# A photometric and astrometric investigation of the brown dwarfs in Blanco 1

S. L. Casewell, D. E. A. Baker, R. F. Jameson, S. T. Hodgkin, P. D. Dobbie, E. Moraux

► To cite this version:

S. L. Casewell, D. E. A. Baker, R. F. Jameson, S. T. Hodgkin, P. D. Dobbie, et al. A photometric and astrometric investigation of the brown dwarfs in Blanco 1. Monthly Notices of the Royal Astronomical Society, 2012, 425, pp.3112-3118. 10.1111/j.1365-2966.2012.21535.x. insu-03612402

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# A photometric and astrometric investigation of the brown dwarfs in Blanco 1

S. L. Casewell,<fnref n="1" />*D, D. E. A. Baker,<fnref n="1" />,<fnref n="2" /> R. F. Jameson,<fnref n="1" /> S. T. Hodgkin,<fnref n="3" /> P. D. Dobbie<fnref n="4" /> and E. Moraux<fnref n="5" />

Accepted 2012 June 15. Received 2012 June 15; in original form 2012 April 4

## ABSTRACT

We present the results of a photometric and astrometric study of the low-mass stellar and substellar population of the young open cluster Blanco 1. We have exploited J-band data, obtained recently with the Wide Field Camera (WFCAM) on the United Kingdom Infrared Telescope (UKIRT), and 10-year-old I- and z-band optical imaging from CFH12k on the Canada-France-Hawaii Telescope (CFHT), to identify 44 candidate low-mass stellar and substellar members, in an area of 2 deg<fnref n="2" />, on the basis of their colours and proper motions. This sample includes five sources which are newly discovered. We also confirm the lowest mass candidate member of Blanco 1 unearthed so far (29*M*<sub>Jup</sub>). We determine the cluster mass function to have a slope of $\alpha = +0.93$, assuming it to have a power-law form. This is high, but nearly consistent with previous studies of the cluster (to within the errors), and also that of its much better studied Northern hemisphere analogue, the Pleiades.

**Key words:** brown dwarfs – stars: low-mass – open clusters and associations: individual: Blanco 1.

## 1 INTRODUCTION

Open clusters are often acclaimed as excellent laboratories with which to study star formation. This is due to the coeval nature of their members and estimates of their age being comparatively robust. Many open star clusters have been studied to date, yielding a large number of low-mass members (e.g. Casewell et al. 2007; Lodieu et al. 2007; Baker et al. 2010) which have been used to refine our knowledge about the low-mass end of star formation via mapping the initial mass function (IMF). The IMF, the number of objects per unit mass interval, is an observable outcome of star formation and can be used to critically examine theoretical models of this process. The IMF is commonly measured using an $\alpha$ parameter given by $dN/dM \propto M^{-\alpha}$, where N is the number of objects and M is mass. For most open star clusters (ages 100 Myr), $\alpha$ is roughly consistent across all samples and $\approx 0.6$ (Bouvier, Moraux & Stauffer 2005). This value is also consistent with field values such as those of Chabrier (2003), although recently it has been suggested that for very low mass field brown dwarfs, the IMF may have a different form. Indeed, Burningham et al. (2010) suggest that in this case $\alpha$ may even have a negative value.

In recent years there has been a particular emphasis on building a solid comprehension of the mechanisms by which very low mass brown dwarfs and free-floating planetary mass objects form (e.g. Bate 2011). Nevertheless, key questions remain to be answered, e.g. 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 fragmentation and collapse of molecular clouds, then there is a strict lower mass limit to their manufacture of 0.007–0.010 M<sub>☉</sub>. 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, hypothetical magnetically mediated rebounds in collapsing cloud cores might lead 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 $\sim 0.001$ M<sub>☉</sub> (Boss 2001).

However, while many surveys of open star clusters have been performed to search for substellar members, the majority of these are in the heavily populated Northern hemisphere clusters. The lack of southern coverage from surveys [e.g. Sloan Digital Sky Survey: York et al. 2000; United Kingdom Infrared Telescope (UKIRT) Infrared Deep Sky Survey: Warren et al. 2007] has impeded detailed studies of the substellar population of a plethora of potentially interesting southern open clusters.

<footnote n="1">¹Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH</footnote>

<footnote n="2">²Science & Technology Research Institute, University of Hertfordshire, College Lane, Hatfield AL10 9AB</footnote>

<footnote n="3">³CASU, Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA</footnote>

<footnote n="4">⁴University of Tasmania, Private Bag 50, Hobart, Tasmania 7001, Australia</footnote>

<footnote n="5">⁵UJF-Grenoble 1/CNRS-INSU, Institut de Planétologie et d'Astrophysique de Grenoble (IPAG) UMR 5274, Grenoble F-38041, France</footnote>

<footnote>*E-mail: slc25@le.ac.uk</footnote>
Blanco 1 is a 90 ± 25 Myr (Panagi & O'dell 1997) open cluster with an age similar to that of the 125-Myr Pleiades cluster (Stauffer, Schultz & Kirkpatrick 1998) at a distance of 207 ± 12 pc as determined from *Hipparcos* measurements (van Leeuwen 2009). Recent work on the cluster includes spectroscopy of F- and G-type stars (Ford, Jeffries & Smalley 2005) which show that the metallicity is [Fe/H] = +0.04, with subsolar abundances for [Ni/Fe], [Si/Fe], [Mg/Fe] and [Ca/Fe]. Cargile, James & Jeffries (2010) have determined a lithium age for the cluster of 132 ± 24 Myr, which is closer to the age of the Pleiades than that determined by Panagi & O'dell (1997). We have taken the age of the cluster to be 120 Myr which is close to both measured values, and is present in the Chabrier et al. (2000) DUSTY models.

Recently, Platais et al. (2011) surveyed 11 deg² of the cluster to provide a comprehensive proper-motion catalogue for all stellar objects down to M5V. Moraux et al. (2007) performed the first study of the cluster to search for brown dwarfs using CFH12k on the Canada-France-Hawaii Telescope in the optical z and I bands to image 2.3 deg² of the cluster centre. They discovered ≈300 cluster members; 30–40 were estimated to be brown dwarfs, some of which had additional K-band photometry and optical spectroscopy. Three of these objects were subsequently confirmed as members by Cargile et al. (2010).

We have used the *I*- and *z*-band images from Moraux et al. (2007) and have combined them with additional deep (*J* ≈ 22) *J*-band photometry obtained using Wide Field Camera (WFCAM) on UKIRT, allowing us not only to select fainter candidate cluster members, but also to measure the proper motion for some of the previously identified objects to prove if they are indeed associated with the cluster.

## 2 OBSERVATIONS AND DATA REDUCTION

### 2.1 CFH12k data

The initial Blanco 1 data were taken with the CFH12k optical mosaic camera during two separate runs as detailed in Moraux et al. (2007). The first of the two runs occurred between 1999 September 30 and October 2, with the second occurring between 2000 December 18 and 20. A total of seven fields were observed covering an area of 2.3 deg² in a (mostly) non-overlapping pattern. Each separate field covered an area of 28 × 42 arcmin². The area of the sky covered is shown in Fig. 1. For each filter, Mould *I* and *z* Prime (see http://www.cfht.hawaii.edu/Instruments/Filters/cfh12k.html, for filter profiles), a short observation of 10 s was accompanied by two longer 600-s exposures. These were then combined to produce an equivalent image containing 1200 s worth of exposure. The detection limits of the data were *I* ~ *z* ~ 24 (Moraux et al. 2007), well below the stellar/substellar boundary which for Blanco 1 is estimated to lie at *I* ≈ 19.15. The reduction of the initial data by Moraux et al. (2007) followed the same prescription as described in Moraux et al. (2003).

The raw CFH12k data frames were extracted from the Canadian Astrophysical Data Centre (CADC) archive and were re-reduced using the imaging pipeline (Irwin & Lewis 2001) following the procedures described in Casewell et al. (2007). Subsequently, the two 600-s images in each filter at each pointing were co-added prior to source extraction and catalogue generation. Sources were identified as having a minimum of five interconnected pixels sitting at a significance of 1.5σ above the background, with aperture photometry carried out using a radius of 3.5 pixels. In addition, a morphological classification flag was provided, with -1 indicating

a stellar-like profile, 0 noise and +1 non-stellar-like sources. For the field-filter-extension/chip combinations of F4-*I*-10, F4-*z*-10, F6-*I*-6, F6-*z*-6 and F8-*I*-8, the astrometry needed further correction to that supplied by the Cambridge Astronomical Survey Unit (CASU) pipeline which was accomplished by using 'AAA' rated stars in Two Micron All Sky Survey (2MASS).

To refine the photometric calibration used by Moraux et al. (2003) which was based on A0 stars, we calculated a zero-point for each chip in each filter using data from European Southern Observatory (ESO).

Blanco 1 formed part of a study of young open clusters by the Monitor project (e.g. Irwin et al. 2008, 2009, and references therein). The observations were obtained using the MPG/ESO 2.2-m telescope with Wide Field Imager (WFI) in service mode, with around 500 epochs measured between 2005 July and 2007 October for four pointings (see Fig. 2). The instrument provides a field of view of ~34 × 33 arcmin² (0.31 deg²), using a mosaic of eight 2k × 4k pixel CCDs, at a scale of ~0.238 arcsec pixel⁻¹. The filter used was

<figure><img src="image_2.png" /><figcaption>**Figure 1.** Outline of the sky coverage of Blanco 1 from the CFH12k tiles (black) and the WFCAM tiles (red). The LMC and VLMC lists of Moraux et al. (2007) are shown as the small and large blue dots, respectively. The black cross indicates the cluster centre, with circles of radius 0.5, 1.0 and 1.5 being shown by the dashed black lines.</figcaption></figure>

<figure><img src="image_3.png" /></figure>
the ESO WFI broad-band *I* filter (designated BB# I203_ESO879, also known as the *I*<sub>EIS</sub> filter) with a central wavelength of 826.9 nm and a sharp cut-off in the red shortwards of 950 nm.

For a full description of the data reduction steps, the reader is referred to Irwin et al. (2007). Briefly, we used the pipeline for the Isaac Newton Telescope (INT) wide-field survey (Irwin & Lewis 2001) for 2D instrumental signature removal (bias correction, flat-fielding, defringing) and astrometric and photometric calibrations. We then generated a master catalogue for each filter by stacking 20 of the frames taken in the best conditions (seeing, sky brightness and transparency) and running the source detection software on the stacked image. Astrometric calibration is tied into 2MASS and has residuals of better than 0.05 arcsec pointing<sup>-1</sup>. Photometric calibration of our data was carried out using regular observations of Landolt (1992) equatorial standard star fields, measured as part of the standard ESO nightly calibrations. Prior to applying the calibration, we converted the Landolt (Cousins) photometry into the *I*<sub>EIS</sub> system using colour equations from Mike Irwin (private communication; see equation 1):

$$I_{EIS} = I - 0.03(V - I) \quad (1)$$

(and also Irwin et al. 2008). By converting the Landolt standards into the EIS system, and working entirely in the natural system of the instrument, this stage of the calibration of the CCD photometry is independent of the differences in colour between the cluster members which are somewhat redder than the Landolt standards.

The *I* and *z* zero-points used to calibrate the CFHT data were then calculated by comparing the uncalibrated instrumental magnitudes against those from the Monitor project. Data from each of the 12 CFH12k CCDs were binned for each of the two separate runs, i.e. all the objects found on chip 6 (over the different fields) taken in the 1999 run were combined to provide one single photometric zero-point for that chip. As small regions of overlap exist between some of the CFH12k fields, a test of photometric accuracy was conducted for those objects with duplicate detections. This yielded rms values of ≈0.035 and ≈0.040 for the *z* and *I* filters, respectively.

To confirm the calibration we first attempted to use the APASS survey (http://www.aavso.org/apass); however, there was insufficient overlap in magnitudes (the survey objects are saturated in the CFHT image) for us to use these data. We then cross-correlated the objects from Moraux et al. (2007) with our data and used them to check the calibration. We obtain an offset of *I* = +0.066 ± 0.018 and *z* = +0.080 ± 0.012 between the original CFHT data and our reprocessed images, with the original data being fainter. This is marginally larger than the rms scatter, but is in general smaller than the errors on the measurements, and so we are satisfied that our calibration is accurate.

Following the photometric calibration, the separate *I* and *z* catalogues for each CFH12k CCD chip were merged. This was done by using a flux-limited sample of objects that had been morphologically classified as stellar. This subset was used as an input for pattern matching and linear transformation equation generation between the associated *x* and *y* pixel coordinates of the objects. Once a transformation had been established for the 'clean' sample, it was used to match the full sample together helping reduce the number of spurious detections between the two images.

## 2.2 WFCAM data

In addition to the optical data, near-infrared (near-IR) observations were also taken. Three WFCAM (Casali et al. 2007) *J*-band tiles were obtained in UKIRT service mode, two on the night of 2006

October 31 and one on the night of 2009 July 22. Each WFCAM paw print used exposures of 18 s and a nine-point jitter pattern with 2 × 2 microstepping to improve the spatial sampling, making 1 h of observations in total per tile: 600-s exposure per paw print, in seeing of ≈1 arcsec or better. The WFCAM data were processed as for the Pleiades survey of Casewell et al. (2007). The calibration and pipeline for the data reduction are described in Hodgkin et al. (2009). The photometric calibration is tied to 2MASS photometry resulting in accuracies of ~1.5 per cent. The total area covered by the WFCAM fields is 2.25 deg<sup>2</sup>, of which ≈2 deg<sup>2</sup> overlaps with the CFH12k data as shown in Fig. 1.

The CFH12k data were pattern matched to the WFCAM data on an individual field by field, chip by chip basis to minimize multiple detections. Each source also had to be classified as stellar in both the CFH12k *I*- and *z*-band images as well as the WFCAM *J*-band image. The resulting catalogue contained 9853 sources (8440 of which were unique).

We estimated the completeness of both the CFH12k and WFCAM images using the method described in Casewell et al. (2007). We inserted 200 fake stars (12 < *J* < 22, 15 < *I*, *z* < 30) generated by IRAF into each chip, 10 times to enable us to have sufficient objects on which to perform the statistics. The sky level, detector gain, seeing, exposure times and zero-point of the images were taken into account when creating the fake stars. The CASU routine IMCORE was used to extract the objects from each image and then the numbers of inserted and extracted objects were compared per magnitude bin. The data were found on average to be 90 per cent complete at 19.73 in the *J* band and 21.5 and 20.6 in the *I* and *z* bands, respectively, and 50 per cent complete at 20.8 in the *J* band and 22.2 and 21.5 in the *I* and *z* bands, respectively. In general, it was found that chip number 4 of WFCAM was about 0.2 mag less sensitive than the other three chips.

# 3 RESULTS

## 3.1 Photometric selection

For consistency with previous studies of objects within this effective temperature range we chose to use the DUSTY models of Chabrier et al. (2000). We selected all sources with *I* − *J* within 0.5 of each side of the model for 120 Myr at 207 pc. This selection allows for uncertainty in distance and the equal-mass binary sequence. We then applied additional selection criteria of *I* − *J* > 1.95 and *z* − *J* > 1.15 to extract the sequence from the bulk of the field stars (Fig. 3). We selected a total of 83 objects using this method.

To determine the accuracy of our selection criteria, we compared how many of the objects presented in Moraux et al. (2007) were recovered by our survey. Moraux et al. (2007) present 764 unique sources, titled low-mass candidates (LMCs) and very low mass candidates (VLMCs), 578 of which are located within our survey area. Many of these objects were discovered using short exposures of 10 s and so are saturated in our data (1200 s in the *I* and *z* bands). To allow for this, and to better exploit our deeper data, we then applied a bright limit of *I* = 18.5 in our selection criteria. We recovered 522 objects in our survey area before the selection criteria were applied. The missing ≈50 objects have not been recovered due to falling between chip gaps, or not being detected in our *J*-band data as they are not red enough to be cluster members. Of the VLMCs, 81 are covered by our survey, and we recover all of these objects in our data.

However, after the selection cuts were made, only 27 LMC objects remained and 38 VLMC objects although there is some overlap
<figure><img src="image_4.png" /><figcaption>**Figure 3.** Colour–magnitude selections in *I*, *I* − *z*. The black points are the stellar CFH12k–WFCAM sources. Red ‘+’ marks the selected objects, while large blue filled circles indicate objects that remained after the proper-motion selection. The objects identified as candidate members from Moraux et al. (2007) are marked by boxes. Representative error bars and masses are also shown, as is the DUSTY model isochrone for 120 Myr (Chabrier et al. 2000).</figcaption></figure>

between the lists (Table A1; Appendix A). The remainder were lost
as they were brighter than I = 18.5, or fell outside the strip de-
fined by the model, i.e. despite appearing to belong to the cluster
sequence in I − z, they do not appear to belong to the sequence in
I − J or z − J (generally being too blue), and so are probably not
members of the cluster.

**3.2 Proper motions and membership probabilities**

We measured the proper motions of the 83 selected objects using
the *z* and *J* bands, which gave an epoch difference of 10 years for
the majority of objects, although a handful had a shorter epoch
difference of only 8 years. The *z* band was chosen over the *I* band
to minimize any effects of differential chromatic refraction as all
images were taken at high air mass, due to Blanco 1 being near
the observing limits of both UKIRT and CFHT. We used a pixel–
pixel transformation routine that uses a set of stationary reference
sources in each image as described by Casewell et al. (2007). The
reference objects were selected to have magnitudes 16 ≤ *mag* < 20
in *z*, similar to that of the candidates, but not so faint as to have poor
pixel centroiding. It was also required that they have an ellipticity
of less than 0.2 in the *z*-band image and be located within 10 arcmin
of the candidate to minimize radial distortion effects. In regions of
overlap it was also ensured that the candidate be on the same chip
as the reference stars in each image.

Centroiding errors were estimated using fake stars as for the
completeness calculations; only this time the difference in pixel po-
sitions between the inserted and recovered stars was measured. This
difference was measured in magnitude bins, as it was anticipated
that fainter objects would have larger centroiding errors. These er-
rors were 0.01 pixels in *J* for *J* < 17.0 and 0.07 for *J* > 17.0, where
the detector pixel size is 0.4 arcsec. For the *z* band, the errors were
0.03 for *z* < 23.0, where the detector pixel size is 0.2 arcsec. These

pixel measurements were added quadratically to the rms error on
the pixel–pixel transforms to generate the proper-motion errors.

Once we had measured the proper motions, the data were binned
in 10 mas yr⁻¹ bins in both RA and dec., and a 2D Gaussian was
fitted to the data in proper-motion space. The σ derived was then
used to reject objects outside the 2σ boundary to remove outliers,
and the fit was then recalculated. This gave a Gaussian width of σ
~ 9.0 mas yr⁻¹ to be used for candidate selection.

It is obvious from Fig. 4 that the average proper motion of our
selected objects (μα cosδ = 8.93 mas yr⁻¹, μδ = 6.70 mas yr⁻¹) is
significantly different from the literature value of the cluster proper
motion (μα cosδ = 20.11 mas yr⁻¹, μδ = 2.43 mas yr⁻¹; van Leeuwen
2009). Platais et al. (2011) determined that the mean motion of field
stars is not at 0,0 as is generally used for relative proper motions,
but at μα cosδ = 8.0 mas yr⁻¹, μδ = -6.0 mas yr⁻¹. To determine if
this offset was applicable to our data we modified our photometric
selection criteria to obtain everything 0.5 mag bluer in I - J than
the DUSTY model (Chabrier et al. 2000). We then measured proper
motions for these 250 objects, and fitted a 2D Gaussian to their
proper motions as before. We determined that the centre of this
distribution is at μα cosδ = 1.96 mas yr⁻¹, μδ = -3.64 mas yr⁻¹, with
a width of 12 mas yr⁻¹. This motion is smaller than that measured
by Platais et al. (2011), but they measured many more stars and
to a better accuracy than this work which has concentrated on the
fainter members of the cluster. This mean motion explains the offset
between our candidate distribution and the cluster motion. Taking
into account the offset makes our mean proper motion μα cosδ =
16.93 mas yr⁻¹, μδ = 0.70 mas yr⁻¹ which is much closer to the
reported value for the cluster. We used the literature value of the
cluster proper motion, minus the Platais et al. (2011) estimation of
the field star motion, μα cosδ = 12.11 mas yr⁻¹, μδ = 8.43 mas yr⁻¹,
as the cluster centre for selection purposes. It should also be noted
that the field and cluster stars are not that far apart in terms of the
relative errors which Platais et al. (2011) discuss in more depth in
<figure><img src="image_5.png" /><figcaption>**Figure 4.** Proper-motion diagram for the Blanco 1 cluster. Objects marked by a '+' are all objects we measured proper motions for. Objects marked by a filled circle are selected objects; they fall within a 27 mas yr⁻¹ circle centred on the cluster proper motion once it has been adjusted for the field star motions (μα cosδ = 12.11 mas yr⁻¹, μδ = 8.43 mas yr⁻¹). The centre of the dotted circle shows the unshifted cluster proper motion (μα cosδ = 20.11 mas yr⁻¹, μδ = 2.43 mas yr⁻¹).</figcaption></figure>

their work on Blanco 1. Despite the small difference between the proper motion of the cluster and field stars, the narrower dispersion in proper motion for the cluster objects, and the colour selections in *I* − *J* and *z* − *J* mean we can be confident that we are selecting true cluster members.

Of the 83 candidates for which we obtained astrometry, 44 had proper motions within (or with errors within) 3σ of the cluster value adjusted to take into account the field star relative motion (μα cosδ = 12.11 ± 0.38, μδ = 8.43 ± 0.25; van Leeuwen 2009; Platais et al. 2011). Of these 44 members, 33 are present in the VLMC list and 24 are present on the LMC list, with 18 objects common to both lists of candidate sets. This leaves five new low-mass candidate members to the cluster (Table 1). The previously identified objects that were rejected are LMC 694/VLMC 66, VLMC 64, VLMC 68, VLMC 69, VLMC 71 and VLMC 74. All VLMC objects with spectra remain in our candidates apart from objects 29, 38, 48 and 49 which are not in our survey area, and 16 and 22 which did not meet our photometric selection criteria (they are too bright). It should be noted that the three objects identified by Moraux et al. (2007) as non-members based upon their spectroscopy, VLMC 28, VLMC 37 and VLMC 44 were not selected (VLMC 28 also does not have a spectrum indicative of it being a cluster member; Cargile et al. 2010). The success of this method in recovering the previously identified spectroscopic members, despite large errors on the small cluster proper motion, leads us to believe this is a robust method for determining cluster members.

Examining the proper-motion vector point diagram (Fig. 4) it is clear that towards the location of the cluster there is an overdensity of objects when compared with regions at a similar distance from 0,0 but on the opposing side of the field star distribution. Unfortunately, the low number of sources coupled with a cluster proper motion comparable to the average proper-motion error means that the two-

Gaussian approach to calculating membership probability as used by Baker et al. (2010) is not applicable. Instead, the simpler annulus method was used as in Casewell et al. (2007). All non-selected objects with measured proper motions (barring those where there has obviously been a problem with the fit – seven cases) were used per *I* magnitude bin to assess the contamination and membership probabilities. We were unable to use an annulus centred on 0,0 as there are very few objects there not deemed to be members of Blanco 1 due to the low proper motion of the cluster. This may mean that some of the membership probabilities have been underestimated (Table 2).

## 4 MASS SPECTRUM

We were not able to generate mass functions for our new candidate members alone, due to there being too few objects to be statistically significant. However, we have used the Moraux et al. (2007) members combined with the *I* − *z* colour and the NextGen model (Baraffe et al. 1998) for objects brighter than *I* = 20.0, and the DUSTY model (Chabrier et al. 2000) for objects fainter than this. Both models are for 120 Myr. We present three mass spectra, one the original data from Moraux et al. (2007) and the second with the non-members, as determined from this work, excluded. The third data set is all the members, including our five new objects (Fig. 5). These new objects have masses between 35 and 46*M*<sub>Jup</sub>, whereas the whole mass range of candidates is between 29 and 80*M*<sub>Jup</sub>. The third data set (filled circles) is fitted by the straight line in Fig. 5, with an α value of 0.93 ± 0.11. However, to obtain this fit, we have omitted the point at log *M* = −1.85: the faintest objects and lowest mass bin where we know the incompleteness is largest. We do not have a good enough photometry to accurately separate the single and binary star sequences, and thus suspected binaries have been assigned a mass equivalent to that of a single object at their recorded magnitude. The point at log *M* = −0.85 is discrepantly high and appears to be affected by binaries and possibly higher multiples at *I* ≈ 20 (Fig. 3).

The α = 0.93 ± 0.11 indicates the slope is higher, but is consistent (to within the errors) with the values given by Moraux et al. (2007), which are 0.67 ± 0.14 and 0.71 ± 0.13 for 100- and 150-Myr models, respectively, especially considering we have small number statistics in some mass bins.

While we have discovered five new low-mass cluster members, our *J*-band data did not allow us to probe deeper into Blanco 1 than the original *I*<sub>*z*</sub> survey by Moraux et al. (2007). This situation will be improved as the VISTA VIKING survey which aims to cover the whole of the Blanco 1 cluster. It will also provide deeper multiband photometry, as well as a far greater baseline between observations to be used for proper-motion analysis. Our work has shown that Blanco 1 does contain brown dwarfs as low in mass as ~30*M*<sub>Jup</sub>, making it similar to the Pleiades.

Follow-up spectroscopic data will allow us to place constraints on the binary fraction of the cluster, as well as confirm membership for the objects without spectra. Decreasing the errors on the proper motions will allow us to determine members with much more confidence than we are currently able to due to the low space motion of the cluster. Once a full census of the cluster has been performed, Blanco 1 can be properly compared to clusters such as the Pleiades. One can then test for the environmental tolerance of the IMF, dynamical evolution and mass segregation effects as well as providing further observational constraints to compare with the results of brown dwarf formation simulations.
<table><caption>**Table 1.** Name, Moraux et al. (2007) name, coordinates, proper motion, and *I*, *z* and *J* magnitudes for our members to the cluster. Previously discovered members also have their other known names listed. An asterisk (*) indicates that membership has been confirmed from the Moraux et al. (2007) spectroscopy.</caption><thead><tr><th>Name</th><th>Alternate name</th><th>RA<br>(J2000.0)</th><th>Dec.</th><th>μ<sub>α</sub>cosδ<br>(mas yr<sup>-1</sup>)</th><th>μ<sub>δ</sub></th><th>I</th><th>z</th><th>J</th></tr></thead><tbody><tr><td>bl2399-4716</td><td></td><td>00 00 6.83</td><td>-30 13 33.72</td><td>+2.27 ± 12.96</td><td>+20.66 ± 11.56</td><td>19.045 ± 0.010</td><td>18.201 ± 0.033</td><td>16.851 ± 0.007</td></tr><tr><td>bl28626-47167</td><td>LMC 571/VLMC 19</td><td>00 07 50.63</td><td>-30 5 9.97</td><td>+7.45 ± 6.86</td><td>+5.90 ± 5.45</td><td>18.515 ± 0.013</td><td>17.688 ± 0.037</td><td>16.391 ± 0.005</td></tr><tr><td>bl2868-4699</td><td>VLMC 89</td><td>00 00 5.86</td><td>-30 20 18.39</td><td>+13.99 ± 8.10</td><td>+14.07 ± 7.33</td><td>21.971 ± 0.050</td><td>20.755 ± 0.073</td><td>18.625 ± 0.027</td></tr><tr><td>bl28691-43204</td><td>VLMC 54</td><td>00 07 41.46</td><td>-29 56 20.39</td><td>+31.93 ± 5.57</td><td>-7.13 ± 3.81</td><td>19.971 ± 0.017</td><td>18.782 ± 0.041</td><td>16.993 ± 0.008</td></tr><tr><td>bl32426-33057</td><td>LMC 629/VLMC 46</td><td>00 05 57.04</td><td>-29 43 48.33</td><td>-1.93 ± 4.90</td><td>-26.66 ± 13.69</td><td>19.328 ± 0.015</td><td>18.431 ± 0.039</td><td>17.128 ± 0.009</td></tr><tr><td>bl32697-33301</td><td>LMC 729/VLMC 72</td><td>00 06 8.96</td><td>-29 44 25.22</td><td>+2.45 ± 22.41</td><td>+16.06 ± 6.60</td><td>21.280 ± 0.029</td><td>20.035 ± 0.054</td><td>18.219 ± 0.019</td></tr><tr><td>bl33600-62347</td><td>LMC 626</td><td>00 06 41.73</td><td>-29 42 52.52</td><td>-13.07 ± 13.81</td><td>-7.03 ± 9.07</td><td>19.310 ± 0.015</td><td>18.466 ± 0.039</td><td>17.248 ± 0.009</td></tr><tr><td>bl33721-62566</td><td>LMC 705/VLMC 63</td><td>00 06 49.36</td><td>-29 40 41.18</td><td>+10.90 ± 8.92</td><td>+12.40 ± 10.72</td><td>20.698 ± 0.022</td><td>19.521 ± 0.046</td><td>17.798 ± 0.013</td></tr><tr><td>bl34536-62053</td><td>LMC 118</td><td>00 06 32.59</td><td>-29 46 05.57</td><td>+11.29 ± 9.71</td><td>+12.87 ± 8.17</td><td>22.782 ± 0.086</td><td>21.357 ± 0.097</td><td>19.111 ± 0.037</td></tr><tr><td>bl37640-49624</td><td>LMC 592/VLMC 30</td><td>00 08 27.38</td><td>-29 43 54.26</td><td>+10.94 ± 18.20</td><td>+17.94 ± 18.40</td><td>18.708 ± 0.014</td><td>17.965 ± 0.037</td><td>16.737 ± 0.006</td></tr><tr><td>bl3800-15178</td><td>LMC 580/VLMC 25*</td><td>00 00 42.74</td><td>-30 17 43.43</td><td>+6.40 ± 13.62</td><td>+10.82 ± 6.75</td><td>18.612 ± 0.010</td><td>17.779 ± 0.032</td><td>16.497 ± 0.006</td></tr><tr><td>bl3819-15319</td><td>LMC 621</td><td>00 00 36.44</td><td>-30 19 15.95</td><td>+22.62 ± 13.28</td><td>+6.14 ± 5.19</td><td>19.184 ± 0.011</td><td>18.288 ± 0.033</td><td>17.002 ± 0.008</td></tr><tr><td>bl43328-55651</td><td>LMC 57</td><td>00 07 22.76</td><td>-30 01 57.32</td><td>-4.16 ± 14.08</td><td>+1.65 ± 8.34</td><td>20.139 ± 0.018</td><td>18.952 ± 0.043</td><td>17.623 ± 0.012</td></tr><tr><td>bl44742-27543</td><td>LMC 82</td><td>00 04 42.00</td><td>-30 04 33.52</td><td>+11.93 ± 4.88</td><td>-3.91 ± 1.70</td><td>21.704 ± 0.041</td><td>20.497 ± 0.058</td><td>18.425 ± 0.034</td></tr><tr><td>bl46456-37665</td><td>LMC 70</td><td>00 05 45.83</td><td>-30 03 46.14</td><td>+24.71 ± 12.74</td><td>+7.79 ± 11.09</td><td>21.221 ± 0.028</td><td>20.031 ± 0.049</td><td>18.142 ± 0.019</td></tr><tr><td>bl51473-23799</td><td>LMC 582/VLMC 26</td><td>00 04 54.98</td><td>-29 46 32.88</td><td>+7.79 ± 10.99</td><td>+11.10 ± 16.02</td><td>18.663 ± 0.010</td><td>17.935 ± 0.031</td><td>16.653 ± 0.006</td></tr><tr><td>bl51714-24041</td><td>LMC 719/VLMC 67</td><td>00 05 05.19</td><td>-29 49 55.81</td><td>+13.41 ± 14.08</td><td>+17.17 ± 6.20</td><td>20.943 ± 0.025</td><td>19.925 ± 0.042</td><td>18.035 ± 0.016</td></tr><tr><td>bl54505-36470</td><td></td><td>00 03 24.26</td><td>-30 00 51.66</td><td>+0.82 ± 16.17</td><td>+18.36 ± 17.75</td><td>18.959 ± 0.010</td><td>18.094 ± 0.031</td><td>16.785 ± 0.007</td></tr><tr><td>bl54514-36502</td><td>LMC 608/VLMC 41*</td><td>00 03 23.62</td><td>-29 55 17.56</td><td>+0.00 ± 0.63</td><td>-0.00 ± 0.63</td><td>19.058 ± 0.010</td><td>18.033 ± 0.031</td><td>16.522 ± 0.006</td></tr><tr><td>bl54613-36058</td><td>LMC 663/VLMC 60</td><td>00 03 40.17</td><td>-30 03 40.85</td><td>-0.57 ± 11.55</td><td>-10.33 ± 4.90</td><td>20.478 ± 0.017</td><td>19.384 ± 0.038</td><td>17.654 ± 0.012</td></tr><tr><td>bl54805-35986</td><td>LMC 74</td><td>00 03 43.78</td><td>-30 04 01.66</td><td>+26.29 ± 3.66</td><td>+32.89 ± 3.99</td><td>21.404 ± 0.030</td><td>20.167 ± 0.047</td><td>18.202 ± 0.019</td></tr><tr><td>bl55543-35398</td><td>LMC 645/VLMC 51*</td><td>00 04 07.62</td><td>-29 59 18.78</td><td>+3.57 ± 6.12</td><td>+3.91 ± 9.53</td><td>19.806 ± 0.014</td><td>18.782 ± 0.036</td><td>17.210 ± 0.009</td></tr><tr><td>bl56159-35803</td><td></td><td>00 03 50.87</td><td>-30 01 58.01</td><td>+17.68 ± 3.98</td><td>+12.55 ± 5.98</td><td>22.329 ± 0.070</td><td>21.084 ± 0.087</td><td>18.855 ± 0.032</td></tr><tr><td>bl56339-2865</td><td>LMC 624/VLMC 45*</td><td>00 01 35.71</td><td>-30 03 10.13</td><td>+18.36 ± 6.54</td><td>+1.41 ± 6.49</td><td>19.268 ± 0.012</td><td>18.326 ± 0.033</td><td>16.915 ± 0.008</td></tr><tr><td>bl57823-3435</td><td>LMC 114</td><td>00 02 5.68</td><td>-30 01 23.65</td><td>+10.59 ± 8.078</td><td>+8.75 ± 5.89</td><td>22.655 ± 0.110</td><td>21.382 ± 0.105</td><td>19.077 ± 0.042</td></tr><tr><td>bl57973-12146</td><td>LMC 647/VLMC 55</td><td>00 02 15.12</td><td>-30 09 52.64</td><td>-1.94 ± 2.57</td><td>+10.56 ± 11.70</td><td>19.835 ± 0.016</td><td>18.843 ± 0.036</td><td>17.235 ± 0.009</td></tr><tr><td>bl58756-8496</td><td>LMC 65</td><td>00 00 7.31</td><td>-29 54 26.81</td><td>-2.81 ± 9.46</td><td>+15.26 ± 14.08</td><td>20.575 ± 0.021</td><td>19.495 ± 0.043</td><td>17.751 ± 0.015</td></tr><tr><td>bl6053-16818</td><td>LMC 53</td><td>00 02 16.10</td><td>-30 18 41.03</td><td>+13.82 ± 14.46</td><td>+8.77 ± 9.38</td><td>19.810 ± 0.013</td><td>18.747 ± 0.034</td><td>17.031 ± 0.008</td></tr><tr><td>bl60868-7134</td><td>LMC 32</td><td>00 01 19.28</td><td>-29 54 06.58</td><td>+57.44 ± 25.62</td><td>-4.17 ± 14.72</td><td>18.779 ± 0.010</td><td>17.991 ± 0.031</td><td>16.762 ± 0.007</td></tr><tr><td>bl63549-4173</td><td>LMC 639/VLMC 50*</td><td>00 00 9.86</td><td>-30 01 59.39</td><td>-9.78 ± 19.95</td><td>+1.10 ± 12.38</td><td>19.612 ± 0.013</td><td>18.651 ± 0.034</td><td>17.229 ± 0.011</td></tr><tr><td>bl65492-14209</td><td>LMC 83</td><td>00 00 34.79</td><td>-30 02 51.04</td><td>+21.52 ± 0.63</td><td>+14.18 ± 0.63</td><td>21.800 ± 0.044</td><td>20.494 ± 0.059</td><td>18.450 ± 0.026</td></tr><tr><td>bl71077-57153</td><td>LMC 609/VLMC 36*</td><td>00 07 08.79</td><td>-30 06 42.35</td><td>+7.29 ± 6.58</td><td>+2.57 ± 6.10</td><td>18.958 ± 0.014</td><td>18.072 ± 0.038</td><td>16.713 ± 0.006</td></tr><tr><td>bl73194-41854</td><td>LMC 632</td><td>00 08 13.55</td><td>-30 16 50.17</td><td>-12.48 ± 11.52</td><td>-10.16 ± 7.71</td><td>19.384 ± 0.015</td><td>18.208 ± 0.038</td><td>16.964 ± 0.008</td></tr><tr><td>bl7507-10415</td><td>LMC 585</td><td>00 00 12.10</td><td>-30 35 56.07</td><td>+6.15 ± 8.99</td><td>+13.13 ± 13.04</td><td>18.676 ± 0.010</td><td>17.833 ± 0.031</td><td>16.599 ± 0.006</td></tr><tr><td>bl76145-40818</td><td>LMC 595/VLMC 33</td><td>00 06 10.79</td><td>-30 21 37.02</td><td>+3.23 ± 11.65</td><td>+5.40 ± 8.18</td><td>18.817 ± 0.014</td><td>17.961 ± 0.037</td><td>16.691 ± 0.007</td></tr><tr><td>bl76187-40858</td><td></td><td>00 06 08.53</td><td>-30 25 42.12</td><td>-2.96 ± 13.80</td><td>-9.37 ± 13.61</td><td>19.767 ± 0.017</td><td>18.851 ± 0.040</td><td>17.483 ± 0.011</td></tr><tr><td>bl76366-40509</td><td>LMC 75</td><td>00 06 29.32</td><td>-30 20 33.54</td><td>+4.38 ± 14.23</td><td>-3.55 ± 9.32</td><td>21.321 ± 0.035</td><td>20.260 ± 0.054</td><td>18.454 ± 0.023</td></tr><tr><td>bl78357-42493</td><td>LMC 631</td><td>00 07 33.84</td><td>-30 24 35.13</td><td>+9.32 ± 9.17</td><td>+7.78 ± 6.73</td><td>19.327 ± 0.015</td><td>18.376 ± 0.039</td><td>16.937 ± 0.008</td></tr><tr><td>bl84053-20433</td><td></td><td>00 03 17.87</td><td>-30 11 40.93</td><td>+6.94 ± 5.57</td><td>+18.65 ± 3.55</td><td>21.663 ± 0.037</td><td>20.350 ± 0.052</td><td>18.382 ± 0.023</td></tr><tr><td>bl85709-22911</td><td>LMC 604/VLMC 43*</td><td>00 04 32.88</td><td>-30 18 41.90</td><td>+9.41 ± 11.43</td><td>+14.39 ± 11.46</td><td>19.127 ± 0.010</td><td>18.238 ± 0.031</td><td>16.894 ± 0.008</td></tr><tr><td>bl862-2266</td><td>LMC 600/VLMC 34*</td><td>00 01 48.76</td><td>-30 38 06.81</td><td>-3.19 ± 7.45</td><td>-2.53 ± 5.11</td><td>18.951 ± 0.011</td><td>17.960 ± 0.032</td><td>16.624 ± 0.007</td></tr><tr><td>bl87593-22425</td><td>LMC 93</td><td>00 04 57.74</td><td>-30 14 02.01</td><td>+46.11 ± 18.36</td><td>-0.75 ± 18.34</td><td>22.186 ± 0.073</td><td>20.903 ± 0.058</td><td>18.822 ± 0.034</td></tr><tr><td>bl89514-11054</td><td>LMC 619</td><td>00 02 50.60</td><td>-30 28 53.89</td><td>-11.83 ± 17.94</td><td>+9.38 ± 11.27</td><td>19.241 ± 0.011</td><td>18.429 ± 0.033</td><td>17.211 ± 0.010</td></tr><tr><td>bl90054-11175</td><td>LMC 85</td><td>00 03 06.63</td><td>-30 29 53.90</td><td>+4.08 ± 6.31</td><td>+14.00 ± 5.78</td><td>21.999 ± 0.049</td><td>20.763 ± 0.067</td><td>18.624 ± 0.029</td></tr></tbody></table>

<table><caption>**Table 2.** Magnitude bins and the associated membership probability and completeness at *I* magnitude.</caption><thead><tr><th>I</th><th>Probability<br>(per cent)</th><th>Completeness<br>(per cent)</th></tr></thead><tbody><tr><td>&lt;19.5</td><td>91</td><td>100</td></tr><tr><td>19.5-20.5</td><td>75</td><td>97</td></tr><tr><td>20.5-21.5</td><td>50</td><td>93</td></tr><tr><td>21.5-22.5</td><td>50</td><td>62</td></tr><tr><td>&gt;22.5</td><td>40</td><td>16</td></tr></tbody></table>

## 5 SUMMARY

We have used near-IR and optical photometry with proper motions derived from CFHT *z*-band and WFCAM *J*-band images to identify

44 candidate cluster members with masses between 29 and $80M_{\text{Jup}}$. Five of these are previously unidentified candidate members and 40 have been identified by Moraux et al. 2007, eight of which have been confirmed as cluster brown dwarfs from spectra. We derive $\alpha = 0.93 \pm 0.11$ from the mass spectrum, which is consistent with the literature for this cluster.

## ACKNOWLEDGMENTS

DEAB acknowledges support from STFC, and SLC is supported by a University of Leicester post-doctoral research position. This paper made use of data from the United Kingdom Infrared Telescope which is operated by the Joint Astronomy Centre on behalf of the Science and Technology Facilities Council of the UK. We would also like to thank France Allard for the use of the Phoenix
<figure><img src="image_6.png" /><figcaption>**Figure 5.** The mass spectrum for Blanco 1. The circles are the original members detailed by Moraux et al. (2007). It should be noted that the errors are Poissonian, and in the case of the lowest mass bins, actually fall within the plotted point. The diamonds are the Moraux et al. (2007) objects, with non-members as determined from this work removed. The filled circles are a complete mass function for all objects we determined to be members. The fit to the data of $\alpha = 0.93 \pm 0.11$ is the solid line.</figcaption></figure>

web simulator at http://phoenix.ens-lyon.fr/simulator/index faces, which has been used in this research. This research has also made use of NASA's Astrophysics Data System Bibliographic Services.

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APPENDIX A: SUPPLEMENTARY TABLE

Table A1 displays a list of LMCs and VLMCs identified by Moraux et al. (2007), which were present in our survey area, but were rejected by our photometric selection.

<table><caption>**Table A1.** LMC and VLMC candidate members in our data, but determined to be non-members of the cluster due to photometry that was incompatible with our selection criteria.</caption><thead><tr><th>Name</th><th>Name</th></tr></thead><tbody><tr><td>LMC 573</td><td>LMC 678</td></tr><tr><td>LMC 574</td><td>LMC 679</td></tr><tr><td>LMC 576</td><td>LMC 682</td></tr><tr><td>LMC 578</td><td>LMC 683</td></tr><tr><td>LMC 581/VLMC 28</td><td>LMC 684</td></tr><tr><td>LMC 586</td><td>LMC 685</td></tr><tr><td>LMC 587</td><td>LMC 686</td></tr><tr><td>LMC 588</td><td>LMC 688</td></tr><tr><td>LMC 590</td><td>LMC 689</td></tr><tr><td>LMC 593</td><td>LMC 690</td></tr><tr><td>LMC 594</td><td>LMC 692</td></tr><tr><td>LMC 596</td><td>LMC 695</td></tr><tr><td>LMC 597</td><td>LMC 699</td></tr><tr><td>LMC 598</td><td>LMC 700</td></tr><tr><td>LMC 603/VLMC 39</td><td>LMC 702</td></tr><tr><td>LMC 610</td><td>LMC 704</td></tr><tr><td>LMC 611</td><td>LMC 707</td></tr><tr><td>LMC 614</td><td>LMC 708</td></tr><tr><td>LMC 616</td><td>LMC 709</td></tr><tr><td>LMC 620</td><td>LMC 710</td></tr><tr><td>LMC 625</td><td>LMC 711</td></tr><tr><td>LMC 627</td><td>LMC 712</td></tr><tr><td>LMC 630</td><td>LMC 713</td></tr><tr><td>LMC 635</td><td>LMC 715</td></tr><tr><td>LMC 637</td><td>LMC 716</td></tr><tr><td>LMC 638</td><td>LMC 717</td></tr><tr><td>LMC 640</td><td>LMC 718</td></tr><tr><td>LMC 642</td><td>LMC 720</td></tr><tr><td>LMC 643</td><td>LMC 721</td></tr><tr><td>LMC 644</td><td>LMC 722</td></tr><tr><td>LMC 646</td><td>LMC 724</td></tr><tr><td>LMC 652</td><td>LMC 725</td></tr><tr><td>LMC 655/VLMC 58</td><td>LMC 726</td></tr><tr><td>LMC 656</td><td>LMC 727</td></tr><tr><td>LMC 657</td><td>LMC 728</td></tr><tr><td>LMC 658</td><td>LMC 730</td></tr><tr><td>LMC 659</td><td>LMC 732</td></tr><tr><td>LMC 660</td><td>LMC 735</td></tr><tr><td>LMC 661</td><td>LMC 736</td></tr><tr><td>LMC 662</td><td>LMC 738</td></tr><tr><td>LMC 666</td><td>LMC 739</td></tr><tr><td>LMC 668</td><td>LMC 741</td></tr><tr><td>LMC 669</td><td>LMC 747</td></tr><tr><td>LMC 671</td><td>LMC 749</td></tr><tr><td>LMC 673</td><td>LMC 752</td></tr><tr><td>LMC 674</td><td>LMC 755</td></tr><tr><td>LMC 675</td><td>LMC 762</td></tr><tr><td>LMC 677</td><td></td></tr></tbody></table>

This paper has been typeset from a TeX/LA TeX file prepared by the author.