Proper motion L and T dwarf candidate members of the Pleiades
Mon. Not. R. Astron. Soc. 000, 1–11 (2007) Printed 3 December 2018 (MN LATEX style file v2.2)
Proper motion L and T dwarf candidate members of the Pleiades
S. L. Casewell1⋆, P. D. Dobbie1,2, S. T. Hodgkin3, E. Moraux4, R. F. Jameson1,
N. C. Hambly5, J. Irwin3 and N. Lodieu6,1
1Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK
2Anglo-Australian Observatory, PO Box 296, Epping NSW 1710 Australia
3CASU, Institute of Astronomy,University of Cambridge, Maddingley Road, Cambridge, CB3 0HA, UK
4Laboratoire d’Astrophysique, Observatoire de Grenoble, Université Joseph Fourier, BP 53, 38041 Grenoble Cedex 9, France
5Scottish Universities Physics Alliance (SUPA), Institute for Astronomy, School of Physics, University of Edinburgh,
Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ
6 Instituto de Astrofı́sica de Canarias, Vı́a Láctea s/n, E-38205 La Laguna, Tenerife, Spain
January 2007
ABSTRACT
We present the results of a deep optical-near-infrared multi-epoch survey covering 2.5 square
degrees of the Pleiades open star cluster to search for new very-low-mass brown dwarf mem-
bers. A significant (∼ 5 year) epoch difference exists between the optical (CFH12k I-, Z-
band) and near infrared (UKIRT WFCAM J-band) observations. We construct I,I-Z and Z,Z-J
colour magnitude diagrams to select candidate cluster members. Proper motions are computed
for all candidate members and compared to the background field objects to further refine the
sample. We recover all known cluster members within the area of our survey. In addition,
we have discovered 9 new candidate brown dwarf cluster members. The 7 faintest candidates
have red Z-J colours and show blue near-infrared colours. These are consistent with being L
and T-type Pleiads. Theoretical models predict their masses to be around 11MJup.
Key words: stars: low-mass, brown dwarfs, open clusters and associa-
tions:individual:Pleiades
1 INTRODUCTION
The initial mass spectrum (IMS), the number of objects manufac-
tured per unit mass interval, is an outcome of the star formation
process which can be constrained via observation. Consequently,
empirical determinations of the form of the IMS can be used to crit-
ically examine our theoretical understanding of the complexities of
star formation. In recent years there has been a particular empha-
sis on building a solid comprehension of the mechanisms by which
very-low-mass stars, brown dwarfs and free-floating planetary mass
objects form (e.g. Boss 2001; Bate 2004; Goodwin, Whitworth &
Ward-Thompson 2004; Whitworth & Goodwin 2005). Neverthe-
less, one key question which remains unanswered is what is the
lowest possible mass of object that can be manufactured by the star
formation process? From a theoretical stance, traditional models
predict that if substellar objects form like stars, via the fragmenta-
tion and collapse of molecular clouds, then there is a strict lower
mass limit to their manufacture of 0.007-0.010 M⊙ (Padoan &
Nordlund, 2002). This is set by the rate at which the gas can radiate
away the heat released by the compression (e.g. Low & Lynden-
Bell, 1976). However, in more elaborate theories, magnetic fields
could cause rebounds in collapsing cloud cores which might lead
⋆ E-mail: slc25@star.le.ac.uk
to the decompressional cooling of the primordial gas, a lowering of
the Jeans mass and hence the production of gravitationally bound
fragments with masses of only ∼0.001 M⊙ (Boss, 2001). In con-
trast, if feedback from putative winds and outflows driven by the
onset of deuterium burning play a role, the smallest objects which
form via the star formation process may be restricted to masses
equal to or greater than the deuterium burning limit (∼0.013 M⊙;
Adams & Fatuzzo, 1996).
Recent work on very young clusters (τ<10 Myrs) and star
formation regions e.g. σ-Orionis, the Trapezium, IC348 and Up-
per Sco (Béjar et al., 2001; Muench et al., 2002; Muench et al.,
2003; Lodieu et al., 2007a) suggests that the initial mass function
continues slowly rising down to masses of the order M∼0.01 M⊙,
at least in these environments. Indeed, it has been claimed that an
object with a mass as low as 2-3 MJup has been unearthed in σ-
Ori (Zapatero-Osorio et al, 2002). However, the cluster member-
ship of σ-Ori 70 is disputed by Burgasser et al. (2004). Further-
more, mass estimates for such young substellar objects derived by
comparing their observed properties to the predictions of theoreti-
cal evolutionary tracks remain somewhat controversial. Baraffe et
al. (2002) have shown that to robustly model the effective temper-
ature and luminosity of a low mass object with an age less than
∼1 Myr, evolutionary calculations need to be coupled to detailed
simulations of the collapse and accretion phase of star formation.
c© 2007 RAS
http://arxiv.org/abs/0704.1578v2
2 S. L. Casewell et al.
As the current generation of evolutionary models start from arbi-
trary initial conditions, theoretical predictions for ages less than a
few Myrs must be treated with a fair degree of caution. Indeed, the
few available dynamical mass measurements of pre-main sequence
objects indicate that models tend to underestimate mass by a few
tens of percent in the range 0.3<∼M<∼1.0 M⊙ (see Hillenbrand &
White, 2004 for review). A recent dynamical mass measurement of
the 50-125 Myrs old object AB Dor C (spectral type ∼M8), the first
for a pre-main sequence object with M<0.3 M⊙, suggests that the
discrepancy between model predictions and reality might be even
larger at lower masses, with the former underestimating mass by a
factor 2-3 at M∼0.1 M⊙ (Close et al., 2005). However, this conclu-
sion is dependent on the assumed age of AB Dor, which is currently
a matter of great contention (Luhman, Stauffer & Mamajek, 2005;
Janson et al., 2006). On the positive side, Zapatero-Osorio et al.,
(2004) have determined the masses of the brown dwarf binary com-
ponents of GJ 569 Bab and their luminosities and effective temper-
atures are in agreement with theoretical predictions, for an age of
300 Myr. More recently, Stassun, Mathieu & Valenti (2006) discuss
an eclipsing brown dwarf binary in the Orion nebula star forming
region and find the large radii predicted by theory for a very young
dwarf. Surprisingly, they find that the secondary is hotter than the
more massive primary. Clearly further work is still needed to sup-
port the predictions of theoretical models.
It is clearly important to search for the lowest mass objects,
not only in the young clusters, but also in more mature clusters,
such as the Pleiades. The results of previous surveys of the Pleiades
indicate that the present day cluster mass function, across the stel-
lar/substellar boundary and down to M∼0.02 M⊙ (based on the
evolutionary models of the Lyon Group), can be represented by a
slowly rising power law model, dN/dM∝M−α . For example, from
their Canada-France-Hawaii Telescope (CFHT) survey conducted
at R and I and covering 2.5 sq. degrees, Bouvier et al. (1998) identi-
fied 17 candidate brown dwarfs (IC>17.8) and derived a power law
index of α=0.6. From their 1.1 sq degrees Isaac Newton Telescope
(INT) survey conducted at I and Z, with follow-up work undertaken
at K, Dobbie et al. (2002) unearthed 16 candidate substellar mem-
bers and found a power law of index α=0.8 to be compatible with
their data. Jameson et al. (2002) showed that a powerlaw of index
α=0.41±0.08 was consistent with the observed mass function over
the range 0.3>∼M>∼0.035 M⊙. This study used a sample of 49 prob-
able brown dwarf members assembled from the four most extensive
CCD surveys of the cluster available at the time, the International
Time Project survey (Zapatero Osorio et al., 1998), the CFHT sur-
vey (Bouvier et al., 1998; Moraux, Bouvier & Stauffer, 2001), the
Burrell Schmidt survey (Pinfield et al., 2000) and the INT survey
(Dobbie et al., 2002). The CFHT survey was subsequently extended
to an area of 6.4 sq. degrees (at I and Z) and unearthed a total of
40 candidate brown dwarfs. Moraux et al. (2003) applied statistical
arguments to account for non-members in their sample and derived
a power law index of α=0.6. Most recently, Bihain et al. (2006)
have used deep R, I, J and K band photometry and proper motion
measurements to unearth 6 robust L type Pleiades members in an
area of 1.8 sq. degrees with masses in the range 0.04-0.02 M⊙ and
derived a power law index of α=0.5±0.2.
Here we report the results of a new optical/infrared survey of
2.5 sq. degrees of the Pleiades, the aim of which is to extend empir-
ical constraints on the cluster mass function down to the planetary
mass regime (M∼0.01 M⊙). In the next section we describe the ob-
servations acquired/used as part for this study, their reduction, their
calibration and their photometric completeness. In subsequent sec-
tions we describe how we have identified candidate brown dwarf
members on the basis of colours and proper motions. We use our
new results to constrain the form of the cluster mass function and
conclude by briefly discussing our findings in the context of star
formation models.
2 OBSERVATIONS, DATA REDUCTION AND SURVEY
COMPLETENESS
2.1 The J band imaging and its reduction
Approximately 3.0 square degrees of the Pleiades cluster was ob-
served in the J band using the Wide Field Camera (WFCAM)
on the United Kingdom Infrared Telescope (UKIRT) between the
dates of 29/09/2005 and 08/01/2006. WFCAM is a near infrared
imager consisting of 4 Rockwell Hawaii-II (HgCdTe 2048x2048)
arrays with 0.4” pixels, arranged such that 4 separate pointings
(pawprints) can be tiled together to cover a 0.75 sq. degree region
of sky (see http://www.ukidss.org/technical/technical.html for dia-
gram). A total of four tiles were observed in a mixture of photo-
metric and non-photometric conditions but in seeing of typically ≈
1.0 arcsecond or better. To ensure that the images were properly
sampled we employed the 2×2 microstep mode. The locations on
the sky of our tiles (shown in Figure 1) were chosen to provide
maximum overlap with the optical fields surveyed in 2000 by the
Canada-France-Hawaii telescope and CFH12k camera but also to
avoid bright stars and areas of significant interstellar extinction.
The images were reduced at the Cambridge Astronomical Sur-
vey Unit (CASU) using procedures which have been custom writ-
ten for the treatment of WFCAM data. In brief, each frame was de-
biased, dark corrected and then flat fielded. The individual dithered
images were stacked before having an object detection routine run
on them. The detection procedure employs a core radius of 5 pixels,
and identifies objects as islands of more than 4 interconnected pix-
els with flux >1.5σ above the background level. The frames were
astrometrically calibrated using point sources in the Two micron
All Sky Survey (2MASS) catalogue. These solutions, in general,
had a scatter of less than 0.1 arcseconds. The photometric calibra-
tion employed by the CASU pipeline also relies on 2MASS data
(there are typically hundreds of 2MASS calibrators per detector)
and is found to be accurate to ≈2% in good conditions (see Warren
et al., 2007, Hodgkin et al., 2007 for details).
In measuring our photometry we used an aperture of 2”, which
is approximately twice the core radius of point sources. This 2” di-
ameter of the aperture is also twice the seeing FWHM. The reduc-
tion pipeline also attempts to classify each source depending on its
morphology (e.g. galaxy, star, noise). However, at the limit of the
data this classification becomes less reliable. Therefore, in our sub-
sequent analysis we chose to define as stellar all objects which lie
within 3 sigma of the stellar locus, where sigma is defined accord-
ing to Irwin et al. (in prep).
2.2 The far-red optical imaging and a new reduction
As part of this work we have used a subset (2.54 square degrees)
of the far-red optical data obtained in the course of the IZ survey of
the Pleiades conducted in 2000 by Moraux et al. (2003). The rele-
vant CFH12k data were extracted from the Canadian Astrophysical
Data Center archive and were reprocessed at Cambridge University
using the CASU optical imaging pipeline (Irwin & Lewis, 2001).
In brief, these data were bias subtracted and corrected for non-
linearity prior to flat fielding. Fringe maps, which were constructed
c© 2007 RAS, MNRAS 000, 1–11
http://www.ukidss.org/technical/technical.html
Proper motion L and T dwarf candidate members of the Pleiades 3
Figure 1. The regions imaged at I, Z and J with the CFHT and UKIRT. The CFH12k pointings (light rectangular outlines) are labelled alphabetically as in
Moraux et al. (2003), while the WFCAM tiles (bold square outlines) are labelled numerically, ranging from 1 to 4. Note that the observations avoid the region
of high reddening to the south of the Merope and the bright stars in vicinity of the cluster centre.
for each photometric band from images obtained during the observ-
ing run, were used to remove the effects of interference between
night sky lines in the CCD substrate. Subsequently, sources at a
level of significant of 3σ or greater were morphologically classified
and aperture photometry obtained for each. A World Coordinate
System (WCS) was determined for each frame by cross-correlating
these sources with the Automated Plate Measuring (APM) machine
catalogue (Irwin, 1985). The approximately 100 common objects
per CCD chip lead to an internal accuracy of typically better than
0.3 ”. The photometry was calibrated onto a CFH12k I and Z natu-
ral system using stars with near zero colour (B-V-R-I≈0) in Landolt
standard field SA98 (Landolt, 1992) which was observed the same
nights as the science data. The systematic errors in the photome-
try were calculated by comparing the photometry of overlapping
fields as in Moraux et al. (2003). The photometry was found to be
accurate to ≈3%.
2.3 The completeness of datasets
To estimate the completeness of our IR images, we injected fake
stars with magnitudes in the range J=12-22 into each of the 16
chips of every WFCAM frame and re-ran the object detection soft-
ware with the same parameters that were used to detect the real
sources. To avoid significantly increasing the density of all sources
in the data we inserted only 200 fake stars per chip in a given run.
To provide meaningful statistics we repeated this whole procedure
ten times. Subsequently, we calculated percentage completeness at
a given magnitude by taking the ratio of the number of fake stars
recovered to the number of fake stars injected into a given magni-
tude bin (and multiplying by 100). We note that a 100% recovery
rate was never achieved at any magnitude since a small proportion
of the fake stars always fell sufficiently close to other sources to
be overlooked by the object detection algorithm. This method was
also applied to determine the completeness of the I and Z band
CFH12k data. However, the magnitude range of the fake stars was
adjusted to be consistent with the different saturation and faint end
magnitude limits of these data. The results of this procedure for all
3 photometric bands are shown in Table 1.
A glance at this table indicates that the IR data are in general
90% complete to J≈19.7, although Field 3 is slightly less deep,
due to moonlight and poor seeing. In this case the proximity of
the moon led to higher background counts. The I data are typically
90% and 50% complete to I=22.5 and 23.5 respectively. The cor-
responding completeness limits for the Z band data are Z=21.5 and
22.5 respectively.
3 ANALYSIS OF THE DATA
3.1 Photometric selection of candidate cluster members
An initial photometrically culled sample of candidate brown dwarfs
has been obtained from the I,I-Z colour-magnitude diagram (Fig-
c© 2007 RAS, MNRAS 000, 1–11
4 S. L. Casewell et al.
 0  1  2  3  4  5
Figure 2. The I,I-Z CMD for the whole of field 1. The solid line is the NEXTGEN model, and the dotted line the DUSTY model. The small points are all
objects that were classed as stellar in both I and Z data. The crosses are all objects that met the following selection criteria: classed as stellar in both I and Z
data, for 16.5 < I < 22.5, they must lie no more than 0.25 magnitudes to the left of the DUSTY isochrone, for I>22.5, they must lie to the right of the line,
I-Z= (I-19.0)/3.5. The filled squares are the previously identified cluster candidate members from Bihain et al. (2006), Moraux et al. (2003) and Bouvier et al.
(1998), plotted to highlight the cluster sequence.
 0  0.5  1  1.5  2  2.5  3  3.5  4
Figure 3. The Z,Z-J CMD for the whole of the survey. The solid line is the NEXTGEN model, and the dotted line the DUSTY model. The crosses are all the
objects selected from the I,I-Z (crosses on Figure 2.). The filled diamonds are all objects that met our selection criteria from the I,I-Z,and Z, Z-J CMDs. These
were selected for proper motion analysis, and were found to be non members. The filled squares are our candidate cluster members (objects that remained
after proper motion analysis). The squares are our ZJ only candidates for all four fields that remained after proper motion analysis. The previously identified
probable members from Bihain et al. (2006), Moraux et al. (2003) and Bouvier et al. (1998) that remained after our proper motion analysis are identified by
open circles around the plotted symbols.
c© 2007 RAS, MNRAS 000, 1–11
Proper motion L and T dwarf candidate members of the Pleiades 5
Table 1. 50 and 90% completeness figures for the optical and infrared fields. The positioning of these fields is shown in Figure 1. Note that while WFCAM
field 1 corresponds to CFHT fields B, C, R and Q, the individual pawprints, do not correspond on a one to one basis - i.e. field1 00 does not correspond to field
Field name I Z WFCAM tile name WFCAM pawprint name J
50% 90% 50% 90% 50% 90%
B 23.2 22.5 22.3 21.5 field1 00 20.9 19.9
C 23.7 22.6 22.6 21.6 field1 01 20.9 20.1
R 24.0 23.0 22.9 21.6 field1 10 20.9 19.8
Q 23.7 22.5 22.7 21.6 field1 11 20.9 19.7
K 23.6 22.5 23.0 21.9 field2 00 20.9 19.7
L 24.0 22.7 23.0 21.8 field2 01 20.9 19.9
D 23.7 22.4 23.0 21.7 field2 10 21.0 19.9
field2 11 20.9 19.7
U 23.5 22.5 22.9 21.7 field3 00 19.5 18.8
V 23.8 22.5 22.7 21.7 field3 01 19.0 17.7
T 23.6 22.5 22.6 21.5 field3 10 19.6 18.6
field3 11 18.9 17.7
I 23.9 22.3 23.1 22.0 field4 00 20.8 19.7
G 23.7 22.7 23.4 22.3 field4 01 20.8 19.7
field4 10 20.8 19.7
field4 11 20.8 19.7
 1  2  3  4  5  6  7
Figure 4. The J,I-J CMD for the whole of the survey. The solid line is the NEXTGEN model, and the dotted line the DUSY model. The crosses are all the
objects selected from the I,I-Z (crosses on Figure 2.). The filled diamonds are all objects that met our selection criteria from the I,I-Z,and Z, Z-J CMDs. These
were selected for proper motion analysis, and were found to be non members. The filled squares are our candidate cluster members (objects that remained
after proper motion analysis). The previously identified members from Bihain et al. (2006), Moraux et al. (2003) and Bouvier et al. (1998) that remained after
our proper motion analysis are identified by open circles around the plotted symbols.
ure 2) where the 120 Myr NEXTGEN (Baraffe et al., 1998) and
DUSTY (Chabrier et al., 2000) model isochrones (modified to take
into account the Pleiades distance of 134 pc e.g. Percival, Salaris
& Groenewegen, 2005) served as a guide to the location of the
Pleiades sequence. With the uncertainties in both the photometry
and the age of the cluster in mind, we selected all objects classed
as stellar in both the I and Z data, which in the magnitude range
16.5 < I < 22.5 lay no more than 0.25 magnitudes to the left of
the DUSTY isochrone. All the candidate Pleiads found by Moraux
et al. (2003) and Bihain et al. (2006) lay within ±0.25 magnitudes
of the DUSTY model. Thus our selection criterion is 0.25 magni-
tudes to the left of the DUSTY model. Below I=22.5, the DUSTY
model is not red enough to account for known field stars, and so
is inappropriate in this effective temperature regime. We have cal-
culated an approximate field star sequence from Tinney, Burgasser
& Kirkpatrick (2003) and Hawley et al (2002) and lowered it by
2 magnitudes. This results in the line I-Z= (I-19.0)/3.5. This se-
lection is conservative, and is particularly aimed at removing the
c© 2007 RAS, MNRAS 000, 1–11
6 S. L. Casewell et al.
bulk of the red tail of the background stars. Subsequently, the initial
list of candidates was cross-correlated with our J band photometric
catalogue (using a matching radius of 2 arcseconds) and a refined
photometrically culled sample obtained using the Z,Z-J colour-
magnitude diagram (Figure 3). These objects are also shown on the
J, I-J colour-magnitude diagram (Figure 4). As before, the 120 Myr
model isochrones were used as a guide to the location of the cluster
sequence. With the photometric uncertainties in mind, all candi-
dates with Z620 were retained. All candidates with 20<Z<21 and
Z-J>1.6 were also retained. Finally, all candidates with Z>21 and
Z-J>1.9 were retained. These constraints are conservative and are
based on the field L and T dwarfs sequence (Z-J>3, Chiu et al.,
2006) since the DUSTY models are known to be inappropriate in
this effective temperature regime. Since our survey is limited by the
depth of the I band data, all candidates with Z>20 and no I band
counterpart were also kept.
3.2 Refining the sample using astrometric measurements
To weed out non-members we have measured the proper motion of
each candidate brown dwarf, using the Z and J band data where
the epoch difference was 5 years. In this process, only objects lying
within 2 arcminutes of each candidate were chosen as potential as-
trometric reference stars. This compromise provided a sufficiently
large number of sources but at the same time minimised the effects
of non-linear distortions in the images. Furthermore, objects with
large ellipticity (>0.2), classed by the photometric pipeline as non-
stellar in the Z band data and with Z<16 or Z>20 were rejected.
This ensured that, in the main, the astrometric reference sources
were not of very low S/N in the J band or saturated in the optical
data. These criteria generally provided at least 20 suitable stars per
candidate brown dwarf.
Six coefficient transforms between the epoch 1 Z band im-
ages and the epoch 2 J band images were calculated using routines
drawn from the STARLINK SLALIB package. The iterative fitting
rejects objects having residuals greater than 3σ, where σ is robustly
calculated as the median of absolute deviation of the reference star
residuals, scaled by the appropriate factor (1.48) to yield an equiv-
alent RMS. Once the routine had converged the relative proper mo-
tions in pixels were calculated by dividing the fitting residuals of
each candidate by the epoch difference. For our data the epoch
difference is approximately 5 years. Subsequently, the astrometric
motion in milliarcseconds per year in RA and DEC was derived by
folding these values through the World Coordinate System trans-
form matrix of the relevant WFCAM image.
To estimate the errors on our proper motions measurements,
we have injected fake stars into both the Z and J band data, in a sim-
ilar fashion to that described in section 2.3. However, here we have
determined the difference between the inserted position and the
photometric pipeline estimate of the centroid of each star. Assum-
ing that the differences between these two locations are normally
distributed, we have divided the fake stars into 3 magnitude bins in
each photometric band (Z621, 23>Z>21, 24>Z>23, 21>J>17)
and fit 2d Gaussians to estimate the 1-sigma centroiding uncertainty
as a function of source brightness.
We find that in the Z band data, for objects with magnitudes
Z621, the centroiding uncertainty is equivalent to 3 mas yr−1 in
each axis, while for objects with 23>Z>21 this number increases
to 8 mas yr−1. For our faintest Z band objects, 24>Z>23, the cen-
troiding uncertainty is equivalent to 12 mas yr−1 in each axis. In
the J band data, for objects with magnitudes 21>J>17, the cen-
troiding uncertainty is equivalent to 5 mas yr−1 in each axis. Thus
for our brightest candidates (Z<21, J<19), the quadratic sum of
the Z and J band centroiding errors is less than or comparable to
the RMS of the residuals of the linear transform fit, which is typ-
ically 5-10 mas yr−1 in each axis. We adopt this latter quantity as
the proper motion uncertainty in both the RA and DEC directions
for these objects. It is worth noting at this point that both the stars
and brown dwarfs of the Pleiades appear to be in a state of dynam-
ical relation (e.g. Pinfield et al. 1998, Jameson et al. 2002), where
the velocity dispersion of the members is proportional to 1/M0.5,
where M is mass. Based on an extrapolation of the data in Figure
4 of Pinfield et al. (1998), we would expect our lowest mass brown
dwarf members (0.01-0.02M⊙) to have velocity dispersion of ∼ 7
mas yr−1. This velocity dispersion should be added quadratically
to the above uncertainties. Our final adopted proper motion selec-
tion, effectively a radius of 14 mas yr−1 , is described below, and
the velocity dispersion is small compared to this.
We fitted the proper motions of all of our photometric candi-
dates (excluding the ZJ only candidates) with a 2D Gaussian, which
centred around 1.1, -7 mas yr−1. This Gaussian had a σ of 14.0. We
were not able to fit two Gaussians, one to the background stars and
one to the Pleiades dwarfs, as described in Moraux et al.(2003),
since only ≈ 30 objects have the correct proper motion for cluster
membership. Consequently, we only selected objects to be proper
motion members if they had proper motions that fell within 1σ of
the proper motion of the cluster at +20.0, -40.0 mas yr−1 (Jones
1981; Hambly, Jameson & Hawkins, 1991; Moraux et al., 2001).
We required the selection criteria to be 1σ, as extending this to 2σ,
would seriously overlap with the field stars centred on 0,0. We did
however extend the selection criteria to 1.5σ, which yielded 14 ad-
ditional objects, however all were rejected due to their bright, but
blue (I≈17.0, I-Z<1.0) positions on the I,I-Z CMD, which led us
to believe that they were field objects. We also attempted to tighten
our selection criteria to a circle with radius 10 mas yr−1. This se-
lection meant that we lost as possible members objects PLZJ 78,
9, 77, 23 (see Table 4). PLIZ 79, 9 and 77 have all been identified
and confirmed as proper motion members by Bihain et al. (2006),
Moraux et al. (2001), and Bouvier et al. (1998). Unfortunately, as
we cannot fit two Gaussians to our data, we cannot calculate a prob-
ability of membership for these objects by the standard method as
defined by Sanders (1971). The proper motion measurements may
be found in Table 4, as well as the I, Z, J, H and K magnitudes for
these candidate members to the cluster.
We have attempted to use control data to determine the level
of contamination within our data, however, the numbers involved
are very small, so any calculated probability will be rather uncer-
tain. We used as controls, two circles of radius 14.0 mas yr−1, at
the same distance from 0,0 proper motion as the Pleiades. We then
separated the data into one magnitude bins, and calculated the prob-
ability for each magnitude bin, using equation 1.
Pmembership =
Ncluster −Ncontrol
Ncluster
Where Pmembership is the probability of membership for that mag-
nitude bin, Ncluster is the number of stars and contaminants within
the cluster circle in that magnitude bin. Ncontrol is the number of
dwarfs in the control circle of proper motion space, see Figure 4.
Ncluster - Ncontrol is the number of Pleiads. It can be seen that the
probability depends on where the control circle is located. Thus as
well as using control circles, we use an annulus and scale down the
count to an area equal to that of a control circle. Note that Figure
5 is for all of the magnitude bins together. Figure 6 is the same
as Figure 5, but for the ZJ selected objects only. The statistics are
c© 2007 RAS, MNRAS 000, 1–11
Proper motion L and T dwarf candidate members of the Pleiades 7
-100 -50  0  50  100
µαcos(δ) mas/yr
Figure 5. Proper motion vector diagram of the photometrically selected
candidate members. The filled triangles are candidate and known cluster
members. The filled diamonds and filled circles are the two separate control
clusters used. The annulus used for the radial method is also plotted.
-100 -50  0  50  100
µαcos(δ) mas/yr
Figure 6. Proper motion vector diagram of the photometrically selected
candidate members. The filled triangles are candidate cluster members se-
lected from the Z,Z-J CMD only. The filled diamonds and filled circles
are the two separate control clusters used. The annulus used for the radial
method is also plotted.
much poorer for the individual magnitude bins and the probabili-
ties are correspondingly more uncertain. It can be seen in Figure
4 that there is not a symmetrical distribution of proper motions.
In fact the distribution in the Vector point diagram, is a classical
”velocity ellipsoid” displaced from zero by reflex motion from the
Sun’s peculiar velocity, and happens to be in the direction of the
Pleiades proper motion vector. We have therefore probably under-
estimated the contamination, as the annulus method of calculating
probabilities assumes that the vector point diagram has a circularly
symmetric distribution of objects. These probabilities are shown in
Table 2, and probabilities derived in the same way but for the ZJ
only candidates can be found in Table 3.
An alternative approach to estimating the contamination is
the use the field L and T dwarf luminosity functions. Chabrier
(2005) gives the T dwarf luminosity function as being 10−3
dwarfs/pc3/unit J mag interval. Our 7 L and T dwarf candidates
cover a total of 0.7 mag in the J band. Note PLZJ 323 and 23 may
be late L dwarfs but we include them in this analysis. The volume of
space we use is 836 pc2, based on 2.5 square degrees and a distance
to the Pleiades of 134±30 pc (Percival et al., 2005). This distance
range corresponds to a distance modulus range of ±0.5 magni-
tudes, which is generous, given that the sequence shown in figure 8
is clearly narrower than ±0.5 magnitudes. Thus the expected num-
ber of contaminating field dwarfs is 0.6. In addition to this, field T
dwarfs are unlikely to have the same proper motion as the Pleiades,
thus reducing the 0.6 further. For the field L dwarfs with MJ≈13.0
(i.e. J≈18.5 at the distance of the Pleiades) the luminosity func-
tion is 3×10−4 dwarfs/pc3/unit J mag interval (Chabrier, 2005). A
similar calculation then gives 0.25 contaminating L dwarfs which
should be further reduced by considering proper motions. It is thus
clear that the field luminosity function indicates that contamination
by field L and T dwarfs should be negligible.
4 RESULTS
Most of these objects, except two bright objects and the faintest
seven have been documented before in surveys - Moraux et al
(2003) and Bihain et al (2006). We recovered all of these objects
within our overlapping area, and none were rejected by our IZ pho-
tometric selection. The objects we recovered were BRB 4, 8, 17, 13,
19, 21, 22, 27 and 28 and PLIZ 2, 3, 5, 6, 13, 14, 19, 20, 26, 28, 31,
34, 35 and 36. PLIZ 18, 27 and 39 were found to have no J counter-
part in our catalogues. Of these objects, BRB 19 and PLIZ 14 and
26 met by our selection criteria on the Z, Z-J CMD, however they
were too blue in their Z-J colour for their place on the sequence.
Out of the remaining objects we find that we agree with the proper
motion measurements as calculated by Bihain et al.(2006) for PLIZ
28, which we believe is a member of the cluster. We agree with Bi-
hain et al.(2006) over their candidates BRB 13 and BRB 19 that
they are not proper motion members to the cluster, however we dis-
agree with their proper motion measurement for BRB 19. We also
find that PLIZ 5 is a non member to the cluster - ie its proper mo-
tion measurement is not within 14 mas yr−1 of the cluster proper
motion value. We find that PLIZ 14 and 26 are not proper motion
members to the cluster, as well as not having met our selection cri-
teria. PLIZ 26 was found to have a proper motion measurement of
35.73±9.00, -25.83±6.96, which did not fall within 14 mas y−1
of the cluster, and also missed the selection made with the wider
circle (21 mas yr−1) as well. We find that PLIZ 19, 20, 34 and 36
are not proper motion members to the cluster. However this means
we disagree with Moraux et al. (2003), over their object PLIZ 20.
They find a proper motion of 25.6±7.3, -44.7±7.4 mas yr−1 for it.
Our proper motion measurement is 0.88± 15.86, -0.92±8.42 mas
yr−1. It is possible that this object has been adversely affected by
its position on the edge of one of the WFCAM chips, thus reducing
the number of reference stars used to calculate its proper motion.
An alternative method of measuring the proper motion using all the
objects on the same chip produced a measurement of 19.14±11.06,
-28.989±11.94 mas yr−1. This value does meet our selection cri-
teria, and has been previously accepted as a member. We suggest
PLIZ 20 is likely to be a member because of this.
We find that PLIZ 2, 3, 6, 31 and 35 are all proper motion
members to the cluster. In addition to this, we find 2 brighter new
candidate members to the cluster. These objects are bright enough
to have appeared in previous surveys, and in the UKIDSS Galactic
cluster survey (GCS). We also have 2 fainter new members to the
c© 2007 RAS, MNRAS 000, 1–11
Table 4. Name,coordinates, Z, I, J, H and K magnitudes for our members to the cluster. The errors quoted are internal (from photon counting). The systematic calibration errors are 2% in
the J, H and K wavebands (Warren et al., 2007), and 3% in the I and Z wavebands. The J, H and K magnitudes are on the MKO system. Previously discovered members also also have their
other known names listed from Moraux et al. (2003), Bihain et al. (2006) and Bouvier et al. (1998). The H and K band magnitudes are taken from the UKIDSS Galactic Cluster Survey with
the exceptions of PLZJ 23, 93, 721 and 235 which have their H band magnitudes listed from our H survey. The K band magnitude for PLZJ 93 is from our UFTI photometry, and PLZJ 23 is
from LIRIS service time. The final 5 objects in the table are our candidates selected from the ZJ data only.
Name Alternate RA dec µαcosδ µδ I Z J H K
name J2000.0 mas yr−1
PLZJ 29 BRB4 03 44 23.23 +25 38 45.11 23.40±8.24 -48.51±6.34 17.005 ± 0.001 16.163 ± 0.001 14.732 ± 0.001 14.132±0.004 13.744±0.004
PLZJ 56 03 44 53.51 +25 36 19.46 19.68±7.34 -35.63±5.29 17.012 ± 0.001 16.351 ± 0.001 15.250 ± 0.001 14.650±0.005 14.342±0.006
PLZJ 45 BRB8, CFHT-PL-7 03 52 58.2 +24 17 31.57 19.72±4.95 -42.37±7.44 17.101 ± 0.001 16.417 ± 0.001 15.247 ± 0.001 14.614±0.005 14.251±0.006
PLZJ 50 03 43 55.98 +25 36 25.45 13.48±8.24 -35.65±5.38 17.239 ± 0.001 16.496 ± 0.001 15.268 ± 0.001 14.693±0.006 14.319±0.006
PLZJ 60 CFHT-PL-10 03 44 32.32 +25 25 18.06 16.93±7.76 -43.15±5.72 17.592 ± 0.001 16.810 ± 0.001 15.460 ± 0.001 14.884±0.007 14.465±0.006
PLZJ 78 PLIZ2 03 55 23.07 +24 49 05.18 19.72±10.06 -29.74±10.45 17.719 ± 0.001 16.948 ± 0.001 15.574 ± 0.001 14.963±0.007 14.552±0.007
PLZJ 46 PLIZ3, BRB11 03 52 67.20 +24 16 01.00 19.55±5.15 -42.58±7.57 17.742 ± 0.001 16.945 ± 0.001 15.583 ± 0.001 14.966±0.007 14.503±0.008
PLZJ 9 PLIZ6,BRB9 03 53 55.09 +23 23 36.38 24.13±13.83 -50.10±22.71 17.752 ± 0.001 16.804 ± 0.001 15.222 ± 0.001 14.548±0.005 14.054±0.005
PLZJ 11 PLIZ20 03 54 05.33 +23 33 59.71 9.14±11.06 -28.98±11.94 19.571±0.004 18.563±0.004 16.691±0.005 15.980±0.016 15.436±0.016
PLZJ 77 PLIZ28,BRB18 03 54 10.04 +23 17 52.28 12.01±14.59 -51.60±15.84 20.760 ± 0.010 19.728 ± 0.010 17.647 ± 0.010 16.789±0.031 16.131±0.030
PLZJ 21 PLIZ31 03 51 47.65 +24 39 59.18 17.84±9.41 -44.92±8.03 20.944±0.014 19.762±0.013 17.575±0.012 16.774±0.026 16.089±0.028
PLZJ 10 PLIZ35,BRB15 03 52 31.19 +24 46 29.61 15.84±8.88 -49.34±6.24 21.293±0.018 20.292±0.016 18.181±0.022 17.118±0.041 16.506±0.0416
PLZJ 4 BRB21 03 54 10.25 +23 41 40.67 29.74±13.17 -38.46±8.88 21.322 ± 0.010 20.215 ± 0.013 18.171 ± 0.010 17.141±0.045 16.377±0.039
PLZJ 61 BRB22 03 44 31.27 +25 35 14.97 25.82±7.89 -40.21±8.47 22.043 ± 0.030 20.782 ± 0.026 18.298 ± 0.020 17.393±0.059 16.657±0.04
PLZJ 32 BRB27 03 44 27.27 +25 44 41.99 25.03±11.52 -38.65±23.46 22.235 ± 0.040 20.962 ± 0.029 18.871 ± 0.030 17.793±0.094 16.950±0.070
PLZJ 37 BRB28 03 52 54.92 +24 37 18.85 18.13±11.53 -48.68±11.38 22.452 ± 0.05 21.216 ± 0.041 18.839 ± 0.030 17.742±0.071 16.921±0.058
PLZJ 23 03 51 53.38 +24 38 12.11 20.75±10.51 -50.05±9.96 23.541 ± 0.140 22.187 ± 0.112 19.960 ± 0.100 19.362±0.100 18.510±0.030
PLZJ 93 03 55 13.00 +24 36 15.8 13.11±14.36 -33.77±12.97 24.488 ± 0.370 22.592 ± 0.164 19.968 ± 0.080 19.955±0.100 19.420 ±0.100
PLZJ 323 03 43 55.27 +25 43 26.2 29.87±12.05 -39.37±11.70 - 21.597±0.054 19.613±0.076 - -
PLZJ 721 03 55 07.14 +24 57 22.34 19.18±22.23 -40.70±12.38 - 22.195±0.092 20.248±0.116 20.417±0.123 -
PLZJ 235 03 52 32.57 +24 44 36.64 20.92±12.16 -45.84±11.75 - 22.339±0.115 20.039±0.112 20.245±0.127 -
PLZJ 112 03 53 19.37 +24 53 31.85 8.56±14.08 -34.59±19.99 - 22.532±0.116 20.281±0.143 - -
PLZJ 100 03 47 19.19 +25 50 53.3 20.23±14.27 -37.28±23.82 - 23.563±0.373 20.254±0.114 - -
000,1
Proper motion L and T dwarf candidate members of the Pleiades 9
Table 2. Probability of membership, magnitude range for our methods of calculating probabilities of membership using the annulus as well as the two control
areas.
Probability Probability Probability Magnitude range
annulus µαcosδ=-20 mas yr−1 µδ=-40 mas yr
−1 µαcosδ=+40 mas yr−1 µδ=-20 mas yr
0.67 0.25 0.0 16 - 17
0.82 0.66 0.0 17 - 18
0.88 1.00 0.0 18 - 19
0.84 1.00 0.0 19 - 20
1.00 1.00 1.00 20 - 21
0.88 0.50 1.00 21 - 22
0.61 1.00 0.00 22 - 23
Table 3. Probability of membership, magnitude range for our methods of calculating probabilities of membership using the annulus as well as the two control
areas for our candidates selected from the ZJ data only.
Probability Probability Probability Magnitude range
annulus µαcosδ=-20 mas yr−1 µδ=-40 mas yr
−1 µαcosδ=+40 mas yr−1 µδ=-20 mas yr
0.61 1.00 1.00 21 - 22
0.35 0.67 0.33 22 - 23
-0.16 -2.00 0.00 23 - 24
cluster, and 5 objects selected using the ZJ photometry only. All of
the objects identified as cluster members in this work are presented
in Table 4. Two WFCAM tiles, 1 and 4, (see Figure 1) also had
deep H band photometry. These tiles were observed at the same
time as the J band imaging, and were observed under the same
conditions, but with the exception that microstepping was not used.
These data were reduced using the same pipeline as the J band data,
but the photometry and object detection used a core radius of 2.5
pixels in this case. Fortunately these tiles also covered our faintest,
previously undiscovered Pleiades candidates, PLZJ 23 and PLZJ
93, as well as two of the candidates selected from the ZJ data only,
PLZJ 721 and 235.
The UKIDSS Galactic Cluster survey (GCS) has also covered
the entire area at J, H and K. The UKIDSS data are reduced using
the same pipeline as the WFCAM data (see Dye et al, 2006 for
details of the pipeline).
We also have used UKIRT service time to measure photom-
etry for PLZJ 93 in the K band. This observation was taken on
09/09/2006 in seeing of better than 1.1” using the UKIRT Fast
Track Imager (UFTI), with a five point dither pattern. The data
were reduced using the ORAC-DR pipeline, and the photometry
was calibrated using UKIRT Faint Standard 115.
The K band photometry for PLZJ 23 was obtained on the night
of 05/03/2007 using the long slit intermediate resolution spectro-
graph (LIRIS) on the William Hershel Telescope in service time,
using a nine point dither pattern in seeing of ≈ 0.9”. The data were
reduced using IRAF and astrometrically and photometrically cali-
brated using 2MASS. The colour transforms presented in Carpen-
ter, (2001) were used to calculate the K band magnitude from the
KS magnitude.
Thus we have I, Z, J, H and K band photometry for the major-
ity of our Pleiades candidates. However H or K band photometry
is still needed for PLZJ 323, 721, 235, 112 and 100, (see Table 4).
Figures 7 and 8 show the K, J-K and H, J-H, colour magnitude
diagrams, together with the NEXTGEN (Baraffe et al, 1998) and
DUSTY (Chabrier et al, 2000) models for the Pleiades age of 120
Myrs (Stauffer et al 1998). The candidate members listed in Table
4 are also plotted in Figures 3 and 4 for clarity. In both of these
diagrams the M dwarf tail, the redward L dwarf sequence and the
L to T blueward transition sequence are clear. The L-T transition
sequence of course only has two objects plotted on it on Figure 7
as we have no K band photometry for the ZJ candidates. As ex-
pected the K, J-K diagram gives the best differentiation between
the sequences. The redward L sequence in this diagram agrees with
that found by Lodieu et al,(2007b) derived from a much greater
area of the Pleiades by the UKIDSS GCS. The GCS is not sensi-
tive enough to see the L-T blueward transition sequence however.
The K, J-K diagram also shows the separation between single and
binary dwarfs quite clearly. Note that the DUSTY theoretical track
is too flat compared to our empirical sequence, see figures 7 and 8.
PLZJ 23, 93, 721 and 235 have J-H colours of 0.60, 0.00, -
0.17 and -0.21 respectively. Comparing these colours with the spec-
tral type colour relations of field dwarfs described in Leggett et al.
(2002), yields estimated spectral types of T1.5, T4.5, T6 and T6
respectively. PLZJ 93 has J-K=0.60 which gives a spectral type
of T3 (Leggett et al., 2002), which is consistent with the spectral
type derived from the J-H colour (T4.5), within the errors. We also
can calculate a H-K colour for this dwarf of 0.6, however the H-K
colour is not a good choice for spectral typing, for instance, H-
K=0.6 covers a range of spectral types from L1 to T3 (Chiu et al.,
2006). The Z-J colour is also not a good choice of colour for mea-
suring spectral types until the later T dwarfs (>T2)(Hawley et al.,
2002). PLZJ 23 has J-K=1.45, which gives a spectral type of be-
tween L8 and T1. The H-K colour for this dwarf is 0.85. We may
thus assume that PLZJ 23 has a spectral type between L8 and T1.5,
and likewise that PLZJ 93 has a spectral type of between T3 and
T5 to take into account the photometric errors. It should be noted
that the Z band quoted in Hawley et al., (2002) is for the Sloan
filter system, and so for this reason we have not chosen to use it
to spectral type our objects. We believe that the J-H colour gives
the best estimate available to us of spectral types. Two of the three
candidate members without H band photometry PLZJ 112 and 100
have fainter J magnitudes than PLZJ 23 and 93, and so it is likely
that they are also T dwarfs. PLZJ 323 is brighter and is therefore
c© 2007 RAS, MNRAS 000, 1–11
10 S. L. Casewell et al.
Table 4.
 0.4  0.6  0.8  1  1.2  1.4  1.6  1.8  2
Figure 7. The K,J-K CMD for our candidate cluster members. The solid
line is the NEXTGEN model of Baraffe et al (1998), and the dotted line
is the DUSTY model of Chabrier et al. (2000). The filled squares are the
candidates identified by Moraux et al. (2003), the filled triangles are the
candidates identified by Bihain et al. (2006), the object enclosed by the
open circle is CFHT-PL-10 identified by Bouvier et al. (1998). The objects
marked by small points are our new candidate members. One of our T dwarf
candidates, PLZJ 93, is found to the bottom of the plot, with a J-K of ≈ 0.6.
PLZJ 23 is also present with a J-K of 1.45.
probably a late L dwarf. Indeed our faintest candidate at Z, PLZJ
100, may be a very late T dwarf, however this assumption is made
using its Z-J colour, which is very red. Using J magnitudes and the
COND models of Baraffe et al. (2003) for 120 Myrs (the DUSTY
models are no longer appropriate for calculating masses for objects
this faint in the Pleiades), we calculate masses of ≈ 11 MJup for
PLZJ 23, 93, 323, 721, 235, 112 and 100. More photometry in the
H and K bands is clearly needed to improve and extend these esti-
mates of the spectral types.
5 MASS SPECTRUM
To calculate the mass spectrum, we first divided the sample into
single dwarfs or single dwarfs with possible low mass companions
and dwarfs that are close to 0.75 magnitudes above the single star
sequence in the K, J-K colour magnitude diagram. The latter we
assume to be equal mass binaries and count them as dwarfs with
masses the same as a dwarf on the single dwarf sequence below
them. From Figures 3, 4, 7 and 8 it can be seen that there are
2 such binaries all with J-K ≈ 1. Dwarfs with J-K <1.2 are as-
signed masses using their H magnitudes and the NEXTGEN mod-
els (Baraffe et al. 1998). For 1.2<J-K<2.0 we use the DUSTY
models (Chabrier et al., 2000) and the J-H colour to assign a mass.
Finally the T dwarf masses were calculated from their J magni-
tudes and the COND models (Baraffe et al., 2003). The masses
were binned into three mass intervals, covering the low, medium
and high mass ranges and the numbers per bin are weighted by
the probabilities of membership calculated using the annulus, and
the bin width has been taken into account. The candidate members
with negative probabilities are obviously omitted from the mass
spectrum. The resultant mass spectrum is shown in Figure 9. The
-0.4 -0.2  0  0.2  0.4  0.6  0.8  1  1.2
Figure 8. The H,J-H CMD for our candidate cluster members. The filled
squares are the candidates identified by Moraux et al. (2003), the filled trian-
gles are the candidates identified by Bihain et al. (2006), the object enclosed
by the open circle is CFHT-PL-10 identified by Bouvier et al. (1998). The
objects marked by small points are our new candidate members.The filled
diamonds are the two candidates with H magnitudes selected from the ZJ
data only. The solid line is the NEXTGEN model of Baraffe et al (1998),
and the dotted line is the DUSTY model of Chabrier et al. (2000).
errors are poissonian. Clearly the statistics are very poor, due to
the small number of objects being dealt with. Using linear regres-
sion we have fitted our data to the relationship dN/dM∝M−α , and
calculate α=0.35±0.31. This is lower but still in agreement with
values presented in the literature (within 1σ), however the error on
this value is large, and the statistics are poor due to the small num-
bers involved. If we take into account the fact that the last mass
bin is only 50% complete (using Tables 1 and 4), then the lowest
mass bin can be increased by 50% to compensate. If we then fit
these data, we derive a value for α of 0.62±0.14, which is in agree-
ment with the literature. Alternatively, we can discount this final
low mass bin as being incomplete and simply omit it from the fit.
In this case we calculate a value for α of 0.86. We have only dis-
played the mass spectrum for the cluster in the area and magnitude
surveyed. This is to avoid trying to take into account biases caused
by some areas being more studied than others, and also because we
are only adding a maximum of 9 objects to the mass spectrum, 7
of which have low probabilities of membership and small masses,
and so are not likely to affect previous results a large amount. The
mass spectrum appears to be rising towards the lowest masses, but
this is not statistically significant due to the large error bars.
6 CONCLUSIONS
We have confirmed a number of L dwarf candidates in the Pleiades.
However the main result in this paper is the discovery of seven L
and T dwarf Pleiads of masses ≈ 11 MJup, below the 13 MJup
deuterium burning limit that is often used, somewhat artificially as
the upper bound for planetary masses. Further H and K band pho-
tometry, currently lacking for some of these candidates, will im-
prove confidence in their membership of the cluster. Planetary mass
brown dwarfs have, of course, been claimed in the Orion nebula
c© 2007 RAS, MNRAS 000, 1–11
Proper motion L and T dwarf candidate members of the Pleiades 11
-1.9 -1.8 -1.7 -1.6 -1.5 -1.4 -1.3 -1.2 -1.1
log mass (Solar masses)
Figure 9. The mass spectrum for our Pleiades candidate members. The
mass bin is in units of M⊙. The solid line is the fit to the data,
(α=0.35±0.31).
(Lucas & Roche 2000) and in the σ-Ori cluster (Zapatero-Osorio et
al., 2002). These clusters both have very young ages and may also
have a spread of ages (Béjar et al., 2001), making mass determi-
nations somewhat uncertain. Lodieu et al. (2006, 2007a) have also
found planetary mass brown dwarfs in the Upper Scorpius Associ-
ation which has an age of 5 Myrs (Preibisch & Zinnecker, 2002).
At very young ages the theoretical models may have significant er-
rors when used to assign masses (Baraffe et al., 2002). Our result
is the first detection of planetary mass objects in a mature cluster
whose age is well established. It strengthens the case that the star
formation process can produce very low mass objects.
7 ACKNOWLEDGEMENTS
SLC, NL,and PDD acknowledge funding from PPARC. We also
acknowledge the Canadian Astronomy Data Centre, which is oper-
ated by the Dominion Astrophysical Observatory for the National
Research Council of Canada’s Herzberg Institute of Astrophysics.
This work has been based on observations obtained at the Canada-
France-Hawaii Telescope (CFHT) which is operated by the Na-
tional Research Council of Canada, the Institut National des Sci-
ences de l’Univers of the Centre National de la Recherche Scien-
tifique of France, and the University of Hawaii. Observations were
also made at the United Kingdom Infrared Telescope, which is op-
erated by the Joint Astronomy Centre on behalf of the U.K. Particle
Physics and Astronomy Research Council. This 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 In-
frared Processing and Analysis Center/California Institute of Tech-
nology, funded by the National Aeronautics and Space Administra-
tion and the National Science Foundation. This research has made
use of NASA’s Astrophysics Data System Bibliographic Services,
the WHT service programme, and the UKIRT service programme.
We would like to thank the referee V.J.S. Béjar for his comments
which have improved the paper.
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Mon. Not. R. Astron. Soc. 000, 1–?? (2009) Printed 3 December 2018 (MN LATEX style file v2.2)
Errata:Proper motion candidate members of the Pleiades
S. L. Casewell1⋆, P. D. Dobbie1,2, S. T. Hodgkin3, E. Moraux4, R. F. Jameson1,
N. C. Hambly5, J. Irwin3 and N. Lodieu6,1
1Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK
2Anglo-Australian Observatory, PO Box 296, Epping NSW 1710 Australia
3CASU, Institute of Astronomy,University of Cambridge, Maddingley Road, Cambridge, CB3 0HA, UK
4Laboratoire d’Astrophysique, Observatoire de Grenoble, Université Joseph Fourier, BP 53, 38041 Grenoble Cedex 9, France
5Scottish Universities Physics Alliance (SUPA), Institute for Astronomy, School of Physics, University of Edinburgh,
Royal Observatory, Blackford Hill, Edinburgh EH9 3HJ
6 Instituto de Astrofı́sica de Canarias, Vı́a Láctea s/n, E-38205 La Laguna, Tenerife, Spain
3 December 2018
The paper “Proper motion L and T dwarf candidate members
of the Pleiades” was published in the Monthly Notices of the Royal
Astronomical Society, 2007, 378, 1131. It has come to our attention
that there were errors in Table 4 as regards the Right Ascension of
the candidate coordinates. Table 4 should read as follows. This has
no impact on the scientific results presented in the paper.
c© 2009 RAS
http://arxiv.org/abs/0704.1578v2
Table 4Name,coordinates, Z, I, J, H and K magnitudes for our members to the cluster. The errors quoted are internal (from photon counting). The systematic calibration errors are 2% in the
J, H and K wavebands (Warren et al., 2007), and 3% in the I and Z wavebands. The J, H and K magnitudes are on the MKO system. Previously discovered members also also have their other
known names listed from Moraux et al. (2003), Bihain et al. (2006) and Bouvier et al. (1998). The H and K band magnitudes are taken from the UKIDSS Galactic Cluster Survey with the
exceptions of PLZJ 23, 93, 721 and 235 which have their H band magnitudes listed from our H survey. The K band magnitude for PLZJ 93 is from our UFTI photometry, and PLZJ 23 is from
LIRIS service time. The final 5 objects in the table are our candidates selected from the ZJ data only.
Name Alternate RA dec µαcosδ µδ I Z J H K
name J2000.0 mas yr−1
PLZJ 29 BRB4 03 44 23.23 +25 38 45.11 23.40±8.24 -48.51±6.34 17.005 ± 0.001 16.163 ± 0.001 14.732 ± 0.001 14.132±0.004 13.744±0.004
PLZJ 56 03 44 53.51 +25 36 19.46 19.68±7.34 -35.63±5.29 17.012 ± 0.001 16.351 ± 0.001 15.250 ± 0.001 14.650±0.005 14.342±0.006
PLZJ 45 BRB8, CFHT-PL-7 03 52 05.82 +24 17 31.57 19.72±4.95 -42.37±7.44 17.101 ± 0.001 16.417 ± 0.001 15.247 ± 0.001 14.614±0.005 14.251±0.006
PLZJ 50 03 43 55.98 +25 36 25.45 13.48±8.24 -35.65±5.38 17.239 ± 0.001 16.496 ± 0.001 15.268 ± 0.001 14.693±0.006 14.319±0.006
PLZJ 60 CFHT-PL-10 03 44 32.32 +25 25 18.06 16.93±7.76 -43.15±5.72 17.592 ± 0.001 16.810 ± 0.001 15.460 ± 0.001 14.884±0.007 14.465±0.006
PLZJ 78 PLIZ2 03 55 23.07 +24 49 05.18 19.72±10.06 -29.74±10.45 17.719 ± 0.001 16.948 ± 0.001 15.574 ± 0.001 14.963±0.007 14.552±0.007
PLZJ 46 PLIZ3, BRB11 03 52 06.71 +24 16 00.99 19.55±5.15 -42.58±7.57 17.742 ± 0.001 16.945 ± 0.001 15.583 ± 0.001 14.966±0.007 14.503±0.008
PLZJ 9 PLIZ6,BRB9 03 53 55.09 +23 23 36.38 24.13±13.83 -50.10±22.71 17.752 ± 0.001 16.804 ± 0.001 15.222 ± 0.001 14.548±0.005 14.054±0.005
PLZJ 11 PLIZ20 03 54 05.33 +23 33 59.71 9.14±11.06 -28.98±11.94 19.571±0.004 18.563±0.004 16.691±0.005 15.980±0.016 15.436±0.016
PLZJ 77 PLIZ28,BRB18 03 54 14.04 +23 17 52.28 12.01±14.59 -51.60±15.84 20.760 ± 0.010 19.728 ± 0.010 17.647 ± 0.010 16.789±0.031 16.131±0.030
PLZJ 21 PLIZ31 03 51 47.65 +24 39 59.18 17.84±9.41 -44.92±8.03 20.944±0.014 19.762±0.013 17.575±0.012 16.774±0.026 16.089±0.028
PLZJ 10 PLIZ35,BRB15 03 52 39.13 +24 46 29.61 15.84±8.88 -49.34±6.24 21.293±0.018 20.292±0.016 18.181±0.022 17.118±0.041 16.506±0.0416
PLZJ 4 BRB21 03 54 10.25 +23 41 40.67 29.74±13.17 -38.46±8.88 21.322 ± 0.010 20.215 ± 0.013 18.171 ± 0.010 17.141±0.045 16.377±0.039
PLZJ 61 BRB22 03 44 31.27 +25 35 14.97 25.82±7.89 -40.21±8.47 22.043 ± 0.030 20.782 ± 0.026 18.298 ± 0.020 17.393±0.059 16.657±0.04
PLZJ 32 BRB27 03 44 27.27 +25 44 41.99 25.03±11.52 -38.65±23.46 22.235 ± 0.040 20.962 ± 0.029 18.871 ± 0.030 17.793±0.094 16.950±0.070
PLZJ 37 BRB28 03 52 54.92 +24 37 18.85 18.13±11.53 -48.68±11.38 22.452 ± 0.05 21.216 ± 0.041 18.839 ± 0.030 17.742±0.071 16.921±0.058
PLZJ 23 03 51 53.38 +24 38 12.11 20.75±10.51 -50.05±9.96 23.541 ± 0.140 22.187 ± 0.112 19.960 ± 0.100 19.362±0.100 18.510±0.030
PLZJ 93 03 55 13.00 +24 36 15.8 13.11±14.36 -33.77±12.97 24.488 ± 0.370 22.592 ± 0.164 19.968 ± 0.080 19.955±0.100 19.420 ±0.100
PLZJ 323 03 43 55.27 +25 43 26.2 29.87±12.05 -39.37±11.70 - 21.597±0.054 19.613±0.076 - -
PLZJ 721 03 55 07.14 +24 57 22.34 19.18±22.23 -40.70±12.38 - 22.195±0.092 20.248±0.116 20.417±0.123 -
PLZJ 235 03 52 32.57 +24 44 36.64 20.92±12.16 -45.84±11.75 - 22.339±0.115 20.039±0.112 20.245±0.127 -
PLZJ 112 03 53 19.37 +24 53 31.85 8.56±14.08 -34.59±19.99 - 22.532±0.116 20.281±0.143 - -
PLZJ 100 03 47 19.19 +25 20 53.3 20.23±14.27 -37.28±23.82 - 23.563±0.373 20.254±0.114 - -
000–000
000,1
	Introduction
	Observations, Data Reduction and Survey Completeness
	The J band imaging and its reduction
	The far-red optical imaging and a new reduction
	The completeness of datasets
	Analysis of the data
	Photometric selection of candidate cluster members
	Refining the sample using astrometric measurements
	Results
	Mass spectrum
	Conclusions
	Acknowledgements
