includes stars in the central 2×2 degree region of the cluster, and has
a faint limit of about V = 15.5. Johnson system BV I photometry is provided for most
of the proposed Hertzsprung members in Johnson & Mitchell (1958) and Iriarte (1967).
Additional Johnson B and V photometry plus Kron I photometry for a fairly large number
of the Hertzsprung members can be found in Stauffer (1980), Stauffer (1982), and Stauffer
(1984). Other Johnson BV photometry for a scattering of stars can be found in Jones (1973),
Robinson & Kraft (1974), Messina (2001). Spectroscopic confirmation, primarily via radial
velocities, that these are indeed Pleiades members has been provided in Soderblom et al.
(1993); Queloz et al. (1998) and Mermilliod et al. (1997).
Two other proper motion surveys provide relatively bright candidate members relatively
far from the cluster center: Artyukhina & Kalinina (1970) and van Leeuwen, Alphenaar, & Brand
– 22 –
(1986). Stars from the Artyukhina catalog are designated as AK followed by the region from
which the star was identified followed by an identification number. The new members pro-
vided in the van Leeuwen paper were taken from an otherwise unpublished proper motion
study by Pels, where the first 118 stars were considered probable members and the remaining
75 stars were considered possible members. Van Leeuwen categorized a number of the Pels
stars as non-members based on the Walraven photometry they obtained, and we adopt those
findings. Radial velocities for stars in these two catalogs have been obtained by Rosvick et al
(1992), Mermilliod et al. (1997), and Queloz et al. (1998), and those authors identified a list
of the candidate members that they considered confirmed by the high resolution spectroscopy.
For these outlying candidate members, to be included in Table 2 we require that the star
be a radial velocity member from one of the above three surveys, or be indicated as having
“no dip” in the Coravel cross-correlation (indicating rapid rotation, which at least for the
later type stars is suggestive of membership). Geneva photometry of the Artyukhina stars
considered as likely members was provided by Mermilliod et al. (1997). The magnitude limit
of these surveys was not well-defined, but most of the Artyukhina and Pels stars are brighter
than V=13.
Jones (1973) provided proper motion membership probabilities for a large sample of
proposed Pleiades members, and for a set of faint, red stars towards the Pleiades. A few
star identification names from the sources considered by Jones appear in Table 2, including
MT (McCarthy & Treanor 1964), VM (van Maanen 1946), ALR (Ahmed et al. 1965), and
J (Jones 1973).
The chronologically next significant source of new Pleiades candidate members was
the flare star survey of the Pleiades conducted at several observatories in the 1960s, and
summarized in Haro et al. (1982), hereafter HCG. The logic behind these surveys was that
even at 100 Myr, late type dwarfs have relatively frequent and relatively high luminosity
flares (as demonstrated by Johnson & Mitchell (1958) having detected two flares during
their photometric observations of the Pleiades), and therefore wide area, rapid cadence
imaging of the Pleiades at blue wavelengths should be capable of identifying low mass cluster
members. However, such surveys also will detect relatively young field dwarfs, and therefore
it is best to combine the flare star surveys with proper motions. Dedicated proper motion
surveys of the HCG flare stars were conducted by Jones (1981) and Stauffer et al. (1991),
with the latter also providing photographic V I photometry (Kron system). Photoelectric
photometry for some of the HCG stars have been reported in Stauffer (1982), Stauffer (1984),
Stauffer & Hartmann (1987), and Prosser et al. (1991). High resolution spectroscopy of
many of the HCG stars is reported in Stauffer (1984), Stauffer & Hartmann (1987) and
Terndrup et al. (2000). Because a number of the papers providing additional observational
data for the flare stars were obtained prior to 1982, we also include in Table 2 the original
– 23 –
flare star names which were derived from the observatory where the initial flare was detected.
Those names are of the form an initial letter indicating the observatory – A (Asiago), B
(Byurakan), K (Konkoly), T (Tonantzintla) – followed by an identification number.
Stauffer et al. (1991) conducted two proper motion surveys of the Pleiades over an
approximately 4×4 degree region of the cluster based on plates obtained with the Lick 20′′
astrographic telescope. The first survey was essentially unbiased, except for the requirement
that the stars fall approximately in the region of the V vs. V − I color-magnitude diagram
where Pleiades members should lie. Candidate members from this survey are designated
by SK numbers. The second survey was a proper motion survey of the HCG stars. Photo-
graphic V I photometry of all the stars was provided as well as proper motion membership
probabilities. Photoelectric photometry for some of the candidate members was obtained as
detailed above in the section on the HCG catalog stars. The faint limit of these surveys is
about V=18.
Hambly et al. (1991) provided a significantly deeper, somewhat wider area proper mo-
tion survey, with the faintest members having V ≃ 20 and the total area covered being
of order 25 square degrees. The survey utilized red sensitive plates from the Palomar and
UK Schmidt telescopes. Due to incomplete coverage at one epoch, there is a vertical swath
slightly east of the cluster center where no membership information is available. Stars from
this survey are designated by their HHJ numbers . Hambly et al. (1993) provide RI photo-
graphic photometry on a natural system for all of their candidate members, plus photoelectric
Cousins RI photometry for a small number of stars and JHK photometry for a larger sam-
ple. Some spectroscopy to confirm membership has been reported in Stauffer et al. (1994),
Stauffer et al. (1995), Oppenheimer et al. (1997), Stauffer et al. (1998), and Steele et al.
(1995), though for most of the HHJ stars there is no spectroscopic membership confirma-
tion.
Pinfield et al. (2000) provide the deepest wide-field proper motion survey of the Pleiades.
That survey combines CCD imaging of six square degrees of the Pleiades obtained with the
Burrell Schmidt telescope (as five separate, non-overlapping fields near but outside the cluster
center) with deep photographic plates which provide the 1st epoch positions. Candidate
members are designated by BPL numbers (for Burrell Pleiades), with the faintest stars
having I ≃ 19.5, corresponding to V > 23. Only the stars brighter than about I= 17 have
sufficiently accurate proper motions to use to identify Pleiades members. Fainter than I=
17, the primary selection criteria are that the star fall in an appropriate place in both an I
vs. I − Z and an I vs. I −K CMD.
Adams et al. (2001) combined the 2MASS and digitized POSS databases to produce a
very wide area proper motion survey of the Pleiades. By design, that survey was very inclu-
– 24 –
sive - covering the entire physical area of the cluster and extending to the hydrogen burning
mass limit. However, it was also very “contaminated”, with many suspected non-members.
The catalog of possible members was not published. We have therefore not included stars
from this study in Table 2; we have used the proper motion data from Adams et al. (2001)
to help decide cases where a given star has ambiguous membership data from the other
surveys.
Deacon & Hambly (2004) provided another deep and very wide area proper motion
survey of the Pleiades. The survey covers a circular area of approximately five degrees
radius to R ∼ 20, or V ∼ 22. Candidate members are designated by DH. Deacon & Hambly
(2004) also provide membership probabilities based on proper motions for many candidate
cluster members from previous surveys. For stars where Deacon & Hambly (2004) derive
P < 0.1 and where we have no other proper motion information or where another proper
motion survey also finds low membership probability, we exclude the star from our catalog.
For cases where two of our proper motion catalogs differ significantly in their membership
assessment, with one survey indicating the star is a probable member, we retain the star
in the catalog as the conservative choice. Examples of the latter where Deacon & Hambly
(2004) derive P < 0.1 include HII 1553, HII 2147, HII 2278 and HII 2665 – all of which
we retain in our catalog because other surveys indicate these are high probability Pleiades
members.
A.2. Photometry
Photometry for stars in open cluster catalogs can be used to help confirm cluster mem-
bership and to help constrain physical properties of those stars or of the cluster. For a
variety of reasons, photometry of stars in the Pleiades has been obtained in a panoply of
different photometric systems. For our own goals, which are to use the photometry to help
verify membership and to define the Pleiades single-star locus in color magnitude diagrams,
we have attempted to convert photometry in several of these systems to a common sys-
tem (Johnson BV and Cousins I). We detail below the sources of the photometry and the
conversions we have employed.
Photoelectric photometry of Pleiades members dates back to at least 1921 (Cummings
1921). However, as far as we are aware the first “modern” photoelectric photometry for the
Pleiades, using a potassium hydride photoelectric cell, is that of Calder & Shapley (1937).
Eggen (1950) provided photoelectric photometry using a 1P21 phototube (but calibrated
to a no-longer-used photographic system) for most of the known Pleiades members within
one degree of the cluster center and with magnitudes < 11. The first phototube photom-
– 25 –
etry of Pleiades stars calibrated more-or-less to the modern UBV system was provided by
Johnson & Morgan (1951). An update of that paper, and the oldest photometry included
here was reported in Johnson & Mitchell (1958), which provided UBV Johnson system pho-
tometry for a large sample of HII and Trumpler candidate Pleiades members. Iriarte (1967)
later reported Johnson system V − I colors for most of these stars. We have converted
Iriarte’s V − I photometry to estimated Cousins V − I colors using a formula from Bessell
(1979):
V − I(Cousins) = 0.778× V − I(Johnson). (A1)
BV RI photometry for most of the Hertzsprung members fainter than V= 10 has been
published by Stauffer (1980), Stauffer (1982), Stauffer (1984), and Stauffer & Hartmann
(1987). The BV photometry is Johnson system, whereas the RI photometry is on the Kron
system. The Kron V − I colors were converted to Cousins V − I using a transformation
provided by Bessell & Weis (1987):
V − I(Cousins) = 0.227 + 0.9567(V − I)k + 0.0128(V − I)
k − 0.0053(V − I)
k (A2)
Other Kron system V−I colors have been published for Pleiades candidates in Stauffer et al.
(1991) (photographic photometry) and in Prosser et al. (1991). These Kron-system colors
have also been converted to Cousins V − I using the above formula.
Johnson/Cousins UBV R photometry for a set of low mass Pleiades members was pro-
vided by Landolt (1979). We only use the BV magnitudes from that study. Additional John-
son system UBV photometry for small numbers of stars is provided in Robinson & Kraft
(1974), Messina (2001) and Jones (1973).
van Leeuwen, Alphenaar, & Meys (1987) provided Walraven V BLUW photometry for
nearly all of the Hertzsprung members brighter than V ∼ 13.5 and for the Pels candidate
members. Van Leeuwen provided an estimated Johnson V derived from the Walraven V
in his tables. We have transformed the Walraven V − B color into an estimate of Johnson
B − V using a formula from Rosvick et al (1992):
B − V (Johnson) = 2.571(V − B)− 1.02(V −B)2 + 0.5(V − B)3 − 0.01 (A3)
Hambly et al. (1993) provided photographic V RI photometry for all of the HHJ candidate
members, and V RI Cousins photoelectric photometry for a small fraction of those stars. We
took all of the HHJ stars with photographic photometry for which we also have photoelectric
V I photometry on the Cousins system, and plotted V (Cousins) vs. V (HHJ) and I(Cousins)
vs. I(HHJ). While there is some evidence for slight systematic departures of the HHJ photo-
graphic photometry from the Cousins system, those departures are relatively small and we
have chosen simply to retain the HHJ values and treat them as Cousins system.
– 26 –
Pinfield et al. (2000) reported their I magnitudes in an instrumental system which they
designated as Ikp. We identified all BPL candidate members for which we had photoelectric
Cousins I estimates, and plotted Ikp vs. IC. Figure 21 shows this correlation, and the
piecewise linear fit we have made to convert from Ikp to IC. Our catalog lists these converted
IC measures for the BPL stars for which we have no other photoelectric I estimates.
Deacon & Hambly (2004) derived RI photometry from the scans of their plates, and
calibrated that photometry by reference to published photometry from the literature. When
we plotted their the difference between their I band photometry and literature values (where
available), we discovered a significant dependence on right ascension. Unfortunately, because
the DH survey extended over larger spatial scales than the calibrating photometry, we could
not derive a correction which we could apply to all the DH stars. We therefore developed
the following indirect scheme. We used the stars for which we have estimated IC magnitudes
(from photoelectric photometry) to define the relation between J and (IC−J) for Pleiades
members. For each DH star, we combined that relation and the 2MASS J magnitude to yield
a predicted IC. Figure 22 shows a plot of the difference of this predicted IC and I(DH) with
right ascension. The solid line shows the relation we adopt. Figure 23 shows the relation
between the corrected I(DH) values and Table 2 IC measures from photoelectric sources.
There is still a significant amount of scatter but the corrected I(DH) photometry appears
to be accurately calibrated to the Cousins system.
In a very few cases (specifically, just five stars), we provide an estimate of Ic based
on data from a wide-area CCD survey of Taurus obtained with the Quest-2 camera on the
Palomar 48 inch Samuel Oschin telescope (Slesnick et al. 2006). That survey calibrated their
photometry to the Sloan i system, and we have converted the Sloan i magnitudes to Ic. We
intend to make more complete use of the Quest-2 data in a subsequent paper.
When we have multiple sources of photometry for a given star, we consider how to com-
bine them. In most cases, if we have photoelectric data, that is given preference. However,
if we have photographic V and I, and only a photoelectric measurement for I, we do not
replace the photographic I with the photoelectric value because these stars are variable and
the photographic measurements are at least in some cases from nearly simultaneous expo-
sures. Where we have multiple sources for photoelectric photometry, and no strong reason
to favor one measurement or set of measurements over another, we have averaged the pho-
tometry for a given star. In most cases, where we have multiple photometry the individual
measurements agree reasonably well but with the caveat that the Pleiades low mass stars are
in many cases heavily spotted and “active” chromospherically, and hence are photometrically
variable. In a few cases, even given the expectation that spots and other phenomena may
affect the photometry, there seems to be more discrepancy between reported V magnitudes
– 27 –
than we expect. We note two such cases here. We suspect these results indicate that at least
some of the Pleiades low mass stars have long-term photometric variability larger than their
short period (rotational) modulation.
HII 882 has at least four presumably accurate V magnitude measurements reported in
the literature. Those measures are: V=12.66 Johnson & Mitchell (1958); V=12.95 Stauffer
(1982); V=12.898 van Leeuwen, Alphenaar, & Brand (1986); and V=12.62 Messina (2001).
HII 345 has at least three presumably accurate V magnitude measurements. Those
measurements are: V=11.65 Landolt (1979); V=11.73 van Leeuwen, Alphenaar, & Brand
(1986); V=11.43 Messina (2001).
At the bottom of Table 2, we provide a key to the source(s) of the optical photometry
provided in the table.
This research made use of the SIMBAD database operated at CDS, Strasbourg, France,
and also of the NED and NStED databases operated at IPAC, Pasadena, USA. A large
amount of data for the Pleiades (and other open clusters) can also be found at the open
cluster database WEBDA (http://www.univie.ac.at/webda/), operated in Vienna by Ernst
Paunzen.
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This preprint was prepared with the AAS LATEX macros v5.2.
– 32 –
Table 1. Pleiades Membership Surveys used as Sources
Reference Area Covered Magnitude Range Number Candidates Name
Sq. Deg. (and band) Prefix
Trumpler (1921) 3 2.5< B <14.5 174 Tr
Trumpler (1921)a 24 2.5< B <10 72 Tr
Hertzsprung (1947) 4 2.5< V <15.5 247 HII
Artyukhina (1969) 60 2.5< B <12.5 ∼200 AK
Haro et al. (1982) 20 11< V <17.5 519 HCG
van Leeuwen et al. (1986) 80 2.5< B <13 193 PELS
Stauffer et al. (1991) 16 14< V <18 225 SK
Hambly et al. (1993) 23 10< I <17.5 440 HHJ
Pinfield et al. (2000) 6 13.5< I <19.5 339 BPL
Adams et al. (2001) 300 8< Ks <14.5 1200 ...
Deacon & Hambly (2004) 75 10< R <19 916 DH
aThe Trumpler paper is listed twice because there are two membership surveys included in that
paper, with differing spatial coverages and different limiting magnitudes.
Fig. 1.— Spatial coverage of the six times deeper “2MASS 6x” observations of the Pleiades.
The 2MASS survey region is approximately centered on Alcyone, the most massive member
of the Pleiades. The trapezoidal box roughly indicates the region covered with the shallow
IRAC survey of the cluster core. The star symbols correspond to the brightest B star
members of the cluster. The red points are the location of objects in the 2MASS 6x Point
Source Catalog.
Fig. 2.— Color-magnitude diagram for the Pleiades derived from the 2MASS 6x obser-
vations. The red dots correspond to objects identified as unresolved, whereas the green
dots correspond to extended sources (primarily background galaxies). The lack of green
dots fainter than K = 16 is indicative that there is too few photons to identify sources as
extended - the extragalactic population presumably increases to fainter magnitudes.
Fig. 3.— As for Figure 2, except in this case the axes are J − H and H − Ks. The
extragalactic objects are very red in both colors.
Fig. 4.— FIGURE REMOVED TO FIT WITHIN ASTRO-PH FILESIZE GUIDELINES.
See http://spider.ipac.caltech.edu/staff/stauffer/pleiades07/ for full-res version. Two-color
(4.5 and 8.0 micron) mosaic of the central square degree of the Pleiades from the IRAC
survey. North is approximately vertical, and East is approximately to the left. The bright
star nearest the center is Alcyone; the bright star at the left of the mosaic is Atlas; and the
bright star at the right of the mosaic is Electra.
http://spider.ipac.caltech.edu/staff/stauffer/pleiades07/
Fig. 5.— Finding chart corresponding approximately to the region imaged with IRAC.
The large, five-pointed stars are all of the Pleiades members brighter than V= 5.5. The
small open circles correspond to other cluster members. Several stars with 8 µm excesses are
labelled by their HII numbers, and are discussed further in Section 6. The short lines through
several of the stars indicate the size and position angle of the residual optical polarization
(after subtraction of a constant foreground component), as provided in Figure 6 of Breger
(1986).
Fig. 6.— Comparison of aperture photometry for Pleiades members derived from the
IRAC 3.6 µm mosaic using the Spitzer APEX package and the IRAF implementation of
DAOPHOT.
Fig. 7.— Difference between aperture photometry for Pleiades members for IRAC channels
1 and 2. The [3.6]−[4.5] color begins to depart from essentially zero at magnitudes ∼10.5,
corresponding approximately to spectral type M0 in the Pleiades.
Fig. 8.— Ks vs. Ks −[4.5] CMD for Pleiades candidate members, illustrating why we have
excluded HII 1695 from the final catalog of cluster members. The “X” symbol marks the
location of HII 1695 in this diagram.
Fig. 9.— Spatial plot of the candidate Pleiades members from Table 2. The large star
symbols are members brighter than Ks= 6; the open circles are stars with 6 < Ks < 9; and
the dots are candidate members fainter than Ks= 9. The solid line is parallel to the galactic
plane.
Fig. 10.— The cumulative radial density profiles for Pleiades members in several magnitude
ranges: heavy, long dash – Ks < 6; dots – 6 < Ks < 9; short dash – 9 < Ks < 12; light, long
dash – Ks > 12.
Fig. 11.— V vs. (V − I)c CMD for Pleiades members with photoelectric photometry. The
solid curve is the “by eye” fit to the single-star locus for Pleiades members.
Fig. 12.— Ks vs. Ks − [3.6] CMD for Pleiades candidate members from Table 2 (dots) and
from deeper imaging of a set of Pleiades VLM and brown dwarf candidate members from
Lowrance et al. (2007) (squares). The solid curve is the single-star locus from Table 3.
Fig. 13.— V vs. (V − I)c CMD for Pleiades candidate members from Table 2 for which we
have photoelectric photometry, compared to theoretical isochrones from Siess et al. (2000)
(left) and from Baraffe et al. (1998) (right). For the left panel, the curves correspond to 10,
50, 100 Myr and a ZAMS; the right panel includes curves for 50, 100 Myr and a ZAMS.
Fig. 14.— K vs. (I −K) CMD for Pleiades candidate members from Table 2, compared to
theoretical isochrones from Siess et al. (2000) (left) and from Baraffe et al. (1998) (right).
The curves correspond to 50 Myr, 100 Myr and a ZAMS.
Fig. 15.— Ks vs. Ks−[3.6] CMD for the objects in the central one square degree of the
Pleiades, combining data from the IRAC shallow survey and 2MASS. The symbols are defined
within the figure (and see text for details). The dashed-line box indicates the region within
which we have searched for new candidate Pleiades VLM and substellar members. The solid
curve is a DUSTY 100 Myr isochrone from Chabrier et al. (2000), for masses from 0.1 to
0.03 M⊙.
Fig. 16.— Proper motion vector point diagrams (VPDs) for various stellar samples in the
central one degree field, derived from combining the IRAC and 2MASS 6x observations.
Top left: VPD comparing all objects in the field (small black dots) to Pleiades members
with 11 < Ks < 14 (large blue dots). Top right: same, except the blue dots are the new
candidate VLM and substellar Pleiades members. Bottom left: same, except the blue dots
are a nearby, low mass field star sample from a box just blueward of the trapezoidal region
in 15. Bottom right: VPD just showing a second, distant field star sample as described in
the text.
Fig. 17.— Same as Fig. 15, except that the new candidate VLM and substellar objects from
Table 4 are now indicated as small, red squares.
Fig. 18.— Two plots intended to isolate Pleiades members with excess and/or extended 8
µm emission. The plot with [3.6]−[8.0] micron colors shows data from Table 3 (and hence is
for aperture sizes of 3 pixel and 2 pixel radius, respectively). The increased vertical spread
in the plots at faint magnitudes is simply due to decreasing signal to noise at 8 µm. The
numbers labelling stars with excesses are the HII identification numbers for those stars.
Fig. 19.— Postage stamp images extracted from individual, 8 µm BCDs for the stars with
extended 8 µm emission, from which we have subtracted an empirical PSF. Clockwise from
the upper left, the stars shown are HII1234, HII859, Merope and HII652. The five-pointed
star indicates the astrometric position of the star (often superposed on a few black pixels
where the 8 µm image was saturated. The circle in the Merope image is centered on the
location of IC349 and has diameter about 25” (the size of IC349 in the optical is of order
10” x 10”).
Fig. 20.— Aperture growth curves from the 8 µm mosaic for stars with 24 µm excesses from
Gorlova et al. (2006) and for a set of control objects (dashed curves). All of the objects have
been scaled to common zero-point magnitudes for 9 pixel apertures, with the 24 µm excess
stars offset from the control objects by 0.1 mag. The three Gorlova et al. (2006) stars with
no excess at 8 µm are HII 996, HII 1284 and HII 2195. The Gorlova et al. (2006) star with
a slight excess at 8 µm is HII 489.
Fig. 21.— Calibration derived relating Ikp from Pinfield et al. (2000) and IC. The dots are
stars for which we have both Ikp and IC measurements (small dots: photographic IC; large
dots: photoelectric IC), and the solid line indicates the piecewise linear fit we use to convert
the Ikp values to IC for stars for which we only have Ikp.
Fig. 22.— Difference between the predicted IC and Deacon & Hambly (2004) I magnitude
as a function of right ascension for the DH stars. No obvious dependence is present versus
declination.
Fig. 23.— Comparison of the recalibrated DH I photometry with estimates of IC for stars
in Table 2 with photoelectric data.
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	Introduction
	New Observational Data
	2MASS ``6x" Imaging of the Pleiades
	Shallow IRAC Imaging
	A Catalog of Pleiades Candidate Members
	Empirical Pleiades Isochrones and Comparison to Model Isochrones
	Identification of New Very Low Mass Candidate Members
	Mid-IR Observations of Dust and PAHS in the Pleiades
	Pleiades Dust-Star Encounters Imaged with IRAC
	Circumstellar Disks and IRAC
	Summary and Conclusions
	APPENDIX
	Membership Catalogs
	Photometry

  We make use of new near and mid-IR photometry of the Pleiades cluster in
order to help identify proposed cluster members. We also use the new photometry
with previously published photometry to define the single-star main sequence
locus at the age of the Pleiades in a variety of color-magnitude planes.
  The new near and mid-IR photometry extend effectively two magnitudes deeper
than the 2MASS All-Sky Point Source catalog, and hence allow us to select a new
set of candidate very low mass and sub-stellar mass members of the Pleiades in
the central square degree of the cluster. We identify 42 new candidate members
fainter than Ks =14 (corresponding to 0.1 Mo). These candidate members should
eventually allow a better estimate of the cluster mass function to be made down
to of order 0.04 solar masses.
  We also use new IRAC data, in particular the images obtained at 8 um, in
order to comment briefly on interstellar dust in and near the Pleiades. We
confirm, as expected, that -- with one exception -- a sample of low mass stars
recently identified as having 24 um excesses due to debris disks do not have
significant excesses at IRAC wavelengths. However, evidence is also presented
that several of the Pleiades high mass stars are found to be impacting with
local condensations of the molecular cloud that is passing through the Pleiades
at the current epoch.

Introduction
Because of its proximity, youth, richness, and location in the northern hemisphere, the
Pleiades has long been a favorite target of observers. The Pleiades was one of the first open
clusters to have members identified via their common proper motion (Trumpler 1921), and
the cluster has since then been the subject of more than a dozen proper motion studies.
Some of the earliest photoelectric photometry was for members of the Pleiades (Cummings
1921), and the cluster has been the subject of dozens of papers providing additional optical
photometry of its members. The youth and nearness of the Pleiades make it a particularly
attractive target for identifying its substellar population, and it was the first open cluster
studied for those purposes (Jameson & Skillen 1989; Stauffer et al. 1989). More than 20 pa-
pers have been subsequently published, identifying additional substellar candidate members
of the Pleiades or studying their properties.
– 3 –
We have three primary goals for this paper. First, while extensive optical photometry
for Pleiades members is available in the literature, photometry in the near and mid-IR is
relatively spotty. We will remedy this situation by using new 2MASS JHKs and Spitzer
IRAC photometry for a large number of Pleiades members. We will use these data to help
identify cluster non-members and to define the single-star locus in color-magnitude diagrams
for stars of 100 Myr age. Second, we will use our new IR imaging photometry of the center of
the Pleiades to identify a new set of candidate substellar members of the cluster, extending
down to stars expected to have masses of order 0.04 M⊙. Third, we will use the IRAC data
to briefly comment on the presence of circumstellar debris disks in the Pleiades and the
interaction of the Pleiades stars with the molecular cloud that is currently passing through
the cluster.
In order to make best use of the IR imaging data, we will begin with a necessary
digression. As noted above, more than a dozen proper motion surveys of the Pleiades have
been made in order to identify cluster members. However, no single catalog of the cluster
has been published which attempts to collect all of those candidate members in a single
table and cross-identify those stars. Another problem is that while there have been many
papers devoted to providing optical photometry of cluster members, that photometry has
been bewilderingly inhomogeneous in terms of the number of photometric systems used. In
Sec. 3 and in the Appendix, we describe our efforts to create a reasonably complete catalog
of candidate Pleiades members and to provide optical photometry transformed to the best
of our ability onto a single system.
2. New Observational Data
2.1. 2MASS “6x” Imaging of the Pleiades
During the final months of Two Micron All Sky Survey (2MASS; Skrutskie et al. (2006))
operations, a series of special observations were carried out that employed exposures six times
longer than used for the the primary survey. These so-called “6x” observations targeted 30
regions of scientific interest including a 3 deg x 2 deg area centered on the Pleiades cluster.
The 2MASS 6x data were reduced using an automated processing pipeline similar to that
used for the main survey data, and a calibrated 6x Image Atlas and extracted 6x Point and
Extended Source Catalogs (6x-PSC and 6x-XSC) analogous to the 2MASS All-Sky Atlas,
PSC and XSC have been released as part of the 2MASS Extended Mission. A description
of the content and formats of the 6x image and catalog products, and details about the 6x
– 4 –
observations and data reduction are given by Cutri et al. (2006; section A3). 1 The 2MASS
6x Atlas and Catalogs may be accessed via the on-line services of the NASA/IPAC Infrared
Science Archive (http://irsa.ipac.caltech.edu).
Figure 1 shows the area on the sky imaged by the 2MASS 6x observations in the Pleiades
field. The region was covered by two rows of scans, each scan being one degree long (in
declination) and 8.5’ wide in right ascension. Within each row, the scans overlap by approx-
imately one arcminute in right ascension. There are small gaps in coverage in the declination
boundary between the rows, and one complete scan in the southern row is missing because
the data in that scan did not meet the minimum required photometric quality. The total
area covered by the 6x Pleiades observations is approximately 5.3 sq. degrees.
There are approximately 43,000 sources extracted from the 6x Pleiades observations
in the 2MASS 6x-PSC, and nearly 1,500 in the 6x-XSC. Because there are at most about
1000 Pleiades members expected in this region, only ∼2% of the 6x-PSC sources are cluster
members, and the rest are field stars and background galaxies. The 6x-XSC objects are
virtually all resolved background galaxies. Near infrared color-magnitude and color-color
diagrams of the unresolved sources from the 2MASS 6x-PSC and all sources in the 6x-XSC
sources from the Pleiades region are shown in Figures 2 and 3, respectively. The extragalactic
sources tend to be redder than most stars, and the galaxies become relatively more numerous
towards fainter magnitudes. Unresolved galaxies dominate the point sources that are fainter
than Ks > 15.5 and redder than J −Ks > 1.2 mag.
The 2MASS 6x observations were conducted using the same freeze-frame scanning tech-
nique used for the primary survey (Skrutskie et al. 2006). The longer exposure times were
achieved by increasing the “READ2-READ1” integration to 7.8 sec from the 1.3 sec used for
primary survey. However, the 51 ms “READ1” exposure time was not changed for the 6x
observations. As a result, there is an effective “sensitivity gap” in the 8-11 mag region where
objects may be saturated in the 7.8 sec READ2-READ1 6x exposures, but too faint to be
detected in the 51 ms READ1 exposures. Because the sensitivity gap can result in incom-
pleteness and/or flux bias in the photometric overlap regime, the near infrared photometry
for sources brighter than J=11 mag in the 6x-PSC was taken from the 2MASS All-Sky PSC
during compilation of the catalog of Pleiades candidate members presented in Table 2 (c.f.
Section 3).
1http://www.ipac.caltech.edu/2mass/releases/allsky/doc/explsup.html
http://irsa.ipac.caltech.edu
– 5 –
2.2. Shallow IRAC Imaging
Imaging of the Pleiades with Spitzer was obtained in April 2004 as part of a joint GTO
program conducted by the IRAC instrument team and the MIPS instrument team. Initial
results of the MIPS survey of the Pleiades have already been reported in Gorlova et al. (2006).
The IRAC observations were obtained as two astronomical observing requests (AORs). One
of them was centered near the cluster center, at RA=03h47m00.0s and Dec=24d07m (2000),
and consisted of a 12 row by 12 column map, with “frametimes” of 0.6 and 12.0 seconds
and two dithers at each map position. The map steps were 290′′ in both the column and
row direction. The resultant map covers a region of approximately one square degree, and
a total integration time per position of 24 sec over most of the map. The second AOR used
the same basic mapping parameters, except it was smaller (9 rows by 9 columns) and was
instead centered northwest from the cluster center at RA=03h44m36.0s and Dec=25d24m.
A two-band color image of the AOR covering the center of the Pleiades is shown in Figure 4.
A pictorial guide to the IRAC image providing Greek names for a few of the brightest stars,
and Hertzsprung (1947) numbers for several stars mentioned in Section 6 is provided in
Figure 5.
We began our analysis with the basic calibrated data (BCDs) from the Spitzer pipeline,
using the S13 version of the Spitzer Science Center pipeline software. Artifact mitigation
and masking was done using the IDL tools provided on the Spitzer contributed software
website. For each AOR, the artifact-corrected BCDs were combined into single mosaics
for each channel using the post-BCD “MOPEX” package (Makovoz & Marleau 2005). The
mosaic images were constructed with 1.22×1.22 arcsecond pixels (i.e., approximately the
same pixel size as the native IRAC arrays).
We derived aperture photometry for stars present in these IRAC mosaics using both
APEX (a component of the MOPEX package) and the “phot” routine in DAOPHOT. In
both cases, we used a 3 pixel radius aperture and a sky annulus from 3 to 7 pixels (except
that for Channel 4, for the phot package we used a 2 pixel radius aperture and a 2 to 6 pixel
annulus because that provided more reliable fluxes at low flux levels). We used the flux for
zero magnitude calibrations provided in the IRAC data handbook (280.9, 179.7, 115.0 and
64.1 Jy for Ch 1 through Ch 4, respectively), and the aperture corrections provided in the
same handbook (multiplicative flux correction factors of 1.124, 1.127, 1.143 and 1.584 for Ch
1-4, inclusive. The Ch4 correction factor is much bigger because it is for an aperture radius
of 2 rather than 3 pixels.).
Figure 6 and Figure 7 provide two means to assess the accuracy of the IRAC photometry.
The first figure compares the aperture photometry from APEX to that from phot, and shows
that the two packages yield very similar results when used in the same way. For this reason,
– 6 –
we have simply averaged the fluxes from the two packages to obtain our final reported value.
The second figure shows the difference between the derived 3.6 and 4.5 µm magnitudes
for Pleiades members. Based on previous studies (e.g. Allen et al. (2004)), we expected
this difference to be essentially zero for most stars, and the Pleiades data corroborate that
expectation. For [3.6]<10.5, the RMS dispersion of the magnitude difference between the two
channels is 0.024 mag. Assuming that each channel has similar uncertainties, this indicates
an internal 1-σ accuracy of order 0.017 mag. The absolute calibration uncertainty for the
IRAC fluxes is currently estimated at of order 0.02 mag. Figure 7 also shows that fainter than
[3.6]=10.5 (spectral type later than about M0), the [3.6]−[4.5] color for M dwarfs departs
slightly from zero, becoming increasingly redder to the limit of the data (about M6).
3. A Catalog of Pleiades Candidate Members
If one limits oneself to only stars visible with the naked eye, it is easy to identify which
stars are members of the Pleiades – all of the stars within a degree of the cluster center that
have V < 6 are indeed members. However, if one were to try to identify the M dwarf stellar
members of the cluster (roughly 14 < V < 23), only of order 1% of the stars towards the
cluster center are likely to be members, and it is much harder to construct an uncontaminated
catalog. The problem is exacerbated by the fact that the Pleiades is old enough that mass
segregation through dynamical processes has occurred, and therefore one has to survey a
much larger region of the sky in order to include all of the M dwarf members.
The other primary difficulty in constructing a comprehensive member catalog for the
Pleiades is that the pedigree of the candidates varies greatly. For the best studied stars,
astrometric positions can be measured over temporal baselines ranging up to a century or
more, and the separation of cluster members from field stars in a vector point diagram
(VPD) can be extremely good. In addition, accurate radial velocities and other spectral
indicators are available for essentially all of the bright cluster members, and these further
allow membership assessment to be essentially definitive. Conversely, at the faint end (for
stars near the hydrogen burning mass limit in the Pleiades), members are near the detection
limit of the existing wide-field photographic plates, and the errors on the proper motions
become correspondingly large, causing the separation of cluster members from field stars
in the VPD to become poor. These stars are also sufficiently faint that spectra capable of
discriminating members from field dwarfs can only be obtained with 8m class telescopes, and
only a very small fraction of the faint candidates have had such spectra obtained. Therefore,
any comprehensive catalog created for the Pleiades will necessarily have stars ranging from
certain members to candidates for which very little is known, and where the fraction of
– 7 –
spurious candidate members increases to lower masses.
In order to address the membership uncertainties and biases, we have chosen a sliding
scale for inclusion in our catalog. For all stars, we require that the available photometry yields
location in color-color and color-magnitude diagrams consistent with cluster membership.
For the stars with well-calibrated photoelectric photometry, this means the star should not
fall below the Pleiades single-star locus by more than about 0.2 mag or above that locus
by more than about 1.0 mag (the expected displacement for a hierarchical triple with three
nearly equal mass components). For stars with only photographic optical photometry, where
the 1-σ uncertainties are of order 0.1 to 0.2 mag, we still require the star’s photometry to
be consistent with membership, but the allowed displacements from the single star locus are
considerably larger. Where accurate radial velocities are known, we require that the star
be considered a radial velocity member based on the paper where the radial velocities were
presented. Where stars have been previously identified as non-members based on photometric
or spectroscopic indices, we adopt those conclusions.
Two other relevant pieces of information are sometimes available. In some cases, in-
dividual proper motion membership probabilities are provided by the various membership
surveys. If no other information is available, and if the membership probability for a given
candidate is less than 0.1, we exclude that star from our final catalog. However, often a star
appears in several catalogs; if it appears in two or more proper motion membership lists we
include it in the final catalog even if P < 0.1 in one of those catalogs. Second, an entirely
different means to identify candidate Pleiades members is via flare star surveys towards the
cluster (Haro et al. 1982; Jones 1981). A star with a formally low membership probability
in one catalog but whose photometry is consistent with membership and that was identified
as a flare star is retained in our catalog.
Further details of the catalog construction are provided in the appendix, as are details of
the means by which the B, V , and I photometry have been homogenized. A full discussion
and listing of all of the papers from which we have extracted astrometric and photometric
information is also provided in the appendix. Here we simply provide a very brief description
of the inputs to the catalog.
We include candidate cluster members from the following proper motion surveys: Trumpler
(1921), Hertzsprung (1947), Jones (1981), Pels and Lub – as reported in van Leeuwen, Alphenaar, & Brand
(1986), Stauffer et al. (1991), Artyukhina (1969), Hambly et al. (1993), Pinfield et al. (2000),
Adams et al. (2001) and Deacon & Hambly (2004). Another important compilation which
provides the initial identification of a significant number of low mass cluster members is the
flare star catalog of Haro et al. (1982). Table 1 provides a brief synopsis of the characteristics
of the candidate member catalogs from these papers. The Trumpler paper is listed twice
– 8 –
in Table 1 because there are two membership surveys included in that paper, with differing
spatial coverages and different limiting magnitudes.
In our final catalog, we have attempted to follow the standard naming convention
whereby the primary name is derived from the paper where it was first identified as a cluster
member. An exception to this arises for stars with both Trumpler (1921) and Hertzsprung
(1947) names, where we use the Hertzsprung numbers as the standard name because that
is the most commonly used designation for these stars in the literature. The failure for
the Trumpler numbers to be given precedence in the literature perhaps stems from the fact
that the Trumpler catalog was published in the Lick Observatory Bulletins as opposed to a
refereed journal. In addition to providing a primary name for each star, we provide cross-
identifications to some of the other catalogs, particularly where there is existing photometry
or spectroscopy of that star using the alternate names. For the brightest cluster members, we
provide additional cross-references (e.g., Greek names, Flamsteed numbers, HD numbers).
For each star, we attempt to include an estimate for Johnson B and V , and for Cousins
I (IC). Only a very small fraction of the cluster members have photoelectric photometry in
these systems, unfortunately. Photometry for many of the stars has often been obtained in
other systems, including Walraven, Geneva, Kron, and Johnson. We have used previously
published transformations from the appropriate indices in those systems to Johnson BV or
Cousins I. In other cases, photometry is available in a natural I band system, primarily
for some of the relatively faint cluster members. We have attempted to transform those I
band data to IC by deriving our own conversion using stars for which we already have a IC
estimate as well as the natural I measurement. Details of these issues are provided in the
Appendix.
Finally, we have cross-correlated the cluster candidates catalog with the 2MASS All-Sky
PSC and also with the 6x-PSC for the Pleiades. For every star in the catalog, we obtain
JHKs photometry and 2MASS positions. Where we have both main survey 2MASS data
and data from the 6x catalog, we adopt the 6x data for stars with J >11, and data from
the standard 2MASS catalog otherwise. We verified that the two catalogs do not have any
obvious photometric or astrometric offsets relative to each other. The coordinates we list in
our catalog are entirely from these 2MASS sources, and hence they inherit the very good
and homogeneous 2MASS positional accuracies of order 0.1 arcseconds RMS.
We have then plotted the candidate Pleiades members in a variety of color-magnitude
diagrams and color-color diagrams, and required that a star must have photometry that is
consistent with cluster membership. Figure 8 illustrates this process, and indicates why (for
example) we have excluded HII 1695 from our final catalog.
– 9 –
Table 2 provides the collected data for the 1417 stars we have retained as candidate
Pleiades members. The first two columns are the J2000 RA and Dec from 2MASS; the next
are the 2MASS JHKs photometry and their uncertainties, and the 2MASS photometric
quality flag (“ph-qual”). If the number following the 2MASS quality flag is a 1, the 2MASS
data come from the 2MASS All-Sky PSC; if it is a 2, the data come from the 6x-PSC. The
next three columns provide the B, V and IC photometry, followed by a flag which indicates
the provenance of that photometry. The last column provides the most commonly used
names for these stars. The hydrogen burning mass limit for the Pleiades occurs at about
V=22, I=18, Ks=14.4. Fifty-three of the candidate members in the catalog are fainter than
this limit, and hence should be sub-stellar if they are indeed Pleiades members.
Table 3 provides the IRAC [3.6], [4.5], [5.8] and [8.0] photometry we have derived for
Pleiades candidate members included within the region covered by the IRAC shallow survey
of the Pleiades (see section 2). The brightest stars are saturated even in our short integration
frame data, particularly for the more sensitive 3.6 and 4.5 µm channels. At the faint end,
we provide photometry only for 3.6 and 4.5 µm because the objects are undetected in the
two longer wavelength channels. At the “top” and “bottom” of the survey region, we have
incomplete wavelength coverage for a band of width about 5′, and for stars in those areas
we report only photometry in either the 3.6 and 5.8 bands or in 4.5 and 8.0 bands.
Because Table 2 is an amalgam of many previous catalogs, each of which have different
spatial coverage, magnitude limits and other idiosyncrasies, it is necessarily incomplete and
inhomogeneous. It also certainly includes some non-members. For V < 12, we expect very
few non-members because of the extensive spectroscopic data available for those stars; the
fraction of non-members will likely increase to fainter magnitudes, particularly for stars
located far from the cluster center. The catalog is simply an attempt to collect all of the
available data, identify some of the non-members and eliminate duplications. We hope that
it will also serve as a starting point for future efforts to produce a “cleaner” catalog.
Figure 9 shows the distribution on the sky of the stars in Table 2. The complete
spatial distribution of all members of the Pleiades may differ slightly from what is shown
due to the inhomogeneous properties of the proper motion surveys. However, we believe
that those effects are relatively small and the distribution shown is mostly representative
of the parent population. One thing that is evident in Figure 9 is mass segregation – the
highest mass cluster members are much more centrally located than the lowest mass cluster
members. This fact is reinforced by calculating the cumulative number of stars as a function
of distance from the cluster center for different absolute magnitude bins. Figure 10 illustrates
this fact. Another property of the Pleiades illustrated by Figure 9 is that the cluster appears
to be elongated parallel to the galactic plane, as expected from n-body simulations of galactic
– 10 –
clusters (Terlevich 1987). Similar plots showing the flattening of the cluster and evidence for
mass segregation for the V < 12 cluster members were provided by (Raboud & Mermilliod
1998).
4. Empirical Pleiades Isochrones and Comparison to Model Isochrones
Young, nearby, rich open clusters like the Pleiades can and should be used to provide
template data which can help interpret observations of more distant clusters or to test
theoretical models. The identification of candidate members of distant open clusters is often
based on plots of stars in a color-magnitude diagram, overlaid upon which is a line meant
to define the single-star locus at the distance of the cluster. The stars lying near or slightly
above the locus are chosen as possible or probable cluster members. The data we have
collected for the Pleiades provide a means to define the single-star locus for 100 Myr, solar
metallicity stars in a variety of widely used color systems down to and slightly below the
hydrogen burning mass limit. Figure 11 and Figure 12 illustrate the appearance of the
Pleiades stars in two of these diagrams, and the single-star locus we have defined. The curve
defining the single-star locus was drawn entirely “by eye.” It is displaced slightly above the
lower envelope to the locus of stars to account for photometric uncertainties (which increase
to fainter magnitudes). We attempted to use all of the information available to us, however.
That is, there should also be an upper envelope to the Pleiades locus in these diagrams, since
equal mass binaries should be displaced above the single star sequence by 0.7 magnitudes
(and one expects very few systems of higher multiplicity). Therefore, the single star locus
was defined with that upper envelope in mind. Table 4 provides the single-star loci for the
Pleiades for BV IcJKs plus the four IRAC channels. We have dereddened the empirical loci
by the canonical mean extinction to the Pleiades of AV = 0.12 (and, correspondingly, AB
= 0.16, AI = 0.07, AJ = 0.03, AK = 0.01, as per the reddening law of Rieke & Lebofsky
(1985)).
The other benefit to constructing the new catalog is that it can provide an improved
comparison dataset to test theoretical isochrones. The new catalog provides homogeneous
photometry in many photometric bands for stars ranging from several solar masses down
to below 0.1 M⊙. We take the distance to the Pleiades as 133 pc, and refer the reader to
Soderblom et al. (2005) for a discussion and a listing of the most recent determinations. The
age of the Pleiades is not as well-defined, but is probably somewhere between 100 and 125
Myr (Meynet, Mermilliod, & Maeder 1993; Stauffer et al. 1998). We adopt 100 Myr for the
purposes of this discussion; our conclusions relative to the theoretical isochrones would not be
affected significantly if we instead chose 125 Myr. As noted above, we adopt AV=0.12 as the
– 11 –
mean Pleiades extinction, and apply that value to the theoretical isochrones. A small number
of Pleiades members have significantly larger extinctions (Breger 1986; Stauffer & Hartmann
1987), and we have dereddened those stars individually to the mean cluster reddening.
Figures 13 and 14 compare theoretical 100 Myr isochrones from Siess et al. (2000) and
Baraffe et al. (1998) to the Pleiades member photometry from Table 2 for stars for which
we have photoelectric photometry. Neither set of isochrones are a good fit to the V − I
based color-magnitude diagram. For Baraffe et al. (1998) this is not a surprise because they
illustrated that their isochrones are too blue in V−I for cool stars in their paper, and ascribed
the problem as likely the result of an incomplete line list, resulting in too little absorption in
the V band. For Siess et al. (2000), the poor fit in the V − I CMD is somewhat unexpected
in that they transform from the theoretical to the observational plane using empirical color-
temperature relations. In any event, it is clear that neither model isochrones match the
shape of the Pleiades locus in the V vs. V − I plane, and therefore use of these V − I based
isochrones for younger clusters is not likely to yield accurate results (unless the color-Teff
relation is recalibrated, as described for example in Jeffries & Oliveira (2005)). On the other
hand, the Baraffe et al. (1998) model provides a quite good fit to the Pleiades single star
locus for an age of 100 Myr in the K vs. I − K plane.2. This perhaps lends support to
the hypothesis that the misfit in the V vs. V − I plane is due to missing opacity in their
V band atmospheres for low mass stars (see also Chabrier et al. (2000) for further evidence
in support of this idea). The Siess et al. (2000) isochrones do not fit the Pleiades locus in
the K vs. I − K plane particularly well, being too faint near I − K=2 and too bright for
I −K > 2.5.
5. Identification of New Very Low Mass Candidate Members
The highest spatial density for Pleiades members of any mass should be at the cluster
center. However, searches for substellar members of the Pleiades have generally avoided
the cluster center because of the deleterious effects of scattered light from the high mass
cluster members and because of the variable background from the Pleiades reflection nebulae.
The deep 2MASS and IRAC 3.6 and 4.5 µm imaging provide accurate photometry to well
below the hydrogen burning mass limit, and are less affected by the nebular emission than
shorter wavelength images. We therefore expect that it should be possible to identify a new
2These isochrones are calculated for the standard K filter, rather than Ks. However, the difference in
location of the isochrones in these plots because of this should be very slight, and we do not believe our
conclusions are significantly affected.
– 12 –
set of candidate Pleiades substellar members by combining our new near and mid-infrared
photometry.
The substellar mass limit in the Pleiades occurs at about Ks=14.4, near the limit of
the 2MASS All-Sky PSC. As illustrated in Figure 2, the deep 2MASS survey of the Pleiades
should easily detect objects at least two magnitudes fainter than the substellar limit. The
key to actually identifying those objects and separating them from the background sources
is to find color-magnitude or color-color diagrams which separate the Pleiades members from
the other objects. As shown in Figure 15, late-type Pleiades members separate fairly well
from most field stars towards the Pleiades in a Ks vs. Ks − [3.6] color-magnitude diagram.
However, as illustrated in Figure 2, in the Ks magnitude range of interest there is also a
large population of red galaxies, and they are in fact the primary contaminants to identi-
fying Pleiades substellar objects in the Ks vs. Ks − [3.6] plane. Fortunately, most of the
contaminant galaxies are slightly resolved in the 2MASS and IRAC imaging, and we have
found that we can eliminate most of the red galaxies by their non-stellar image shape.
Figure 15 shows the first step in our process of identifying new very low mass members of
the Pleiades. The red plus symbols are the known Pleiades members from Table 2. The red
open circles are candidate Pleiades substellar members from deep imaging surveys published
in the literature, mostly of parts of the cluster exterior to the central square degree, where
the IRAC photometry is from Lowrance et al. (2007). The blue, filled circles are field M
and L dwarfs, placed at the distance of the Pleiades, using photometry from Patten et al.
(2006). Because the Pleiades is ∼100 Myr, its very low mass stellar and substellar objects
will be displaced about 0.7 mag above the locus of the field M and L dwarfs according to the
Baraffe et al. (1998) and Chabrier et al. (2000) models, in accord with the location in the
diagram of the previously identified, candidate VLM and substellar objects. The trapezoidal
shaped region outlined with a dashed line is the region in the diagram which we define as
containing candidate new VLM and substellar members of the Pleiades. We place the faint
limit of this region at Ks=16.2 in order to avoid the large apparent increase in faint, red
objects for Ks> 16.2, caused largely by increasing errors in the Ks photometry. Also, the
2MASS extended object flags cease to be useful fainter than about Ks= 16.
We took the following steps to identify a set of candidate substellar members of the
Pleiades:
• keep only objects which fall in the trapezoidal region in Figure 15.
• remove objects flagged as non-stellar by the 2MASS pipeline software;
• remove objects which appear non-stellar to the eye in the IRAC images;
– 13 –
• remove objects which do not fall in or near the locus of field M and L dwarfs in a J−H
vs. H −Ks diagram;
• remove objects which have 3.6 and 4.5 µm magnitudes that differ by more than 0.2
• remove objects which fall below the ZAMS in a J vs. J −Ks diagram.
As shown in Figure 15, all stars earlier than about mid-M have Ks − [3.6] colors bluer
than 0.4. This ensures that for most of the area of the trapezoidal region, the primary
contaminants are distant galaxies. Fortunately, the 2MASS catalog provides two types of
flags for identifying extended objects. For each filter, a chi-square flag measures the match
between the objects shape and the instrumental PSF, with values greater than 2.0 generally
indicative of a non-stellar object. In order not to be misguided by an image artifact in one
filter, we throw out the most discrepant of the three flags and average the other two. We
discard objects with mean χ2 greater than 1.9. The other indicator is the 2MASS extended
object flag, which is the synthesis of several independent tests of the objects shape, surface
brightness and color (see Jarrett, T. et al (2000) for a description of this process). If one
simply excludes the objects classified as extended in the 2MASS 6x image by either of these
techniques, the number of candidate VLM and substellar objects lying inside the trapezoidal
region decreases by nearly a half.
We have one additional means to demonstrate that many of the identified objects are
probably Pleiades members, and that is via proper motions. The mean Pleiades proper
motion is ∆RA = 20 mas yr−1 and ∆Dec = −45 mas yr−1 (Jones 1973). With an epoch
difference of only 3.5 years between the deep 2MASS and IRAC imaging, the expected motion
for a Pleiades member is only 0.07 arcseconds in RA and −0.16 arcseconds in Dec. Given
the relatively large pixel size for the two cameras, and the undersampled nature of the IRAC
3.6 and 4.5 µm images, it is not a priori obvious that one would expect to reliably detect the
Pleiades motion. However, both the 2MASS and IRAC astrometric solutions have been very
accurately calibrated. Also, for the present purpose, we only ask whether the data support a
conclusion that most of the identified substellar candidates are true Pleiades members (i.e.,
as an ensemble), rather than that each star is well enough separated in a VPD to derive a
high membership probability.
Figure 16 provides a set of plots that we believe support the conclusion that the majority
of the surviving VLM and substellar candidates are Pleiades members. The first plot shows
the measured motions between the epoch of the 2MASS and IRAC observations for all known
Pleiades members from Table 2 that lie in the central square degree region and have 11 <
Ks < 14 (i.e., just brighter than the substellar candidates). The mean offset of the Pleiades
– 14 –
stellar members from the background population is well-defined and is quantitatively of the
expected magnitude and sign (+0.07 arcsec in RA and −0.16 arcsec in Dec). The RMS
dispersion of the coordinate difference for the field population in RA and Dec is 0.076 and
0.062 arcseconds, supportive of our claim that the relative astrometry for the two cameras
is quite good. Because we expect that the background population should have essentially
no mean proper motion, the non-zero mean “motion” of the field population of about <
∆RA>=0.3 arcseconds is presumably not real. Instead, the offset is probably due to the
uncertainty in transferring the Spitzer coordinate zero-point between the warm star-tracker
and the cryogenic focal plane. Because it is simply a zero-point offset applicable to all the
objects in the IRAC catalog, it has no effect on the ability to separate Pleiades members
from the field star population.
The second panel in Figure 16 shows the proper motion of the candidate Pleiades VLM
and substellar objects. While these objects do not show as clean a distribution as the known
members, their mean motion is clearly in the same direction. After removing 2-σ deviants,
the median offsets for the substellar candidates are 0.04 and −0.11 arcseconds in RA and
Dec, respectively. The objects whose motions differ significantly from the Pleiades mean
may be non-members or they may be members with poorly determined motions (since a few
of the high probability members in the first panel also show discrepant motions).
The other two panels in Figure 16 show the proper motions of two possible control
samples. The first control sample was defined as the set of stars that fall up to 0.3 magnitudes
below the lower sloping boundary of the trapezoid in Figure 15. These objects should be late
type dwarfs that are either older or more distant than the Pleiades or red galaxies. We used
the 2MASS data to remove extended or blended objects from the sample in the same way as
for the Pleiades candidates. If the objects are nearby field stars, we expect to see large proper
motions; if galaxies, the real proper motions would be small – but relatively large apparent
proper motions due to poor centroiding or different centroids at different effective wavelengths
could be present. The second control set was defined to have −0.1 < K − [3.6] < 0.1 and
14.0 < K < 14.5, and to be stellar based on the 2MASS flags. This control sample should
therefore be relatively distant G and K dwarfs primarily. Both control samples have proper
motion distributions that differ greatly from the Pleiades samples and that make sense for,
respectively, a nearby and a distant field star sample.
Figure 17 shows the Pleiades members from Table 2 and the 55 candidate VLM and
substellar members that survived all of our culling steps. We cross-correlated this list with the
stars from Table 2 and with a list of the previously identified candidate substellar members
of the cluster from other deep imaging surveys. Fourteen of the surviving objects correspond
to previously identified Pleiades VLM and substellar candidates. We provide the new list
– 15 –
of candidate members in Table 5. The columns marked as µ(RA) and µ(DEC) are the
measured motions, in arcsec over the 3.5 year epoch difference between the 2MASS-6x and
IRAC observations. Forty-two of these objects have Ks> 14.0, and hence inferred masses
less than about 0.1 M⊙; thirty-one of them have Ks> 14.4, and hence have inferred masses
below the hydrogen burning mass limit.
Our candidate list could be contaminated by foreground late type dwarfs that happen
to lie in the line of sight to the Pleiades. How many such objects should we expect? In
order to pass our culling steps, such stars would have to be mid to late M dwarfs, or early
to mid L dwarfs. We use the known M dwarfs within 8 pc to estimate how many field M
dwarfs should lie in a one square degree region and at distance between 70 and 100 parsecs
(so they would be coincident in a CMD with the 100 Myr Pleiades members). The result is
∼3 such field M dwarf contaminants. Cruz et al. (2006) estimate that the volume density of
L dwarfs is comparable to that for late-M dwarfs, and therefore a very conservative estimate
is that there might also be 3 field L dwarfs contaminating our sample. We regard this (6
contaminating field dwarfs) as an upper limit because our various selection criteria would
exclude early M dwarfs and late L dwarfs. Bihain et al. (2006) made an estimate of the
number of contaminating field dwarfs in their Pleiades survey of 1.8 square degrees; for the
spectral type range of our objects, their algorithm would have predicted just one or two
contaminating field dwarfs for our survey.
How many substellar Pleiades members should there be in the region we have surveyed?
That is, of course, part of the question we are trying to answer. However, previous studies
have estimated that the Pleiades stellar mass function for M < 0.5 M⊙ can be approximated
as a power-law with an exponent of -1 (dN/dM ∝ M−1). Using the known Pleiades members
from Table 2 that lie within the region of the IRAC survey and that have masses of 0.2 <
M/M⊙< 0.5 (as estimated from the Baraffe et al. (1998) 100 Myr isochrone) to normalize
the relation, the M−1 mass function predicts about 48 members in our search region and
with 14 < K < 16.2 (corresponding to 0.1 < M/M⊙< 0.035). Other studies have suggested
that the mass function in the Pleiades becomes shallower below 0.1 M⊙, dN/dM ∝ M
−0.6.
Using the same normalization as above, this functional form for the Pleiades mass function
for M < 0.1 M⊙ yields a prediction of 20 VLM and substellar members in our survey. The
number of candidates we have found falls between these two estimates. Better proper motions
and low-resolution spectroscopy will almost certaintly eliminate some of these candidates as
non-members.
– 16 –
6. Mid-IR Observations of Dust and PAHS in the Pleiades
Since the earliest days of astrophotography, it has been clear that the Pleiades stars
are in relatively close proximity to interstellar matter whose optical manifestation is the
spider-web like network of filaments seen particularly strongly towards several of the B stars
in the cluster. High resolution spectra of the brightest Pleiades stars as well as CO maps
towards the cluster show that there is gas as well as dust present, and that the (primary)
interstellar cloud has a significant radial velocity offset relative to the Pleiades (White 2003;
Federman & Willson 1984). The gas and dust, therefore, are not a remnant from the forma-
tion of the cluster but are simply evidence of a a transitory event as this small cloud passes by
the cluster in our line of sight (see also Breger (1986)). There are at least two claimed mor-
phological signatures of a direct interaction of the Pleiades with the cloud. White & Bally
(1993) provided evidence that the IRAS 60 and 100 µm image of the vicinity of the Pleiades
showed a dark channel immediately to the east of the Pleiades, which they interpreted as
the “wake” of the Pleiades as it plowed through the cloud from the east. Herbig & Simon
(2001) provided a detailed analysis of the optically brightest nebular feature in the Pleiades
– IC 349 (Barnard’s Merope nebula) – and concluded that the shape and structure of that
nebula could best be understood if the cloud was running into the Pleiades from the south-
east. Herbig & Simon (2001) concluded that the IC 349 cloudlet, and by extension the rest
of the gas and dust enveloping the Pleiades, are relatively distant outliers of the Taurus
molecular clouds (see also Eggen (1950) for a much earlier discussion ascribing the Merope
nebulae as outliers of the Taurus clouds). White (2003) has more recently proposed a hybrid
model, where there are two separate interstellar cloud complexes with very different space
motions, both of which are colliding simultaneously with the Pleiades and with each other.
Breger (1986) provided polarization measurements for a sample of member and back-
ground stars towards the Pleiades, and argued that the variation in polarization signatures
across the face of the cluster was evidence that some of the gas and dust was within the clus-
ter. In particular, Figure 6 of that paper showed a fairly distinct interface region, with little
residual polarization to the NE portion of the cluster and an L-shaped boundary running
EW along the southern edge of the cluster and then north-south along the western edge of
the cluster. Stars to the south and west of that boundary show relatively large polarizations
and consistent angles (see also our Figure 5 where we provide a few polarization vectors from
Breger (1986) to illustrate the location of the interface region and the fact that the position
angle of the polarization correlates well with the location in the interface).
There is a general correspondence between the polarization map and what is seen with
IRAC, in the sense that the B stars in the NE portion of the cluster (Atlas and Alcyone)
have little nebular emission in their vicinity, whereas those in the western part of the cluster
– 17 –
(Maia, Electra and Asterope) have prominent, filamentary dust emission in their vicinity.
The L-shaped boundary is in fact visible in Figure 4 as enhanced nebular emission running
between and below a line roughly joining Merope and Electra, and then making a right
angle and running roughly parallel to a line running from Electra to Maia to HII1234 (see
Figure 5).
6.1. Pleiades Dust-Star Encounters Imaged with IRAC
The Pleiades dust filaments are most strongly evident in IRAC’s 8 µm channel, as
evidenced by the distinct red color of the nebular features in Figure 4. The dominance at 8
µm is an expected feature of reflection nebulae, as exemplified by NGC 7023 (Werner et al.
2004), where most of the mid-infrared emission arises from polycyclic aromatic hydrocarbons
(PAHs) whose strongest bands in the 3 to 10 µm region fall at 7.7 and 8.6 µm. One might
expect that if portions of the passing cloud were particularly near to one of the Pleiades
members, it might be possible to identify such interactions by searching for stars with 8.0
µm excesses or for stars with extended emission at 8 µm. Figure 18 provides two such plots.
Four stars stand out as having significant extended 8 µm emission, with two of those stars
also having an 8 µm excess based on their [3.6]−[8.0] color. All of these stars, plus IC 349,
are located approximately along the interface region identified by Breger (1986).
We have subtracted a PSF from the 8 µm images for the stars with extended emission,
and those PSF-subtracted images are provided in Figure 19. The image for HII 1234 has
the appearance of a bow-shock. The shape is reminiscent of predictions for what one should
expect from a collision between a large cloud or a sheet of gas and an A star as described
in Artymowicz & Clampin (1997). The Artymowicz & Clampin (1997) model posits that
A stars encountering a cloud will carve a paraboloidal shaped cavity in the cloud via radi-
ation pressure. The exact size and shape of the cavity depend on the relative velocity of
the encounter, the star’s mass and luminosity and properties of the ISM grains. For typical
parameters, the predicted characteristic size of the cavity is of order 1000 AU, quite compa-
rable to the size of the structures around HII 652 and HII 1234. The observed appearance
of the cavity depends on the view angle to the observer. However, in any case, the direction
from which the gas is moving relative to the star can be inferred from the location of the
star relative to the curved rim of the cavity; the “wind” originates approximately from the
direction connecting the star and the apex of the rim. For HII 1234, this indicates the cloud
which it is encountering has a motion relative to HII 1234 from the SSE, in accord with a
Taurus origin and not in accord for where a cloud is impacting the Pleiades from the west
as posited in White (2003). The nebular emission for HII 652 is less strongly bow-shaped,
– 18 –
but the peak of the excess emission is displaced roughly southward from the star, consistent
with the Taurus model and inconsistent with gas flowing from the west.
Despite being the brightest part of the Pleiades nebulae in the optical, IC 349 appears
to be undetected in the 8 µm image. This is not because the 8 µm image is insensitive to
the nebular emission - there is generally good agreement between the structures seen in the
optical and at 8 µm, and most of the filaments present in optical images of the Pleiades are
also visible on the 8 µm image (see Figures 4 and 19) and even the psf-subtracted image of
Merope shows well-defined nebular filaments. The lack of enhanced 8 µm emission from the
region of IC 349 is probably because all of the small particles have been scoured away from
this cloudlet, consistent with Herbig’s model to explain the HST surface photometry and
colors. There is no PAH emission from IC 349 because there are none of the small molecules
that are the postulated source of the PAH emission.
IC349 is very bright in the optical, and undetected to a good sensitivity limit at 8 µm; it
must be detectable via imaging at some wavelength between 5000 Å and 8 µm. We checked
our 3.6 µm data for this purpose. In the standard BCD mosaic image, we were unable to
discern an excess at the location of IC349 either simply by displaying the image with various
stretches or by doing cuts through the image. We performed a PSF subtraction of Merope
from the image in order to attempt to improve our ability to detect faint, extended emission
30” from Merope - unfortunately, bright stars have ghost images in IRAC Ch. 1, and in this
case the ghost image falls almost exactly at the location of IC349. IC349 is also not detected
in visual inspection of our 2MASS 6x images.
6.2. Circumstellar Disks and IRAC
As part of the Spitzer FEPS (Formation and Evolution of Planetary Systems) Legacy
program, using pointed MIPS photometry, Stauffer et al. (2005) identified three G dwarfs
in the Pleiades as having 24 µm excesses probably indicative of circumstellar dust disks.
Gorlova et al. (2006) reported results of a MIPS GTO survey of the Pleiades, and identified
nine cluster members that appear to have 24 µm excesses due to circumstellar disks. However,
it is possible that in a few cases these apparent excesses could be due instead to a knot of
the passing interstellar dust impacting the cluster member, or that the 24 µm excess could
be flux from a background galaxy projected onto the line of sight to the Pleiades member.
Careful analysis of the IRAC images of these cluster members may help confirm that the
MIPS excesses are evidence for debris disks rather than the other possible explanations.
Six of the Pleiades members with probable 24 µm excesses are included in the region
– 19 –
mapped with IRAC. However, only four of them have data at 8 µm – the other two fall
near the edge of the mapped region and only have data at 3.6 and 5.8 µm. None of the
six stars appear to have significant local nebular dust from visual inspection of the IRAC
mosaic images. Also, none of them appear problematic in Figure 18. For a slightly more
quantitative analysis of possible nebular contamination, we also constructed aperture growth
curves for the six stars, and compared them to other Pleiades members. All but one of the
six show aperture growth curves that are normal and consistent with the expected IRAC
PSF. The one exception is HII 489, which has a slight excess at large aperture sizes as is
illustrated in Figure 20. Because HII 489 only has a small 24 µm excess, it is possible that
the 24 µm excess is due to a local knot of the interstellar cloud material and is not due to a
debris disk. For the other five 24 µm excess stars we find no such problem, and we conclude
that their 24 µm excesses are indeed best explained as due to debris disks.
7. Summary and Conclusions
We have collated the primary membership catalogs for the Pleiades to produce the first
catalog of the cluster extending from its highest mass members to the substellar limit. At the
bright end, we expect this catalog to be essentially complete and with few or no non-member
contaminants. At the faint end, the data establishing membership are much sparser, and we
expect a significant number of objects will be non-members. We hope that the creation of
this catalog will spur efforts to obtain accurate radial velocities and proper motions for the
faint candidate members in order to eventually provide a well-vetted membership catalog
for the stellar members of the Pleiades. Towards that end, it would be useful to update
the current catalog with other data – such as radial velocities, lithium equivalent widths,
x-ray fluxes, Hα equivalent widths, etc. – which could be used to help accurately establish
membership for the low mass cluster candidates. It is also possible to make more use of
“negative information” present in the proper motion catalogs. That is, if a member from
one catalog is not included in another study but does fall within its areal and luminosity
coverage, that suggests that it likely failed the membership criteria of the second study. For
a few individual stars, we have done this type of comparison, but a systematic analysis of
the proper motion catalogs should be conducted. We intend to undertake these tasks, and
plan to establish a website where these data would be hosted.
We have used the new Pleiades member catalog to define the single-star locus at 100
Myr for BV IcKs and the four IRAC bands. These curves can be used as empirical calibration
curves when attempting to identify members of less well-studied, more distant clusters of
similar age. We compared the Pleiades photometry to theoretical isochrones from Siess et al.
– 20 –
(2000) and Baraffe et al. (1998). The Siess et al. (2000) isochrones are not, in detail, a good
fit to the Pleiades photometry, particularly for low mass stars. The Baraffe et al. (1998) 100
Myr isochrone does fit the Pleiades photometry very well in the I vs. I −K plane.
We have identified 31 new substellar candidate members of the Pleiades using our com-
bined seven-band infrared photometry, and have shown that the majority of these objects
appear to share the Pleiades proper motion. We believe that most of the objects that may
be contaminating our list of candidate brown dwarfs are likely to be unresolved galaxies, and
therefore low resolution spectroscopy should be able to provide a good criterion for culling
our list of non-members.
The IRAC images, particularly the 8 µm mosaic, provide vivid evidence of the strong in-
teraction of the Pleiades stars and the interstellar cloud that is passing through the Pleiades.
Our data are supportive of the model proposed by Herbig & Simon (2001) whereby the pass-
ing cloud is part of the Taurus cloud complex and hence is encountering the Pleiades from
the SSE direction. White & Bally (1993) had proposed a model whereby the cloud was
encountering the Pleiades from the west and used this to explain features in the IRAS 60
and 100 µm images of the region as the wake of the Pleiades moving through the cloud.
Our data appear to not be supportive of that hypothesis, and therefore leaves the apparent
structure in the IRAS maps as unexplained.
Most of the support for this work was provided by the Jet Propulsion Laboratory, Cal-
ifornia Institute of Technology, under NASA contract 1407. This research has made use of
NASA’s Astrophysics Data System (ADS) Abstract Service, and of the SIMBAD database,
operated at CDS, Strasbourg, France. This research has made use of data products from
the Two Micron All-Sky Survey (2MASS), which is a joint project of the University of Mas-
sachusetts and the Infrared Processing and Analysis Center, funded by the National Aero-
nautics and Space Administration and the National Science Foundation. These data were
served by the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion
Laboratory, California Institute of Technology, under contract with the National Aeronautics
and Space Administration. The research described in this paper was partially carried out at
the Jet Propulsion Laboratory, California Institute of Technology, under contract with the
National Aeronautics and Space Administration.
– 21 –
A. APPENDIX
A.1. Membership Catalogs
Membership lists of the Pleiades date back to antiquity if one includes historical and
literary references to the Seven Sisters (Alcyone, Maia, Merope, Electra, Taygeta, Asterope
and Celeno) and their parents (Atlas and Pleione). The first paper discussing relative proper
motions of a large sample of stars in the Pleiades (based on visual observations) was published
by Pritchard (1884). The best of the early proper motion surveys of the Pleiades derived
from photographic plate astrometry was that by Trumpler (1921), based on plates obtained
at Yerkes and Lick observatories. The candidate members from that survey were presented
in two tables, with the first being devoted to candidate members within about one degree
from the cluster center (operationally, within one degree from Alcyone) and the second table
being devoted to candidates further than one degree from the cluster center. Most of the
latter stars were denoted by Trumpler by an S or R, followed by an identification number.
We use Tr to designate the Trumpler stars (hence Trnnn for a star from the 1st table and
the small number of stars in the second table without an “S” or an “R”, and TrSnnn or
TrRnnn for the other stars). For the central region, Trumpler’s catalog extends to V ∼ 13,
while the outer region catalog includes stars only to about V ∼ 9.
The most heavily referenced proper motion catalog of the Pleiades is that provided by
Hertzsprung (1947). That paper makes reference to two separate catalogs: a photometric
catalog of the Pleiades published by Hertzsprung in 1923 (Hertzsprung 1923), whose members
are commonly referred to by HI numbers, and the new proper motion catalog from the
1947 paper, commonly referenced as the HII catalog. While both HI and HII numbers
have been used in subsequent observational papers, it is the HII identification numbers
that predominate. That catalog – derived from Carte du Ciel blue-sensitive plates from
14 observatories – includes stars in the central 2×2 degree region of the cluster, and has
a faint limit of about V = 15.5. Johnson system BV I photometry is provided for most
of the proposed Hertzsprung members in Johnson & Mitchell (1958) and Iriarte (1967).
Additional Johnson B and V photometry plus Kron I photometry for a fairly large number
of the Hertzsprung members can be found in Stauffer (1980), Stauffer (1982), and Stauffer
(1984). Other Johnson BV photometry for a scattering of stars can be found in Jones (1973),
Robinson & Kraft (1974), Messina (2001). Spectroscopic confirmation, primarily via radial
velocities, that these are indeed Pleiades members has been provided in Soderblom et al.
(1993); Queloz et al. (1998) and Mermilliod et al. (1997).
Two other proper motion surveys provide relatively bright candidate members relatively
far from the cluster center: Artyukhina & Kalinina (1970) and van Leeuwen, Alphenaar, & Brand
– 22 –
(1986). Stars from the Artyukhina catalog are designated as AK followed by the region from
which the star was identified followed by an identification number. The new members pro-
vided in the van Leeuwen paper were taken from an otherwise unpublished proper motion
study by Pels, where the first 118 stars were considered probable members and the remaining
75 stars were considered possible members. Van Leeuwen categorized a number of the Pels
stars as non-members based on the Walraven photometry they obtained, and we adopt those
findings. Radial velocities for stars in these two catalogs have been obtained by Rosvick et al
(1992), Mermilliod et al. (1997), and Queloz et al. (1998), and those authors identified a list
of the candidate members that they considered confirmed by the high resolution spectroscopy.
For these outlying candidate members, to be included in Table 2 we require that the star
be a radial velocity member from one of the above three surveys, or be indicated as having
“no dip” in the Coravel cross-correlation (indicating rapid rotation, which at least for the
later type stars is suggestive of membership). Geneva photometry of the Artyukhina stars
considered as likely members was provided by Mermilliod et al. (1997). The magnitude limit
of these surveys was not well-defined, but most of the Artyukhina and Pels stars are brighter
than V=13.
Jones (1973) provided proper motion membership probabilities for a large sample of
proposed Pleiades members, and for a set of faint, red stars towards the Pleiades. A few
star identification names from the sources considered by Jones appear in Table 2, including
MT (McCarthy & Treanor 1964), VM (van Maanen 1946), ALR (Ahmed et al. 1965), and
J (Jones 1973).
The chronologically next significant source of new Pleiades candidate members was
the flare star survey of the Pleiades conducted at several observatories in the 1960s, and
summarized in Haro et al. (1982), hereafter HCG. The logic behind these surveys was that
even at 100 Myr, late type dwarfs have relatively frequent and relatively high luminosity
flares (as demonstrated by Johnson & Mitchell (1958) having detected two flares during
their photometric observations of the Pleiades), and therefore wide area, rapid cadence
imaging of the Pleiades at blue wavelengths should be capable of identifying low mass cluster
members. However, such surveys also will detect relatively young field dwarfs, and therefore
it is best to combine the flare star surveys with proper motions. Dedicated proper motion
surveys of the HCG flare stars were conducted by Jones (1981) and Stauffer et al. (1991),
with the latter also providing photographic V I photometry (Kron system). Photoelectric
photometry for some of the HCG stars have been reported in Stauffer (1982), Stauffer (1984),
Stauffer & Hartmann (1987), and Prosser et al. (1991). High resolution spectroscopy of
many of the HCG stars is reported in Stauffer (1984), Stauffer & Hartmann (1987) and
Terndrup et al. (2000). Because a number of the papers providing additional observational
data for the flare stars were obtained prior to 1982, we also include in Table 2 the original
– 23 –
flare star names which were derived from the observatory where the initial flare was detected.
Those names are of the form an initial letter indicating the observatory – A (Asiago), B
(Byurakan), K (Konkoly), T (Tonantzintla) – followed by an identification number.
Stauffer et al. (1991) conducted two proper motion surveys of the Pleiades over an
approximately 4×4 degree region of the cluster based on plates obtained with the Lick 20′′
astrographic telescope. The first survey was essentially unbiased, except for the requirement
that the stars fall approximately in the region of the V vs. V − I color-magnitude diagram
where Pleiades members should lie. Candidate members from this survey are designated
by SK numbers. The second survey was a proper motion survey of the HCG stars. Photo-
graphic V I photometry of all the stars was provided as well as proper motion membership
probabilities. Photoelectric photometry for some of the candidate members was obtained as
detailed above in the section on the HCG catalog stars. The faint limit of these surveys is
about V=18.
Hambly et al. (1991) provided a significantly deeper, somewhat wider area proper mo-
tion survey, with the faintest members having V ≃ 20 and the total area covered being
of order 25 square degrees. The survey utilized red sensitive plates from the Palomar and
UK Schmidt telescopes. Due to incomplete coverage at one epoch, there is a vertical swath
slightly east of the cluster center where no membership information is available. Stars from
this survey are designated by their HHJ numbers . Hambly et al. (1993) provide RI photo-
graphic photometry on a natural system for all of their candidate members, plus photoelectric
Cousins RI photometry for a small number of stars and JHK photometry for a larger sam-
ple. Some spectroscopy to confirm membership has been reported in Stauffer et al. (1994),
Stauffer et al. (1995), Oppenheimer et al. (1997), Stauffer et al. (1998), and Steele et al.
(1995), though for most of the HHJ stars there is no spectroscopic membership confirma-
tion.
Pinfield et al. (2000) provide the deepest wide-field proper motion survey of the Pleiades.
That survey combines CCD imaging of six square degrees of the Pleiades obtained with the
Burrell Schmidt telescope (as five separate, non-overlapping fields near but outside the cluster
center) with deep photographic plates which provide the 1st epoch positions. Candidate
members are designated by BPL numbers (for Burrell Pleiades), with the faintest stars
having I ≃ 19.5, corresponding to V > 23. Only the stars brighter than about I= 17 have
sufficiently accurate proper motions to use to identify Pleiades members. Fainter than I=
17, the primary selection criteria are that the star fall in an appropriate place in both an I
vs. I − Z and an I vs. I −K CMD.
Adams et al. (2001) combined the 2MASS and digitized POSS databases to produce a
very wide area proper motion survey of the Pleiades. By design, that survey was very inclu-
– 24 –
sive - covering the entire physical area of the cluster and extending to the hydrogen burning
mass limit. However, it was also very “contaminated”, with many suspected non-members.
The catalog of possible members was not published. We have therefore not included stars
from this study in Table 2; we have used the proper motion data from Adams et al. (2001)
to help decide cases where a given star has ambiguous membership data from the other
surveys.
Deacon & Hambly (2004) provided another deep and very wide area proper motion
survey of the Pleiades. The survey covers a circular area of approximately five degrees
radius to R ∼ 20, or V ∼ 22. Candidate members are designated by DH. Deacon & Hambly
(2004) also provide membership probabilities based on proper motions for many candidate
cluster members from previous surveys. For stars where Deacon & Hambly (2004) derive
P < 0.1 and where we have no other proper motion information or where another proper
motion survey also finds low membership probability, we exclude the star from our catalog.
For cases where two of our proper motion catalogs differ significantly in their membership
assessment, with one survey indicating the star is a probable member, we retain the star
in the catalog as the conservative choice. Examples of the latter where Deacon & Hambly
(2004) derive P < 0.1 include HII 1553, HII 2147, HII 2278 and HII 2665 – all of which
we retain in our catalog because other surveys indicate these are high probability Pleiades
members.
A.2. Photometry
Photometry for stars in open cluster catalogs can be used to help confirm cluster mem-
bership and to help constrain physical properties of those stars or of the cluster. For a
variety of reasons, photometry of stars in the Pleiades has been obtained in a panoply of
different photometric systems. For our own goals, which are to use the photometry to help
verify membership and to define the Pleiades single-star locus in color magnitude diagrams,
we have attempted to convert photometry in several of these systems to a common sys-
tem (Johnson BV and Cousins I). We detail below the sources of the photometry and the
conversions we have employed.
Photoelectric photometry of Pleiades members dates back to at least 1921 (Cummings
1921). However, as far as we are aware the first “modern” photoelectric photometry for the
Pleiades, using a potassium hydride photoelectric cell, is that of Calder & Shapley (1937).
Eggen (1950) provided photoelectric photometry using a 1P21 phototube (but calibrated
to a no-longer-used photographic system) for most of the known Pleiades members within
one degree of the cluster center and with magnitudes < 11. The first phototube photom-
– 25 –
etry of Pleiades stars calibrated more-or-less to the modern UBV system was provided by
Johnson & Morgan (1951). An update of that paper, and the oldest photometry included
here was reported in Johnson & Mitchell (1958), which provided UBV Johnson system pho-
tometry for a large sample of HII and Trumpler candidate Pleiades members. Iriarte (1967)
later reported Johnson system V − I colors for most of these stars. We have converted
Iriarte’s V − I photometry to estimated Cousins V − I colors using a formula from Bessell
(1979):
V − I(Cousins) = 0.778× V − I(Johnson). (A1)
BV RI photometry for most of the Hertzsprung members fainter than V= 10 has been
published by Stauffer (1980), Stauffer (1982), Stauffer (1984), and Stauffer & Hartmann
(1987). The BV photometry is Johnson system, whereas the RI photometry is on the Kron
system. The Kron V − I colors were converted to Cousins V − I using a transformation
provided by Bessell & Weis (1987):
V − I(Cousins) = 0.227 + 0.9567(V − I)k + 0.0128(V − I)
k − 0.0053(V − I)
k (A2)
Other Kron system V−I colors have been published for Pleiades candidates in Stauffer et al.
(1991) (photographic photometry) and in Prosser et al. (1991). These Kron-system colors
have also been converted to Cousins V − I using the above formula.
Johnson/Cousins UBV R photometry for a set of low mass Pleiades members was pro-
vided by Landolt (1979). We only use the BV magnitudes from that study. Additional John-
son system UBV photometry for small numbers of stars is provided in Robinson & Kraft
(1974), Messina (2001) and Jones (1973).
van Leeuwen, Alphenaar, & Meys (1987) provided Walraven V BLUW photometry for
nearly all of the Hertzsprung members brighter than V ∼ 13.5 and for the Pels candidate
members. Van Leeuwen provided an estimated Johnson V derived from the Walraven V
in his tables. We have transformed the Walraven V − B color into an estimate of Johnson
B − V using a formula from Rosvick et al (1992):
B − V (Johnson) = 2.571(V − B)− 1.02(V −B)2 + 0.5(V − B)3 − 0.01 (A3)
Hambly et al. (1993) provided photographic V RI photometry for all of the HHJ candidate
members, and V RI Cousins photoelectric photometry for a small fraction of those stars. We
took all of the HHJ stars with photographic photometry for which we also have photoelectric
V I photometry on the Cousins system, and plotted V (Cousins) vs. V (HHJ) and I(Cousins)
vs. I(HHJ). While there is some evidence for slight systematic departures of the HHJ photo-
graphic photometry from the Cousins system, those departures are relatively small and we
have chosen simply to retain the HHJ values and treat them as Cousins system.
– 26 –
Pinfield et al. (2000) reported their I magnitudes in an instrumental system which they
designated as Ikp. We identified all BPL candidate members for which we had photoelectric
Cousins I estimates, and plotted Ikp vs. IC. Figure 21 shows this correlation, and the
piecewise linear fit we have made to convert from Ikp to IC. Our catalog lists these converted
IC measures for the BPL stars for which we have no other photoelectric I estimates.
Deacon & Hambly (2004) derived RI photometry from the scans of their plates, and
calibrated that photometry by reference to published photometry from the literature. When
we plotted their the difference between their I band photometry and literature values (where
available), we discovered a significant dependence on right ascension. Unfortunately, because
the DH survey extended over larger spatial scales than the calibrating photometry, we could
not derive a correction which we could apply to all the DH stars. We therefore developed
the following indirect scheme. We used the stars for which we have estimated IC magnitudes
(from photoelectric photometry) to define the relation between J and (IC−J) for Pleiades
members. For each DH star, we combined that relation and the 2MASS J magnitude to yield
a predicted IC. Figure 22 shows a plot of the difference of this predicted IC and I(DH) with
right ascension. The solid line shows the relation we adopt. Figure 23 shows the relation
between the corrected I(DH) values and Table 2 IC measures from photoelectric sources.
There is still a significant amount of scatter but the corrected I(DH) photometry appears
to be accurately calibrated to the Cousins system.
In a very few cases (specifically, just five stars), we provide an estimate of Ic based
on data from a wide-area CCD survey of Taurus obtained with the Quest-2 camera on the
Palomar 48 inch Samuel Oschin telescope (Slesnick et al. 2006). That survey calibrated their
photometry to the Sloan i system, and we have converted the Sloan i magnitudes to Ic. We
intend to make more complete use of the Quest-2 data in a subsequent paper.
When we have multiple sources of photometry for a given star, we consider how to com-
bine them. In most cases, if we have photoelectric data, that is given preference. However,
if we have photographic V and I, and only a photoelectric measurement for I, we do not
replace the photographic I with the photoelectric value because these stars are variable and
the photographic measurements are at least in some cases from nearly simultaneous expo-
sures. Where we have multiple sources for photoelectric photometry, and no strong reason
to favor one measurement or set of measurements over another, we have averaged the pho-
tometry for a given star. In most cases, where we have multiple photometry the individual
measurements agree reasonably well but with the caveat that the Pleiades low mass stars are
in many cases heavily spotted and “active” chromospherically, and hence are photometrically
variable. In a few cases, even given the expectation that spots and other phenomena may
affect the photometry, there seems to be more discrepancy between reported V magnitudes
– 27 –
than we expect. We note two such cases here. We suspect these results indicate that at least
some of the Pleiades low mass stars have long-term photometric variability larger than their
short period (rotational) modulation.
HII 882 has at least four presumably accurate V magnitude measurements reported in
the literature. Those measures are: V=12.66 Johnson & Mitchell (1958); V=12.95 Stauffer
(1982); V=12.898 van Leeuwen, Alphenaar, & Brand (1986); and V=12.62 Messina (2001).
HII 345 has at least three presumably accurate V magnitude measurements. Those
measurements are: V=11.65 Landolt (1979); V=11.73 van Leeuwen, Alphenaar, & Brand
(1986); V=11.43 Messina (2001).
At the bottom of Table 2, we provide a key to the source(s) of the optical photometry
provided in the table.
This research made use of the SIMBAD database operated at CDS, Strasbourg, France,
and also of the NED and NStED databases operated at IPAC, Pasadena, USA. A large
amount of data for the Pleiades (and other open clusters) can also be found at the open
cluster database WEBDA (http://www.univie.ac.at/webda/), operated in Vienna by Ernst
Paunzen.
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This preprint was prepared with the AAS LATEX macros v5.2.
– 32 –
Table 1. Pleiades Membership Surveys used as Sources
Reference Area Covered Magnitude Range Number Candidates Name
Sq. Deg. (and band) Prefix
Trumpler (1921) 3 2.5< B <14.5 174 Tr
Trumpler (1921)a 24 2.5< B <10 72 Tr
Hertzsprung (1947) 4 2.5< V <15.5 247 HII
Artyukhina (1969) 60 2.5< B <12.5 ∼200 AK
Haro et al. (1982) 20 11< V <17.5 519 HCG
van Leeuwen et al. (1986) 80 2.5< B <13 193 PELS
Stauffer et al. (1991) 16 14< V <18 225 SK
Hambly et al. (1993) 23 10< I <17.5 440 HHJ
Pinfield et al. (2000) 6 13.5< I <19.5 339 BPL
Adams et al. (2001) 300 8< Ks <14.5 1200 ...
Deacon & Hambly (2004) 75 10< R <19 916 DH
aThe Trumpler paper is listed twice because there are two membership surveys included in that
paper, with differing spatial coverages and different limiting magnitudes.
Fig. 1.— Spatial coverage of the six times deeper “2MASS 6x” observations of the Pleiades.
The 2MASS survey region is approximately centered on Alcyone, the most massive member
of the Pleiades. The trapezoidal box roughly indicates the region covered with the shallow
IRAC survey of the cluster core. The star symbols correspond to the brightest B star
members of the cluster. The red points are the location of objects in the 2MASS 6x Point
Source Catalog.
Fig. 2.— Color-magnitude diagram for the Pleiades derived from the 2MASS 6x obser-
vations. The red dots correspond to objects identified as unresolved, whereas the green
dots correspond to extended sources (primarily background galaxies). The lack of green
dots fainter than K = 16 is indicative that there is too few photons to identify sources as
extended - the extragalactic population presumably increases to fainter magnitudes.
Fig. 3.— As for Figure 2, except in this case the axes are J − H and H − Ks. The
extragalactic objects are very red in both colors.
Fig. 4.— FIGURE REMOVED TO FIT WITHIN ASTRO-PH FILESIZE GUIDELINES.
See http://spider.ipac.caltech.edu/staff/stauffer/pleiades07/ for full-res version. Two-color
(4.5 and 8.0 micron) mosaic of the central square degree of the Pleiades from the IRAC
survey. North is approximately vertical, and East is approximately to the left. The bright
star nearest the center is Alcyone; the bright star at the left of the mosaic is Atlas; and the
bright star at the right of the mosaic is Electra.
http://spider.ipac.caltech.edu/staff/stauffer/pleiades07/
Fig. 5.— Finding chart corresponding approximately to the region imaged with IRAC.
The large, five-pointed stars are all of the Pleiades members brighter than V= 5.5. The
small open circles correspond to other cluster members. Several stars with 8 µm excesses are
labelled by their HII numbers, and are discussed further in Section 6. The short lines through
several of the stars indicate the size and position angle of the residual optical polarization
(after subtraction of a constant foreground component), as provided in Figure 6 of Breger
(1986).
Fig. 6.— Comparison of aperture photometry for Pleiades members derived from the
IRAC 3.6 µm mosaic using the Spitzer APEX package and the IRAF implementation of
DAOPHOT.
Fig. 7.— Difference between aperture photometry for Pleiades members for IRAC channels
1 and 2. The [3.6]−[4.5] color begins to depart from essentially zero at magnitudes ∼10.5,
corresponding approximately to spectral type M0 in the Pleiades.
Fig. 8.— Ks vs. Ks −[4.5] CMD for Pleiades candidate members, illustrating why we have
excluded HII 1695 from the final catalog of cluster members. The “X” symbol marks the
location of HII 1695 in this diagram.
Fig. 9.— Spatial plot of the candidate Pleiades members from Table 2. The large star
symbols are members brighter than Ks= 6; the open circles are stars with 6 < Ks < 9; and
the dots are candidate members fainter than Ks= 9. The solid line is parallel to the galactic
plane.
Fig. 10.— The cumulative radial density profiles for Pleiades members in several magnitude
ranges: heavy, long dash – Ks < 6; dots – 6 < Ks < 9; short dash – 9 < Ks < 12; light, long
dash – Ks > 12.
Fig. 11.— V vs. (V − I)c CMD for Pleiades members with photoelectric photometry. The
solid curve is the “by eye” fit to the single-star locus for Pleiades members.
Fig. 12.— Ks vs. Ks − [3.6] CMD for Pleiades candidate members from Table 2 (dots) and
from deeper imaging of a set of Pleiades VLM and brown dwarf candidate members from
Lowrance et al. (2007) (squares). The solid curve is the single-star locus from Table 3.
Fig. 13.— V vs. (V − I)c CMD for Pleiades candidate members from Table 2 for which we
have photoelectric photometry, compared to theoretical isochrones from Siess et al. (2000)
(left) and from Baraffe et al. (1998) (right). For the left panel, the curves correspond to 10,
50, 100 Myr and a ZAMS; the right panel includes curves for 50, 100 Myr and a ZAMS.
Fig. 14.— K vs. (I −K) CMD for Pleiades candidate members from Table 2, compared to
theoretical isochrones from Siess et al. (2000) (left) and from Baraffe et al. (1998) (right).
The curves correspond to 50 Myr, 100 Myr and a ZAMS.
Fig. 15.— Ks vs. Ks−[3.6] CMD for the objects in the central one square degree of the
Pleiades, combining data from the IRAC shallow survey and 2MASS. The symbols are defined
within the figure (and see text for details). The dashed-line box indicates the region within
which we have searched for new candidate Pleiades VLM and substellar members. The solid
curve is a DUSTY 100 Myr isochrone from Chabrier et al. (2000), for masses from 0.1 to
0.03 M⊙.
Fig. 16.— Proper motion vector point diagrams (VPDs) for various stellar samples in the
central one degree field, derived from combining the IRAC and 2MASS 6x observations.
Top left: VPD comparing all objects in the field (small black dots) to Pleiades members
with 11 < Ks < 14 (large blue dots). Top right: same, except the blue dots are the new
candidate VLM and substellar Pleiades members. Bottom left: same, except the blue dots
are a nearby, low mass field star sample from a box just blueward of the trapezoidal region
in 15. Bottom right: VPD just showing a second, distant field star sample as described in
the text.
Fig. 17.— Same as Fig. 15, except that the new candidate VLM and substellar objects from
Table 4 are now indicated as small, red squares.
Fig. 18.— Two plots intended to isolate Pleiades members with excess and/or extended 8
µm emission. The plot with [3.6]−[8.0] micron colors shows data from Table 3 (and hence is
for aperture sizes of 3 pixel and 2 pixel radius, respectively). The increased vertical spread
in the plots at faint magnitudes is simply due to decreasing signal to noise at 8 µm. The
numbers labelling stars with excesses are the HII identification numbers for those stars.
Fig. 19.— Postage stamp images extracted from individual, 8 µm BCDs for the stars with
extended 8 µm emission, from which we have subtracted an empirical PSF. Clockwise from
the upper left, the stars shown are HII1234, HII859, Merope and HII652. The five-pointed
star indicates the astrometric position of the star (often superposed on a few black pixels
where the 8 µm image was saturated. The circle in the Merope image is centered on the
location of IC349 and has diameter about 25” (the size of IC349 in the optical is of order
10” x 10”).
Fig. 20.— Aperture growth curves from the 8 µm mosaic for stars with 24 µm excesses from
Gorlova et al. (2006) and for a set of control objects (dashed curves). All of the objects have
been scaled to common zero-point magnitudes for 9 pixel apertures, with the 24 µm excess
stars offset from the control objects by 0.1 mag. The three Gorlova et al. (2006) stars with
no excess at 8 µm are HII 996, HII 1284 and HII 2195. The Gorlova et al. (2006) star with
a slight excess at 8 µm is HII 489.
Fig. 21.— Calibration derived relating Ikp from Pinfield et al. (2000) and IC. The dots are
stars for which we have both Ikp and IC measurements (small dots: photographic IC; large
dots: photoelectric IC), and the solid line indicates the piecewise linear fit we use to convert
the Ikp values to IC for stars for which we only have Ikp.
Fig. 22.— Difference between the predicted IC and Deacon & Hambly (2004) I magnitude
as a function of right ascension for the DH stars. No obvious dependence is present versus
declination.
Fig. 23.— Comparison of the recalibrated DH I photometry with estimates of IC for stars
in Table 2 with photoelectric data.
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	Introduction
	New Observational Data
	2MASS ``6x" Imaging of the Pleiades
	Shallow IRAC Imaging
	A Catalog of Pleiades Candidate Members
	Empirical Pleiades Isochrones and Comparison to Model Isochrones
	Identification of New Very Low Mass Candidate Members
	Mid-IR Observations of Dust and PAHS in the Pleiades
	Pleiades Dust-Star Encounters Imaged with IRAC
	Circumstellar Disks and IRAC
	Summary and Conclusions
	APPENDIX
	Membership Catalogs
	Photometry
