A&A 470, 281-294 (2007)
DOI: 10.1051/0004-6361:20066500

A WFI survey in the Chamaeleon II dark cloud[*],[*],[*]

L. Spezzi1,2,3 - J. M. Alcalá2 - A. Frasca1 - E. Covino2 - D. Gandolfi1,3

1 - INAF - Osservatorio Astrofisico di Catania, via S. Sofia 78, 95123 Catania, Italy
2 - INAF - Osservatorio Astronomico di Capodimonte, Salita Moiariello 16, 80131 Napoli, Italy
3 - Dipartimento di Fisica e Astronomia, Universitá di Catania, via S. Sofia 78, 95123 Catania, Italy

Received 3 October 2006 / Accepted 2 February 2007

Aims. We present the results of an optical multi-band survey for low-mass pre-main sequence (PMS) stars and young brown dwarfs (BDs) in the Chamaeleon II (Cha II) dark cloud. This survey constitutes the complementary optical data to the c2d Spitzer Legacy survey in Cha II.
Methods. Using the wide-field imager (WFI) at the ESO 2.2 m telescope, we surveyed a sky area of about 1.75 square degrees in Cha II. The region was observed in the $R_{\rm c}$, $I_{\rm c}$, and z broad-bands, in H$\alpha $ and in two medium-band filters centred at 856 and 914 nm. We selected PMS star and young BD candidates using colour-magnitude diagrams (CMDs) and theoretical isochrones reproduced ad-hoc for the WFI at the ESO2.2m telescope system. We reinforced the selection criteria by using the previously known PMS stars in Cha II to define the PMS locus on the CMDs and by investigating the infrared (IR) colours of the candidates. By exploiting the WFI intermediate-band photometry we also estimated the effective temperature and the level of H$\alpha $ emission of the candidates.
Results. Our survey, which is one of the largest and deepest optical surveys conducted so far in Cha II, recovered the majority of the PMS stars and 10 member candidates of the cloud from previous IR surveys. In addition, the survey revealed 10 new potential members. From our photometric characterisation, we estimate that some 50% of the 20 candidates will result in true Cha II members. Based on our temperature estimates, we conclude that several of these objects are expected to be sub-stellar and give a first estimate of the fraction of sub-stellar objects.

Key words: stars: low-mass, brown dwarfs - stars: formation - stars: pre-main sequence - ISM: clouds - ISM: individual objects: Chamaeleon II

1 Introduction

Recent investigations in star-forming regions (SFRs) point out that the fraction of young brown dwarfs (BDs) relative to low-mass and more massive pre-main sequence (PMS) stars may vary significantly among different SFRs (Kroupa 2002). In particular, the fraction is different in T and in OB associations (Muench et al. 2003; Hillenbrand & Carpenter 2000; Luhman et al. 2000; Preibisch et al. 2003; Briceño et al. 2002). Different environments may give rise to different initial conditions for star and planet formation and, hence, to differences in the observed spectrum of masses (Kroupa 2001,2002). This may have strong consequences on the initial mass function (IMF), in particular in the sub-stellar domain. It has also been proposed that BDs may form as members of small groups of objects, which may be ejected by dynamical interactions before they can grow to stellar masses (Reipurth & Clarck 2001); as a result, many low-mass and sub-stellar objects might have escaped detection in surveys that concentrate sharply in the cores of SFRs. Thus, one possible reason for the differences in the fraction of young sub-stellar objects in T and OB associations can in principle be ascribed to the photometric and spatial incompleteness of the imaging surveys. Spatially complete deep-imaging surveys are thus crucial in order to single out low-mass star and BD candidates to be investigated by follow-up spectroscopy. Only then, problems like mass segregation in SFRs and the low-mass end of the IMF can be addressed.

In this paper, optical wide-field imaging observations in the Chamaeleon II dark cloud (hereafter Cha II), complemented with ${\it JHK}$ photometry from 2MASS (Cutri et al. 2003), are used to search for low-mass PMS stars and young BDs. Our survey covers almost 2 square degrees in the Cha II cloud.

Because of its proximity to the Sun ( $d\approx178$ pc, Whittet et al. 1997), young age (0.1-10 Myr, Hughes & Hartigan 1992), and relatively high galactic latitude ( $b\approx -15$ deg, Schwartz 1991), which decreases the effects of contamination by background stars, the Cha II dark cloud is particularly well-suited to studies of low-mass PMS stars and young BDs. It is indeed characterised by the presence of objects with H$\alpha $ emission (Hartigan 1993; Hughes & Hartigan 1992), as well as of embedded Class-I and Class-II IR sources (Larson et al. 1998; Whittet et al. 1991; Prusti et al. 1992) and X-rays sources (Alcalá et al. 2000).

Investigations in the near-IR by the DENIS survey have revealed several candidate young BDs in Cha II (Vuong et al. 2001); however, a spectroscopic follow-up by Barrado y Navascués & Jayawardhana (2004) failed to confirm the young BD candidates, although it revealed the least massive classical T Tauri star in the cloud. Persi et al. (2003) performed ISOCAM observations and IR spectroscopy of several objects in the core of Cha II and found a number of sources with IR excess. Their most promising candidate, ISO-CHA II 13, has been recently confirmed as the first BD in the region known to be surrounded by a disk (Alcalá et al. 2006). An optical wide-field imaging survey by López-Martí et al. (2005) proposes two young BD candidates. However, that survey covers only about 10% of the cloud area.

The Cha II cloud has also been included in the Spitzer Legacy survey "From Molecular Cores to Planet Forming Disks'' or c2d (Evans et al. 2003) as a test case of a cloud with moderate star-formation activity. Results of the c2d survey in Cha II have been published recently by Young et al. (2005), Allers et al. (2006), and Porras et al. (2006) . The data presented in this paper constitute the optical ancillary data for the c2d survey in Cha II. As such, they are part of a multi-wavelength study of the Cha II cloud that will be presented in a forthcoming paper (Alcalá et al., in preparation).

The outline of the paper is as follows. In Sect. 2 the observations, data reduction and calibration procedures are described. In Sect. 3 we present specific tools for the analysis of the photometric data. Section 4 describes the criteria for the selection of PMS star and BD candidates in Cha II, using the tools developed in Sect. 3. The results of the survey and, in particular, the fraction of sub-stellar objects estimated in Cha II are discussed in Sect. 5. Our conclusions are presented in Sect. 6 and, finally, a few notes on some individual objects are presented in Appendix C (available only in electronic form).

\end{figure} Figure 1: IRAS 100 $\mu $m dust emission map of the Cha II dark cloud. The contours, from 10 to 40 MJy $\times $ sr-1 in steps of 2.5 MJy $\times $ sr-1, are also drawn (white lines). The big polygons define the area covered by the seven ESO2.2m+WFI pointings. The confirmed members of the cloud and the candidates selected in previous surveys are represented with open circles and squares, respectively; the black dots mark the objects recovered by our selection criteria (see Sect. 4).

2 Observations and data reduction

The observations were carried out in two observing runs (27-30 April 2001 and 20-23 March 2002) using the wide-field imaging (WFI) mosaic camera attached to the ESO 2.2m telescope at La Silla (Chile). The mosaic consists of eight 2k $\times $ 4k CCDs forming an 8k $\times $ 8k array with a pixel scale of 0.238 $^{\prime\prime}$/pix; hence, a single WFI pointing covers a sky area of about 30$^\prime$ $\times $ 30$^\prime$. The Cha II dark cloud has an extension of about 2 square degrees (Hughes & Hartigan 1992). Thus, seven adjacent WFI pointings allowed us to cover about 70% of the cloud area. The distribution on the sky of the WFI pointings is shown in Fig. 1. The overlap of about 2$^\prime$  between adjacent pointings allowed us to check the consistency in the photometry and in the astrometry.

The observations were performed in the RC, IC and z broad bands, in two H$\alpha $ filters, narrow (H$\alpha_7$, ${\rm\lambda_c}=658$ nm and ${\it FWHM}=7.4$ nm) and wide (H $\alpha _{12}$, ${\rm\lambda_c}=665$ nm and ${\it FWHM}=12.1$ nm), and in two intermediate-band filters centred at 865 and 914 nm. In order to cover the gaps between the WFI CCDs and to correct for moving objects and cosmic ray hits, a sequence of typically five ditherings was performed for every pointing in each filter. The journal of the observations is presented in Table 1. The RA (in hh:mm:ss) and Dec. (in dd:mm:ss) of each pointing are indicated between parenthesis; total exposure time ( $T_{\rm exp}$), seeing and air mass (X) values correspond to the final staked images (see Sect. 2.1).

2.1 Pre-reduction and astrometry

The pre-reduction of the raw images, passing through bias subtraction, flat-fielding and correction for fringing, was performed on a nightly basis, using the mscred package under IRAF[*], following the guidelines described in Alcalá et al. (2002). The resulting pre-reduced images are uniform to about 1%. To correct for large-scale illumination variations, a super-flat was used.

The astrometric calibration and relative flux scaling between ditherings were done using the ASTROMETRIX tool[*]. It performs a global astrometric solution that takes overlapping sources falling on adjacent CCDs in different ditherings into account. For each pointing, the astrometric solution was first computed for the R-band dithering set using the USNO-A2.0 catalogue (Monet et al. 1998) as a reference. A catalogue of sources was then extracted from the re-sampled R-band image and used as the reference catalogue for all the other bands. Within the global astrometry process, the astrometric solution was constrained for each CCD by both the positions from the R-band catalogue and those from overlapping sources in all the other CCDs. The co-addition of the different dithered images for a given filter and pointing was performed using the SWARP tool[*]. The resulting 8k $\times $ 8k stacked images were normalised by the exposure time. The absolute astrometric precision of our images is about 0.35 arcsec, corresponding to the RMS accuracy of the USNO-A2.0 catalogue, while the internal RMS, computed from overlapping sources in different exposures, is within 0.05 arcsec, indicating the good performance of ASTROMETRIX.

2.2 The photometric calibration

In order to transform the R and I instrumental magnitudes to the standard Cousins system, the Landolt standard fields SA 98, SA 101, and SA 107 (Landolt 1992) were observed nightly. The standard RC and IC magnitudes were determined using the transformation equations:

\begin{displaymath}R_C=r_0 + c_R \cdot (r_0-i_0)+ZP_R
\end{displaymath} (1)

\begin{displaymath}I_C=i_0 +c_I \cdot (r_0-i_0)+ZP_I
\end{displaymath} (2)

where r0 and i0 are the instrumental magnitudes corrected for atmospheric extinction and cR and ZPR and cI and ZPI the colour terms and zero points for the R and I bands, respectively. The standard magnitudes from Stetson (2000) were used to determine the zero points and colour terms.

The intermediate-band instrumental photometry was transformed to the standard AB photometric system following the prescriptions by Jacoby et al. (1987) and Alcalá et al. (2002) by using the equation

\begin{displaymath}m_{AB}(\lambda)=m_{0}(\lambda) + ZP_{\lambda}
\end{displaymath} (3)

where m0 is the instrumental magnitude in a given intermediate-band filter corrected for atmospheric extinction, and $ZP_{\lambda}$ the zero point derived by using the spectrophotometric standard stars Hiltner 600, LTT 4364, and Eg 274 (Hamuy et al. 1992), which were observed by us in the same run. The resulting magnitudes in the intermediate band filters are in the AB system.

The nightly calibration coefficients (atmospheric extinction coefficient, K, zero point, ZP, and colour terms, c), are reported in Table 2.

Table 1: Journal of the observations.
Field Date Filter $T_{\rm exp}$ Seeing X
(RA, Dec.) (d/m/y)   (min) ( $^{\prime\prime}$)  
ChaII_1$^\dagger$ 22/03/02 $R_{\rm c}$/162 $2\times5$ 1.2 1.48
(13:06:18, -76:44:21) 22/03/02 $I_{\rm c}$/Iwp $2\times5$ 1.2 1.48
  22/03/02 z+/61 $2\times5$ 1.2 1.49
  21/03/02 H$\alpha $/7 $5\times4$ 1.0 1.54
  21/03/02 H$\alpha $/12 $3\times5$ 1.0 1.51
ChaII_2 22/03/02 $R_{\rm c}$/162 $2\times5$ 1.2 1.55
(12:57:31, -76:44:20) 22/03/02 $I_{\rm c}$/Iwp $2\times5$ 1.2 1.53
  22/03/02 z+/61 $2\times5$ 1.2 1.51
  21/03/02 H$\alpha $/7 $5\times5$ 1.0 1.53
  21/03/02 H$\alpha $/12 $3\times5$ 1.0 1.56
  23/03/02 856/14 $4\times5$ 1.0 1.50
  23/03/02 914/27 $5\times5$ 1.0 1.48
ChaII_3 29/04/01 $R_{\rm c}$/162 $10\times5$ 2.7 1.84
(13:08:13, -77:14:27) 29/04/01 $I_{\rm c}$/Iwp $10\times5$ 2.7 1.69
  29/04/01 z+/61 $10\times5$ 2.7 1.60
  21/03/02 H$\alpha $/7 $5\times5$ 1.2 1.64
  21/03/02 H$\alpha $/12 $3\times5$ 1.2 1.59
  23/03/02 856/14 $4\times5$ 1.2 1.51
  23/03/02 914/27 $5\times5$ 1.2 1.53
ChaII_4 22/03/02 $R_{\rm c}$/162 $2\times5$ 1.4 1.49
(12:59:03, -77:13:58) 22/03/02 $I_{\rm c}$/Iwp $2\times5$ 1.4 1.50
  22/03/02 z+/61 $2\times5$ 1.4 1.50
  21/03/02 H$\alpha $/7 5$\times $5 1.2 1.50
  21/03/02 H$\alpha $/12 3$\times $4 1.2 1.51
  23/03/02 856/14 $4\times5$ 1.2 1.57
  24/03/02 914/27 $5\times5$ 1.2 1.54
ChaII_5 22/03/02 $R_{\rm c}$/162 $2\times5$ 1.2 1.60
(12:49:53, -77:13:59) 22/03/02 $I_{\rm c}$/Iwp $2\times5$ 1.2 1.62
  22/03/02 z+/61 $2\times5$ 1.2 1.66
  21/03/02 H$\alpha $/7 $5\times5$ 1.2 1.75
  21/03/02 H$\alpha $/12 $3\times5$ 1.2 1.66
  23/03/02 856/14 $4\times5$ 1.0 1.49
  23/03/02 914/27 $5\times5$ 1.0 1.50
ChaII_6 30/04/01 $R_{\rm c}$/162 $10\times6$ 2.0 1.67
(13:10:07, -77:45:08) 30/04/01 $I_{\rm c}$/Iwp $10\times6$ 2.0 1.57
  30/04/01 z+/61 $10\times6$ 2.0 1.52
  21/03/02 H$\alpha $/7 $5\times5$ 1.4 1.61
  21/03/02 H$\alpha $/12 $3\times5$ 1.4 1.66
  23/03/02 856/14 $4\times5$ 1.2 1.67
  23/03/02 914/27 $5\times5$ 1.2 1.74
ChaII_7 22/03/02 $R_{\rm c}$/162 $2\times4$ 1.2 1.78
(13:00:28, -77:44:39) 22/03/02 $I_{\rm c}$/Iwp $2\times5$ 1.2 1.74
  22/03/02 z+/61 $2\times5$ 1.2 1.69
  21/03/02 H$\alpha $/7 $5\times5$ 1.0 1.52
  21/03/02 H$\alpha $/12 $3\times5$ 1.0 1.51
  23/03/02 856/14 $4\times5$ 1.0 1.58
  23/03/02 914/27 $5\times5$ 1.0 1.62

$^\dagger$ Not observed in the 856-nm and 914-nm bands.

Table 2: Photometric calibration coefficients.
Filter Date K$^\dagger$ ZP c
R 04/27/2001 0.096 $24.418\pm0.004$ $-0.038\pm0.010$
" 04/28/2001 0.096 $24.388\pm0.001$ $-0.056\pm0.008$
" 04/29/2001 0.096 $24.445\pm0.004$ $-0.023\pm0.003$
" 03/21/2002 0.096 $24.397\pm0.003$ $-0.030\pm0.003$
I 04/27/2001 0.082 $23.413\pm0.012$ $0.222\pm0.002$
" 04/28/2001 0.082 $23.399\pm0.001$ $0.216\pm0.003$
" 04/29/2001 0.082 $23.431\pm0.001$ $0.221\pm0.007$
" 03/21/2002 0.082 $23.335\pm0.001$ $0.222\pm0.002$
z 04/27/2001 0.080 $21.546\pm0.804$ $0.336\pm0.311$
" 04/28/2001 0.080 $21.246\pm0.303$ $0.230\pm0.143$
" 04/29/2001 0.080 $21.248\pm0.353$ $0.214\pm0.140$
" 03/21/2002 0.080 $21.460\pm0.014$ $0.373\pm0.001$
$\rm H\alpha _7$ 03/20/2002 0.096 $21.249\pm0.095$ /
$\rm H\alpha _{12}$ 03/20/2002 0.096 $21.733\pm0.133$ /
$856 ~\rm nm$ 03/22/2002 0.080 $20.894\pm0.089$ /
$914~\rm nm$ 03/22/2002 0.080 $21.018\pm0.082$ /

$^\dagger$ The mean atmospheric extinction coefficients for La Silla have been adopted.

2.2.1 The z-band photometric calibration

The calibration of the z magnitudes required more effort, because there are no data for standard stars available in the literature. Since we aimed at determining the colour-magnitude diagrams (CMDs) of the sources in Cha II (Sect. 4) and their spectral energy distributions (SEDs, Sect. 5), we needed both to tie the z magnitudes with some "reference'' photometric system and to determine the flux at Earth of a star with magnitude z=0. To this aim, we observed the same Landolt standard fields in the z filter. In the Landolt SA 98 field, we selected the A0-type star SA 98 653, whose visual magnitude and colours are reported in Table 3. This star is characterised very well and its absolute spectrophotometry well-determined (Gutiérrez Moreno et al. 1988). For this A0-type star, we can assume (IC-z)=0, thus defining the "standard'' z-magnitude system as that for which the colour (IC-z) is zero for A0 type stars. The zero point of this calibration is then determined as

Zp=z-z0=IC-z0 (4)

where IC=9.522 is the standard magnitude of SA 98 653 in the Cousins system and z0 its z-band instrumental magnitude corrected for atmospheric extinction. We then applied this zero-point correction to the z0 instrumental magnitude of all the stars in the Landolt fields. In this way we obtained a catalogue of Landolt stars with z magnitudes and (I-z) colours in our internal WFI-Cousins system. We then used this catalogue to find the least-square solution to the equation

\begin{displaymath}z=z_0 + c_z \cdot (i_0-z_0) + ZP_z
\end{displaymath} (5)

where z0 and i0 are the standard stars instrumental magnitudes corrected for atmospheric extinction and ZPz and cz the z calibration coefficients for the newly-defined WFI-Cousins photometric system. The nightly mean values obtained for ZPz and cz are reported in Table 2.


Table 3: Standard photometry of SA 98 653 (RA = 06:52:05, Dec. = -00:18:18) from Stetson (2000).
V (B-V) (U-B) (V-RC) (R-I)C
9.539 -0.004 -0.097 0.007 0.008

Since the absolute spectrophotometry of SA 98 653 is available from Gutiérrez Moreno et al. (1988), we can use this star to obtain the absolute flux calibration in the z filter, i.e. to obtain the flux at Earth of a star with z=0. Following the procedure outlined in Appendix A (available only in electronic form) we derive

                    F(z=0) = $\displaystyle (8.4 \pm 0.1) \times 10^{-10}~ {\rm erg~ cm}^{-2}~ {\rm s}^{-1}~ \AA^{-1}$  
  = $\displaystyle (2608 \pm 31)~ {\rm Jy}.$  

This value can be used to derive the flux in the z-band of the sources in our catalogue using their corresponding z-band magnitudes.

2.3 The catalogue extraction

\end{figure} Figure 2: Photometric errors versus magnitudes and relative exponential fits for all the point-like sources detected in Cha II and for all the available filters. The "double sequence'', clearly visible in R and H $\alpha _{12}$ bands, is due to different seeing conditions in the individual exposures.

2.3.1 PSF-fitting procedure

The PSF-fitting photometry was performed by using the IRAF/daophot package (Stetson 1987). Since we are interested in the faintest stellar objects, the threshold level was defined in order to select all the sources having a signal-to-noise ratio ($\sigma $) greater than 3. This choice may lead to the extraction of many spurious detections, however daophot allows us to distinguish between stars, extended sources (galaxies), saturated objects, and other spurious detections. The PSF fitting procedure was re-iterated twice for each Cha II field; given the moderate crowding of these fields, this choice allowed us to detect the faintest neighbours to bright stars while keeping the number of spurious detections relatively small. The typical residual of the PSF-fitting is less than 2% of the peak brightness. A single catalogue comprising all the stellar sources detected in all the available photometric bands was finally produced. In Fig. 2 the internal photometric errors of all the detected point-like sources are plotted against the magnitude for all the available filters; the relative exponential fits are over-plotted. Table 4 summarises the number of point-like sources detected in the surveyed area in each filter and the limiting magnitudes achieved at the 10$\sigma $, 5$\sigma $, and 3$\sigma $ levels, respectively.

2.3.2 Completeness

The completeness of our catalogues was estimated in the standard way by inserting artificial stars into the images and recovering them using the same extraction parameters as for the real objects; the fraction of recovered artificial objects provides a measure of the completeness. We used the IRAF/addstar package to perform the exercise. We inserted 3000 artificial sources; this number should not alter the crowding statistics in the images significantly. The profile for the artificial sources was generated by using the same PSF model as used for the source extraction; the positions of the artificial objects are randomly distributed over the entire area of the mosaic, and their magnitudes range uniformly between the detection and the saturation limits in each band. As an example we show the results for the ChaII_2 field (see Table 1). Figure 3 shows the fraction of recovered artificial objects as a function of magnitude for each filter. The corresponding magnitude limits at 95% completeness level (C=95%) are reported in Table 4.


Table 4: Number of stellar sources (N*) detected in each filter and relative limiting magnitudes at the 10$\sigma $, 5$\sigma $, 3$\sigma $ and 95% completeness level.
Filter N* Mag 10$\sigma $ Mag 5$\sigma $ Mag 3$\sigma $ Mag (C=95%)
Rc/162 141 400 21.8 22.6 23.2 21.3
Ic/Iwp 141 000 21.0 21.7 22.3 21.0
z+/61 79 000 18.7 19.6 20.2 18.8
H$\alpha $/7 72 500 17.5 19.8 20.4 18.0
H$\alpha $/12 84 200 17.8 19.9 20.8 18.4
856-nm 76 700 18.6 19.4 20.0 18.0
914-nm 80 400 18.9 19.5 20.1 18.0

\end{figure} Figure 3: Completeness plots for extraction of artificial stars from the ChaII_2 mosaic (see Table 1) for all the photometric bands used in this work (see text).

3 Tools for candidate selection

Our primary criterion for selecting low-mass PMS star and BD candidates from the extracted catalogue was based on the comparison of the object location in CMDs with theoretical isochrones. In addition, our WFI data in the 865-nm and 914-nm intermediate-bands allowed us to obtain a first estimate of the effective temperature of the candidates, whereas the measurements in H$\alpha $ provided us with a diagnostics for possible H$\alpha $ emission. Our selection method exploits the tools now described.

3.1 Theoretical isochrones

Theoretical isochrones for low-mass stars and BDs down to 0.001 $M_{\odot}$ are provided by Baraffe et al. (1998) and Chabrier et al. (2000) in the Cousins photometric system (Bessel 1990). Since the transmission curves of the WFI filters are somewhat different from the original Cousins ones, in particular for the I-band, it is crucial to transform the colours and magnitudes into the appropriate photometric system. We thus transformed the isochrones by Baraffe et al. (1998) and Chabrier et al. (2000) into the WFI-Cousins system as described in Appendix B (available only in electronic form). In this way, we can use the isochrones to define the PMS locus in the CMDs (e.g. Fig. 9) in a photometrically consistent way.

In Fig. 4 the theoretical RC vs. (R-I)C CMD is shown. The 95% completeness limit was determined from the values reported in Table 4. For AV=0, our survey is complete down to 0.02 $M_{\odot}$ at the 95% level for objects younger than 10 Myr. For an average extinction $A_V\approx2$ mag for Cha II (Cambrèsy 1999), a 95% completeness limit corresponds to 0.03 $M_{\odot}$.

\end{figure} Figure 4: Theoretical RC vs. (R-I)C diagram. The isochrones (continuous curves) and PMS tracks (dashed curves), shifted to the distance modulus of Cha II (6.25 mag, Whittet et al. 1997) are in the WFI-Cousins photometric system. The dashed straight line represents the 95% completeness limit. The $A_{\rm V}=2$ mag reddening vector is shown.

3.2 The ( m856-m914) index vs. temperature relation

Late-type dwarfs are characterised by strong molecular absorption bands essentially due to metallic oxides (TiO and VO). The ESO-WFI medium-band filter centred at 856 nm covers important TiO absorption features that deepen with decreasing temperature, while the medium-band filter centred at 914 nm lies in a wavelength range that is relatively featureless in late-type objects (see Fig. B.1 lower panel, in Appendix B). Thus, the ( m856-m914) colour index is sensitive to the effective temperature for very cool objects (2000 K $\la$  $T_{\rm eff}\la 3800$ K), where TiO dominates the opacity (Burrows & Liebert 1993; Allard 1990; O'Neal et al. 1998). In the absence of interstellar extinction, this index can provide a first reliable guess for the temperature of the cool candidates.

Using the filter transmission curves for the corresponding intermediate-band filters and the synthetic low-resolution StarDusty spectra for low-mass stars and BDs by Allard et al. (2000), we derived the relationship between the ( m856-m914) index and the effective temperature, applying the methods described in Appendix B. The synthetic ( m856-m914) index, derived by integrating the StarDusty spectra under the filter transmissions and CCD quantum efficiency curves, is in the instrumental WFI system. As such, it was transformed into the standard AB photometric system by applying the zero points derived from the spectrophotometric standard stars Hiltner 600, LTT 4364, and Eg 274, observed with WFI in the 856-nm and 914-nm bands (see Sect. 2.2). From the standard fluxes (Hamuy et al. 1992) integrated under the WFI passbands, likewise the StarDusty spectra, we found an average zero-point correction for the ( m856-m914) colour of -0.13 mag, which allowed us to obtain the calibrated synthetic ( m856-m914) index as

\begin{displaymath}(m_{856}-m_{914})^{\rm synt} _{\rm calib}=(m_{856}-m_{914})^{\rm synt} _{\rm instr}-0.13.
\end{displaymath} (6)

In Fig. 5 we show the resulting ( m856-m914) index versus temperature relation. This relation depends on the stellar gravity, as indicated by the shaded area in Fig. 5, which encompasses all the calibrations for models in the range $3.5\leq \log g \leq5$. We also derived this relationship using other stellar atmosphere models (e.g. NextGen, COND, etc.); the StarDusty models were chosen because they match better the empirical relationship between the ( m856-m914) index and the effective temperature by López-Martí et al. (2004) (straight line in Fig. 5). Furthermore, Allard et al. (2001) found that silicate dust grains can form abundantly in the outer atmospheric layers of the latest M dwarfs and BDs.

In the case of high interstellar extinction ($A_V \ga 5$ mag), the temperature derived from this relation may be underestimated because the ( m856-m914) index becomes larger than in the absence of extinction. Thus, without any reliable evaluation of $A_{\rm V}$, we can use these temperature estimates only for candidate selection purposes (see Sect. 4).

\end{figure} Figure 5: The ( m856-m914) index versus effective temperature relation derived from the corresponding WFI intermediate-band filters and the synthetic low-resolution StarDusty spectra for low-mass stars and BDs by Allard et al. (2000). The two solid curves show the relation for $\log g=5.0$ (lower boundary) and $\log g=3.5$ (upper boundary); the dashed curve represents the $T_{\rm eff}$-colour relation for $\log g=4.0$. The straight line is the empirical relationship derived by López-Martí et al. (2004).

3.3 H$\alpha $ equivalent width calibration

The WFI H $\alpha_7$ filter covers the H$\alpha $ line, while the H $\alpha _{12}$ filter lies in a wavelength range that is affected by neither the H$\alpha $ wings nor strong photospheric lines (see Fig. B.1 mid panel, in Appendix B). The ( $\rm H\alpha _{12}-H\alpha _7$) index is thus sensitive to the intensity of the H$\alpha $ line. We used the StarDusty spectra by Allard et al. (2000) to define a calibration relation between this index and the H$\alpha $ equivalent width ( ${\it EW}_{\rm H\alpha}$). In order to simulate stellar spectra with different values of ${\it EW}_{\rm H\alpha}$, we added a Gaussian emission profile centred at H$\alpha $ (656.2 nm) and having a FWHM of 3 Å  to the synthetic spectra. The latter value was fixed by considering the H$\alpha $ line profiles of a sample of 75 PMS stars presented by Fernández et al. (1995). We found that the typical FWHM of the H$\alpha $ emission lines of these objects is around 3 Å, although it can be higher for classical T Tauri stars (5-6 Å). However, from tests made by using profiles with different FWHM, we have verified that the FWHM value has no significant effect on the (H $\alpha_{12}-{\rm H}\alpha_7$) vs. ${\it EW}_{\rm H\alpha}$ calibration.

By applying the method described in Appendix B and correcting the synthetic ( $\rm H\alpha _{12}-H\alpha _7$) colour by the mean ( $\rm H\alpha _{12}-H\alpha _7$) colour of the spectrophotometric standard stars used to transform the intermediate-band photometry into the AB system (see Sect. 2.2), we obtained the calibration relation ( $\rm H\alpha _{12}-H\alpha _7$) vs. ${\it EW}_{\rm H\alpha}$ displayed in Fig. 6. This relation does not depend on the FWHM of the Gaussian emission core but strongly depends on the effective temperature for objects cooler than about 4500 K. As expected from the contribution of the H$\alpha $ absorption wings (see Fig. B.1, mid panel), for a fixed ${\it EW}_{\rm H\alpha}$ the ( $\rm H\alpha _{12}-H\alpha _7$) index is higher for cooler objects; however, between 3000 and 2000 K, this trend is inverted because the $\rm H\alpha _{12}$ filter covers strong molecular absoption bands that heavily affect the continuum of very cool stars.

\end{figure} Figure 6: Calibration relation between the ( $\rm H\alpha _{12}-H\alpha _7$) index and the H$\alpha $ equivalent width derived from the corresponding WFI intermediate-band filters and the synthetic low-resolution StarDusty spectra for low-mass stars and BDs by Allard et al. (2000).

The relationship between the H$\alpha $ index and ${\it EW}_{\rm H\alpha}$ derived above was checked by using unpublished WFI H$\alpha $ photometry and the spectroscopy of already confirmed PMS stars in the L1616 cometary cloud obtained with HYDRA@WIYN3.5m (Alcalá et al. 2004) and FORS@VLT (Alcalá et al., in preparation). The spectra of the PMS stars in L1616 were integrated under the transmission curves of the WFI H$\alpha $ filters to derive their "synthetic'' ( $\rm H\alpha _{12}-H\alpha _7$) index. This index was then transformed into the AB system by using the mean H$\alpha $ colour index of the corresponding spectrophotometric standard stars observed in the L1616 run. The WFI H$\alpha $ photometry of these PMS stars was then used to compare the observed ( $\rm H\alpha _{12}-H\alpha _7$) index with the one derived from the spectra. Figure 7 shows the results. The best-fitting line (continuous line in Fig. 7) has a slope close to 1 and an offset close to zero, and this implies that no significant correction must be applied to the ( $\rm H\alpha _{12}-H\alpha _7$) vs. ${\it EW}_{\rm H\alpha}$ relation deduced from the synthetic spectra. Based on the errors of the observed ( $\rm H\alpha _{12}-H\alpha _7$) index, we conclude that bona-fide emission line objects will have an H$\alpha $ index greater than 0.1, corresponding to ${\it EW}_{\rm H\alpha} \approx 10 ~\AA$ for an object with $T_{\rm eff} \approx 4000$ K.

The H$\alpha $ index versus ${\it EW}_{\rm H\alpha}$ relation was also checked using data from the literature for the Cha II PMS stars. In Fig. 8 the H$\alpha $ equivalent width derived from spectroscopic data (Hughes & Hartigan 1992), ${\it EW}_{\rm H\alpha}\ (\rm spec)$, is compared with the one derived from the H$\alpha $ photometric index, ${\it EW}_{\rm H\alpha}\ (\rm phot)$. The effective temperatures reported by Hughes & Hartigan (1992) were used to estimate the latter, while the temperature reported by Barrado y Navascués & Jayawardhana (2004) was used for the PMS star C 41. From Fig. 8, the ${\it EW}_{\rm H\alpha}$ values derived from the H$\alpha $ photometric index are in good agreement with those derived from spectroscopic data. The computed RMS suggests that the relationship shown in Fig. 6 can be used to estimate the ${\it EW}_{\rm H\alpha}$ from the WFI H$\alpha $ photometry for objects with ${\it EW}_{\rm H\alpha} \ga 10 ~\AA$ and ${\it EW}_{\rm H\alpha} \ga 30 ~\AA$, at the 1$\sigma $ and 3$\sigma $ levels, respectively, provided we have an estimate of the stellar temperature.

4 Selection of the candidates

For the selection of low-mass PMS stars and young BD candidates, we exploited both the optical data from our survey and the JHK photometry available from the 2MASS catalogue (Cutri et al. 2003), using the tools described in the previous section. A matching radius of 0.5 $^{\prime\prime}$ was used to merge the WFI and the 2MASS catalogues, and this value was set by taking the astrometric accuracy of both catalogues into account.

Low-mass PMS star and BD candidates in SFRs can be identified in optical CMDs. In these diagrams, the colour increases rapidly for late spectral types; hence, the contamination from foreground stars is expected to decrease.

\end{figure} Figure 7: ( $\rm H\alpha _{12}-H\alpha _7$) observed colours for PMS stars in L1616 as a function of the synthetic ( $\rm H\alpha_{12}~-~H\alpha_{7}$) colours (see text). The continuous line represents the best fit of the data; the dashed line has the offset a=0 and slope b=1.

A first sample of candidates was then selected based on the IC vs. (IC-z) and IC vs. (R-I)C CMDs. All the point-like sources detected in Cha II above the 3$\sigma $ level were placed in these CMDs (Fig. 9) and their position was compared with the theoretical isochrones. The isochrones transformed into the WFI-Cousins system, as described in Sect. 3.1, were scaled to the Cha II distance of 178 pc (Whittet et al. 1997) in both CMDs. We then selected all the objects falling above the 10 Myr isochrone in both CMDs (Fig. 9). This criterion takes into account the age spread of the previously known PMS stars in Cha II (0.1-10 Myr, Hughes & Hartigan 1992) and the PMS locus that they define on the CMDs, as well as the uncertainties in interstellar extinction. In this way we selected 114 objects out of which 17 were previously known Cha II members.

\end{figure} Figure 8: Comparison between the H$\alpha $ equivalent width estimated in this work from the ( $\rm H\alpha _{12}-H\alpha _7$) photometric index and the one measured by Hughes & Hartigan (1992) and Barrado y Navascués & Jayawardhana (2004) for confirmed PMS stars in Cha II (open circles). The smaller black dots mark the objects recovered by our selection criteria (see Sect. 4).

\end{figure} Figure 9: IC vs. (IC-z) and IC vs. (R-I)C diagrams for the point-like objects in our survey (small gray dots). The lines represent the theoretical isochrones derived as explained in Sect. 3.1, shifted to the distance modulus of Cha II (6.25 mag, Whittet et al. 1997). The continuous line in each diagram represents the saturation limit. The big black dots represent the 37 sources selected from our criteria, while the open circles and squares represent the known PMS stars and candidates selected from previous surveys, respectively. The magnitudes reported by Hughes & Hartigan (1992) were used to plot the objects above and/or close to our saturation limits in any of the three broad-bands. The $A_{\rm V}=2$ mag reddening vector is shown in both diagrams.

\end{figure} Figure 10: (J-H) vs. (H-K) diagram for the sources detected in Cha II (small gray dots). The solid curve shows the relation between these indexes for main sequence stars (lower branch) and giants (upper branch). The dashed line indicates the classical T Tauri stars locus by Meyer et al. (1997). The continuous line is the dividing line between giants and dwarfs. The shaded area indicates the locus presumably uncontaminated by giants. The previously known PMS stars and candidates are indicated with open circles and squares respectively. The big black dots represent the objects selected in this work. The $A_{\rm V}=2$ mag reddening vector is shown. The two PMS stars which fall outside the selection area are the heavily veiled star Sz 47 (below the shaded area) and the binary Sz 60E (above the shaded area).

However, the 97 new candidates from the selection above may still be contaminated by background giants and highly-reddened objects. We therefore investigated the properties of the selected sample in the ${\it JHK}$ bands using their 2MASS magnitudes. It has been found by different authors that the intrinsic (J-H) and (H-K) colours of young late-M objects are dwarf-like with, in some cases, an (H-K) excess, mainly arising from a circumstellar disk or an in-falling envelope (Luhman & Rieke 1999; Meyer et al. 1997; Lee et al. 2005). Furthermore, Meyer et al. (1997) found that classical T Tauri stars, with prominent IR excesses due to circumstellar accretion disks, exhibit a narrow range of colours in the (J-H) vs. (H-K) diagram (Fig. 10). By using the (J-H) vs. (H-K) diagram we then performed a secondary selection on the sample coming from the optical CMDs, following similar criteria to those in Lee et al. (2005). The sub-sample of all the objects falling below the dividing line between giants and dwarfs (cf. Fig. 10) should be less contaminated by background giant stars. In this way we end up with a sub-sample of 84 objects, which still include the 17 previously known members.

The sub-sample passing the secondary selection may still contain foreground and highly-reddened background objects. We thus performed a third-level selection as follows. We assumed that the objects passing the primary and secondary selections above are all at the distance of Cha II and estimated their temperature and luminosity as described in Sect. 5. Although these are first guesses of the actual stellar temperatures and luminosities, these values already provide good estimates for candidate selection purposes. We then constructed the HR diagram of our candidates sample and selected all the objects falling between the birth-line and the 20 Myr isochrone, using the models by Baraffe et al. (1998) and Chabrier et al. (2000). Although the age distribution of the confirmed members of Cha II ranges between 0.1 and 10 Myr (Hughes & Hartigan 1992), the more relaxed criterion of a 20 Myr cut-off for the candidates with a high membership probability was set by taking the uncertainties on the distance to the Cha II cloud ( $\Delta d=18$ pc, Whittet et al. 1997) and the interstellar extinction into account. The evolutionary tracks of low-mass PMS stars run almost parallel to the luminosity axis of the H-R diagram. Thus, although mass estimates are not greatly affected by interstellar extinction, the age dispersion may be considerably increased if the adopted extinction for each candidate is not well-determined. The mean uncertainty on our visual extinction determinations is $\sim$1.5 mag (see Sect. 5.3.1); this would correspond to a luminosity uncertainty of $\sim$16% and, hence, an age uncertainty of $\sim$2 Myr for an object with mass $\sim$0.5 $M_{\odot}$ and age 2-3 Myr, i.e. the typical values in Cha II (Alcalá et al., in preparation). Thus, the criterion of the 20 Myr isochrone ensures that the selected objects are, to a first approximation, consistent with Cha II membership.

The theoretical isochrones depend on the physics involved in the models. As such, they may be rather uncertain, in particular in the very low-mass domain. We have thus tested the reliability of our selection criteria using as test bench the publicly available samples of confirmed PMS stars and BDs in Taurus and IC348 reported by Briceño et al. (2002) and Luhman et al. (2003), respectively. As can be seen in Fig. 14 (available only in electronic form), the selection criteria recover the vast majority of the previously known PMS stars and BDs in these regions, providing a good check that they work well on selecting these type of objects.

5 Results of the survey

The already confirmed Cha II population members consist of some 36 objects, comprised of 33 PMS stars (Hughes & Hartigan 1992; Alcalá et al. 2000; Young et al. 2005), one BD (Alcalá et al. 2006), and two planetary-mass objects (Allers et al. 2006; Jayawardhana & Ivanov 2006).

The screening of our data lead, on the other hand, to selection of 37 interesting objects. Of these, 17 were already known (16 PMS stars plus the BD) and 10 are candidates with high probability of membership based on their selection in previous surveys (Allers et al. 2006; Persi et al. 2003; Vuong et al. 2001; Young et al. 2005). The remaining 10 selected objects are completely new candidates. The other 13 previously known PMS stars are saturated and/or close to the saturation limit in at least one of our broad-band RIz images, so we could not apply our selection criteria. However, using their magnitudes and colours from the literature (Hughes & Hartigan 1992; Alcalá et al. 1995), these objects would fall within our selection; this means that, adding these to the 17 previously known non-saturated objects, our criteria recover 30 of the 36 previously known Cha II members.

There are a few exceptions in which the above selection criteria fail: the heavily veiled T Tauri stars Sz 47 (Hughes & Hartigan 1992) and C 41 (Barrado y Navascués & Jayawardhana 2004), which exhibit strong UV excess that affects their colours, and the Class-I source IRAS 12500-7658 (Young et al. 2005), which is a deeply embedded object. These three sources are sub-luminous in both CMDs. In addition, the other Class-I source, ISO-CHA II 28, is not detected in any of our images. Although this type of objects would escape our selection, they are expected to be rare in Cha II (Young et al. 2005; Alcalá et al., in preparation). The two planetary-mass candidates reported by Allers et al. (2006) also escaped selection because they were barely detected only in our I-band images. The Herbig-Haro object HH 54 (Giannini et al. 2006) and the Class-0 source BHR 86 (Garay et al. 2002) are also associated with Cha II. HH 54 was detected only in our R-band and H$\alpha $ images, while BHR 86 was not detected in any of the optical bands. These objects will be discussed in detail in a forthcoming paper (Alcalá et al., in preparation).

In conclusion, our selection criteria recover the majority (about 80%) of the confirmed and candidate members of Cha II reported in previous surveys. For the sake of clarity, in all the diagrams we use dots to represent the 37 objects selected with our criteria, open circles to represent the already confirmed members, and open squares to represent the 10 candidates selected in previous surveys. The 10 new candidates appear simply as dots in all the diagrams.

In Table 7 we report all the objects with at least one photometric measurement in the optical bands[*]. The table contains the confirmed members of Cha II, the candidates selected by previous surveys, and the new candidates selected in this work, as well as other sources identified as possible candidates in previous studies but that were rejected by our selection criteria. The latter sample includes objects like the candidates reported by López-Martí et al. (2005)[*] and some of the sources identified by Persi et al. (2003). Comments on some individual objects are given in Appendix C.

In the following we concentrate on the sample of 20 candidates, namely the 10 candidates selected from previous surveys, but also recovered from our selection criteria, and the 10 new candidates. They are reported in Table 6 and their spatial distribution is shown in Fig. 1. The 10 previously known candidates recovered in our selection have also been proposed as young members of Cha II by recent studies (Allers et al. 2006; Vuong et al. 2001; Young et al. 2005), which supports their membership to the cloud and, at the same time, the reliability of our selection method. Moreover, several of them possess H$\alpha $ emission (see Sect. 5.2). We note that these are the candidates with the highest membership likelihood, but we do not exclude the possibility that a few true Cha II members could have escaped some of our selection criteria, in particular if they are heavily veiled, deeply embedded, or if they are components of unresolved binaries. If all the 20 candidates are confirmed by spectroscopy, the population of Cha II will increase to some 56 members. Two of the 10 new candidates in our sample (WFI J12585611-7630105 and WFI J13005531-7708295) have been spectroscopically confirmed as PMS stars (Alcalá et al., in preparation).

5.1 Possible contaminants

One possible source of contamination of our candidates is represented by field dwarfs. In order to estimate the number of possible foreground dwarfs in the field of Cha II, we followed the prescriptions by Burgasser et al. (2004). We used the low-mass luminosity function simulations from Burgasser (2004) and the absolute magnitudes for late M and L field dwarfs from Dahn et al. (2002). Assuming a limiting magnitude of I=20, i.e. the I magnitude of our faintest candidates, and a mass function ${\rm d}N/{\rm d}M \propto M^\alpha$ with $0.5<\alpha<1.5$, we expect 8-10 foreground dwarfs with $2000< T_{\rm eff}<3000$ K in the $\sim$2 square-degree area observed in Cha II. Therefore, up to about 50% of the candidates might be foreground dwarfs unrelated to the SFR.

Another possible source of contamination is represented by faint galaxies that may have colours similar to those of PMS objects, in particular of BD candidates. From the K-band galaxy number counts toward the celestial south pole, the expected number of background galaxies with $K \la 13$ mag, i.e. the K magnitude of our faintest candidates, in a $\sim$2 square-degree area is $\sim$20 (Minezaki & Kobayashi 1998). Only diffuse interstellar extinction, which is negligible ($\sim$0.02 mag) at the K-band (Jones et al. 1981), is considered for this estimate. In Cha II, the extinction may be as high as $A_{\rm V}\approx 6$-8 mag, i.e. $A_{\rm K}\approx 1$ mag. Thus, in order to contaminate our sample, the background galaxies should have $K\la 12$ mag; the predicted galaxy number count at this magnitude is less than 2 in the surveyed area. In addition, our PSF extraction methods remove the extended objects quite efficiently; thus, the only contaminants may eventually be the point-like extra-galactic objects, mainly QSO's. According to Prescott et al. (2006), only a handful of QSO's are expected to be present in 2 square degrees for magnitudes brighter than R=20 mag. The number increases for fainter magnitudes, but the vast majority of our candidates are brighter. Thus, we do not expect that extra-galactic objects represent a major contamination problem.

5.2 H$\alpha $ emission of the candidates

The calibration relation between the ( $\rm H\alpha _{12}-H\alpha _7$) index and the ${\it EW}_{\rm H\alpha}$ (see Sect. 3.3) allows us to investigate a possible H$\alpha $ emission of the candidates. Based on this calibration, a threshold of ( $\rm H\alpha _{12}-H\alpha _7$) = 0.1 mag would translate into an equivalenth width ${\it EW}_{\rm H\alpha}\approx$ 10 Å at the 1$\sigma $ level of detection (cf. Sect. 3.3). Most of the previously known PMS stars and candidates have a larger $\rm H\alpha$ index than this value (cf. Fig. 11). Forty per cent of the 20 candidates show H$\alpha $ emission at the 1$\sigma $ level, but only 2 of the 20 show it at the 3$\sigma $ level. Considering that the H$\alpha $ line can be strongly variable in young objects, we cannot use the index as a major diagnostic for the selection of the candidates, but only as a consistency check on their possible PMS nature. However, the 1$\sigma $ level detections are still consistent with recent spectroscopic evidence of H$\alpha $ emission (Alcalá et al., in preparation). Moreover, the ${\it EW}_{\rm H\alpha}$ range covered by our candidates (10-35 Å) is in line with what is found by White & Basri (2003) for the slowly rotating, non-accreting stars and for BDs in Taurus-Auriga. For the remaining candidates, the ( $\rm H\alpha _{12}-H\alpha _7$) index does not indicate any clear evidence of H$\alpha $ emission. For three objects (WFI J12533662-7706393, WFI J12583675-7704065, and WFI J13005297-7709478), the level of H$\alpha $ emission could not be estimated because they are not detected in our H$\alpha $ images (the magnitude limit is $\sim$20.5 mag at the 3$\sigma $ level). These very faint candidates are indeed relevant because their optical colours resemble those of ISO-Cha II-13 (Alcalá et al. 2006). Note that ISO-Cha II-13 was also not detected in our H$\alpha $ images. The H$\alpha $ information drawn from our data for the selected candidates is reported in Table 6.

\end{figure} Figure 11: RC magnitudes versus ( $\rm H\alpha _{12}-H\alpha _7$) colours for the point-like objects in Cha II (small gray dots). The previously known PMS stars and candidates are indicated by open circles and squares, respectively, whereas the sources selected in this work are represented with big black dots. The $A_{\rm V}=2$ mag reddening vector is shown.

5.3 Towards the fraction of sub-stellar objects in ChaII

An interesting and important quantity in this investigation is the number of sub-stellar objects compared to the PMS stars in Cha II. In order to determine this quantity, we must single out the sub-stellar objects. According to Chabrier et al. (2000), the sub-stellar limit for 3-4 Myr objects, the approximate age of Cha II (Hughes & Hartigan 1992; Alcalá et al. 2000; Cieza et al. 2005) falls at a temperature of about 2900 K. Using the tools described in Sect. 3 and exploiting the potential of our photometric data, we attempted a first estimate of the temperature of the candidate members of Cha II and thus a guess at the fraction of sub-stellar candidates relative to the PMS stars. We first derived the dereddened SEDs of the candidates and then provided a first estimate of their temperature.


Table 5: Absolute flux calibration constants and effective wavelengths for the optical and near-IR pass-bands.
Filter $F^0 _{\lambda}$ $\lambda_{\rm eff}$ Ref.
  (erg $\cdot$ s-1 $\cdot$ cm-2 $\cdot$ Å-1) ($\mu $m)  
RC $2.25 \times 10^{-9}$ 0.64 Cousins (1976)
IC $1.19 \times 10^{-9}$ 0.79 Cousins (1976)
z $8.40 \times 10^{-10}$ 0.96 This work
J $3.13 \times 10^{-10}$ 1.25 Cutri et al. (2003)
H $1.13 \times 10^{-10}$ 1.62 Cutri et al. (2003)
K $4.28 \times 10^{-11}$ 2.20 Cutri et al. (2003)

5.3.1 Dereddened SEDs

The SEDs were derived by merging the RIz WFI-Cousins photometry with the 2MASS ${\it JHK}$ magnitudes. The standard fluxes at each of the optical and near-IR pass-bands have been derived from the observed magnitudes using the expression $F_{\lambda} = F^0 _{\lambda} \cdot 10^{-0.4 \cdot {\rm mag}^{\rm corr} _{\lambda}}$, where mag $^{\rm corr} _{\lambda}$ is the observed magnitude corrected for interstellar reddening and $F^0 _{\lambda}$ the Earth flux of an A0-type star of magnitude V=0 (Table 5). In order to derive the visual extinction, AV, a SED minimization procedure was used as follows. We assume that the observed magnitude at each wavelength ( $m^{\lambda} _{\rm obs}$) can be derived by the following equation

 \begin{displaymath}m^{\lambda}_{\rm obs} - m^{\lambda}_{\rm ref} = A_V \cdot \frac{{\rm d}A_\lambda}{{\rm d}A_V} + S
\end{displaymath} (7)

where $m^{\lambda}_{\rm ref}$ is a distance-scaled reference magnitude at a given wavelength ($\lambda$) and temperature ( $T_{\rm eff}$). In Eq. (7), AV is the visual extinction, $A_\lambda$ the wavelength-dependent extinction, and S a scaling factor depending only on the stellar radius. Using a grid of reference magnitudes and temperatures, we perform a linear fit to the previous expression where $\frac{{\rm d}A_\lambda}{{\rm d}A_V}$ and $m^{\lambda}_{\rm obs} - m^{\lambda}_{\rm ref}$ are the independent and dependent variables, respectively, and AV and S the two free parameters of the fit. The $\frac{{\rm d}A_\lambda}{{\rm d}A_V}$ term was derived using the extinction law by Cardelli et al. (1989), assuming the standard value RV=3.1. The grid of reference SEDs was constructed by combining the tabulations of colours as functions of spectral types by Kenyon & Hartmann (1995), Luhman et al. (2003), and Ducati et al. (2001). The AV value of the reference SED minimizing the $\chi^2$ of the fit corresponds to the best approximation of the actual visual extinction. As a check, we compared the results from our fitting procedure with the extinction derived from the spectral types and the observed colours of previously known PMS stars in Cha II. The RMS difference in AV is less than 1.5 mag, with only one object having a residual of about 4 mag.

In Fig. 12 the SEDs of the 10 new candidates selected in this work are shown. For six of these, we found a match in the c2d Spitzer catalogues, and their IRAC and/or MIPS fluxes are also included. As appears from their SEDs, no significant IR excess is detected in these six objects. If they are spectroscopically confirmed as Cha II members, they will most likely correspond to weak-line T Tauri stars (WTTS) or to BDs with thin disks. For the four remaining objects lacking Spitzer data, we cannot assess the presence of strong IR excess. However, for three of these we estimate an H$\alpha $ equivalenth width greater than about 20 Å.

It is interesting to note that the number of WTTS relative to CTTS in Cha II is very small. Even assuming that all the candidates in our sample (see Table 6) were to result in WTTS, which is not likely because several of them show strong H$\alpha $ emission, the CTTS would outnumber the WTTS by a factor $\sim$2. This would contrast with what it is found in other clouds, like Cha I and Taurus, where WTTS dominate the PMS population (Feigelson & Montmerle 1999). Our selection based on the (J-H) vs. (H-K) diagram picks up objects later than about K5. This might lead to the conclusion that some earlier type WTTS could have escaped our selection. However, such objects would have been detected in the previous ROSAT X-ray surveys by Alcalá et al. (2000,1995). Only two WTTS in the cloud, namely RXJ1301.0-7654a and RXJ1303.1-7706 were discovered in these surveys. Thus, the low fraction of WTTS in Cha II seems to be real. More details on this will be discussed in a future c2d paper (Alcalá et al., in preparation), with a more complete set of SEDs for the Cha II members.

\end{figure} Figure 12: Spectral energy distributions of the 10 new candidates selected in this work. Dereddened optical data are represented with filled circles; the IR data are from 2MASS (filled squares), IRAC@Spitzer and MIPS@Spitzer (filled triangles and upside down triangles, respectively). The open circles represent the observed fluxes. The best-fitting StarDusty spectra by Allard et al. (2000) with the same temperature as the objects are over-plotted (see Sect. 5.3.2). WFI J12585611-7630105, WFI J13014752-7631023, and WFI J13031616-7629381 fall outside the areas mapped with MIPS and IRAC; although WFI J13071960-7655476 and WFI J13005297-7709478 fall in the area surveyed with MIPS, they were not detected.

5.3.2 Temperature

The temperature of the candidates was estimated using the $T_{\rm eff}$ vs. ( m856-m914) calibration relation reported in Sect. 3.2. This calibration is valid for cool objects with temperatures in the range 2000-3800 K, corresponding to a dereddened ( m856-m914) index in the range 1.16-0.16 mag. Thus, we applied the calibration to objects with a dereddened index ( m856-m914) > 0.16 mag. The $A_{\rm V}$ values determined as explained in Sect. 5.3.1 and the extinction law by Cardelli et al. (1989) were used to deredden the ( m856-m914) index. For objects with ( m856-m914) < 0.16 mag, i.e. hotter than about 3800 K, a temperature estimate was done by fitting a grid of reference SEDs to the dereddened SEDs. The grid of reference SEDs was constructed by combining the tabulations of colours as a function of spectral type by Kenyon & Hartmann (1995), Luhman et al. (2003), and Ducati et al. (2001); the best approximation to the dereddened SED was then obtained by $\chi^2$-minimization of the flux differences between the dereddened SED and the reference SED. The minimization procedure was applied to the short-wavelength portion ( $\lambda \leq \lambda_{J}$) of the SED, which is less affected by an eventual IR excess. The method turns out to be accurate within 250 K relative to the spectroscopic temperature estimates for most of the previously known PMS stars. This is shown in the upper panel of Fig. 13, where the residuals between the temperature derived from the dereddened SED fitting, $T_{\rm fit}$, and that obtained from spectroscopy by Hughes & Hartigan (1992), $T_{\rm spec}$, are shown. For ISO-CHA II 13 we adopt the temperature value determined spectroscopically by Alcalá et al. (2006).

For the cool candidates, we also performed a consistency check between the temperature derived from SED fitting and what is coming from the $T_{\rm eff}$ vs. ( m856-m914) calibration. We find that the two methods yield consistent results within $\sim$200 K, as can be seen from the lower panel of Fig. 13. We can then use the $T_{\rm eff}$ vs. ( m856-m914) calibration confidently for those objects with dereddened ( m856-m914) > 0.16 mag.

The resulting temperature for each of the candidates is reported in Table 6. Using these estimates, we attempted to determine their radius and luminosity as follows:


First the reference SED of the same effective temperature as the object was selected and scaled to the Cha II distance (d=178 pc, Whittet et al. 1997):

 \begin{displaymath}F'_{\rm ref}=F_{\rm ref} \cdot \left ( \frac{R^\star}{d} \right )^2
\end{displaymath} (8)

where $R^\star$ is the stellar radius. Since $R^\star$ is an unknown quantity, we let it vary between 0.1, i.e. the value expected for a BD at the Deuterium burning limit (Baraffe et al. 2003), and 10 $R_{\odot}$; for the objects whose effective temperature could not be estimated from the ( m856-m914) index, both $T_{\rm eff}$ and $R^\star$ were set as free parameters in the assumption that the short-wavelength side of the dereddened SED ( $\lambda \leq \lambda_{J}$) is mainly influenced by $T_{\rm eff}$.

The best estimate of $R^\star$ was then obtained by minimizing the flux differences between the observed dereddened SED and the reference one. The minimization procedure was performed, as above, on the short-wavelength portion ( $\lambda \leq \lambda_{J}$) of the SED;

Finally, we calculated the bolometric luminosities by applying bolometric corrections to both the dereddened I-band and J-band data. Bolometric corrections were taken from Kenyon & Hartmann (1995) for objects earlier than M6 (i.e. $T_{\rm eff} \ga 3000$ K), while for cooler objects we followed the prescriptions by Luhman (1999) and combined the compilations of bolometric corrections as a function of temperature by Bessel (1991), Monet et al. (1992), Tinney et al. (1993), and Leggett et al. (1996).
Having an estimate of the temperature and luminosity of the candidates, we determined their masses and ages by comparison with the set of evolutionary tracks by Baraffe et al. (1998) and Chabrier et al. (2000) on the HR diagram.

\end{figure} Figure 13: Upper panel: comparison between the effective temperature resulting from the SED fitting procedure described in Sect. 5.3.2 and that derived from the spectral classification for the confirmed T Tauri stars in Cha II reported by Hughes & Hartigan (1992) and, for C 41, by Barrado y Navascués & Jayawardhana (2004) (open circles). The black dots mark the objects recovered by our selection criteria. The point at $T_{\rm spec}\approx 2900$ K represents ISO-CHA II 13 as reported by Alcalá et al. (2006). Lower Panel: comparison between the effective temperature resulting from the SED fitting procedure described in Sect. 5.3.2 and that derived from the $T_{\rm eff}$ vs. ( m856-m914) calibration. The open circles and squares represent previously known PMS stars and candidates, respectively. The black dots mark the objects recovered by our selection criteria.

Table 6: The 20 candidates selected in this work.

${\it EW}_{\rm H\alpha}$ $A_{\rm V}$ $T_{\rm eff}$ Alit V $T^{lit} _{\rm eff}$ Ref. Classification
  ($\AA$) (mag) (K) (mag) (K)   in this work
WFI J12533662-7706393 Not Det. 5.58 $\dagger$3000 $\pm$ 250       PMS star cand.
C 17 No Em. 5.95 $\star$3600 $\pm$ 200 5.1   b PMS star cand.
C 33 No Em. 4.48 $\dagger$3700 $\pm$ 250 2.4   b PMS star cand.
IRAS 12535-7623 No Em. 3.38 $\star$3900 $\pm$ 200 5   a PMS star cand.
WFI J12583675-7704065 Not Det. 3.84 $\dagger$2400 $\pm$ 200       BD cand.
WFI J12585611-7630105 30 1.91 $\dagger$2900 $\pm$ 200       BD cand.
ISO-CHA II 29 No Em. 5.94 $\star$3900 $\pm$ 200 11.5   b PMS star cand.
WFI J12591881-7704419 No Em. 4.98 $\star$3800 $\pm$ 200       PMS star cand.
WFI J12592348-7726589 No Em. 3.33 $\dagger$3700 $\pm$ 250       PMS star cand.
WFI J13005297-7709478 Not Det. 5.50 $\dagger$2500 $\pm$ 200       BD cand.
WFI J13005531-7708295 No Em. 4.13 $\star$3800 $\pm$ 200       PMS star cand.
IRAS F12571-7657 No Em. 3.95 $\star\star$>3000 8,13   b, c PMS star cand.
WFI J13014752-7631023 No Em. 0.88 $\dagger$2800 $\pm$ 200       BD cand.
C 50 20 3.46 $\dagger$3000 $\pm$ 200 2.0   b PMS star cand.
C 51 20 1.16 $\dagger$3300 $\pm$ 250 2.3   b PMS star cand.
WFI J13031615-7629381 15 0.61 $\dagger$2900 $\pm$ 200       BD cand.
C 62 35 3.53 $\dagger$3000 $\pm$ 200 4.1, 4 3140 b, c PMS star cand.
WFI J13071960-7655476 25 0.50 $\star$3000 $\pm$ 200       PMS star cand.
C 66 25 3.87 $\dagger$2900 $\pm$ 200 3.7, 0 2793 b, c BD cand.
2MASS13125238-7739182 10 1.36 $\dagger$3400 $\pm$ 250     d PMS star cand.
Notes to Col. 4: $\dagger$ Temperature derived from the dereddened ( m856-m914) colour index (Sect. 3.2). $\star$ Temperature derived from the SED fitting method described in Sect. 5.3.2. $\star\star$ See Appendix C.
References: (a) Larson et al. (1998); (b) Vuong et al. (2001); (c) Allers et al. (2006); (d) Young et al. (2005).

5.4 The fraction of sub-stellar objects

Previous investigations (Alcalá et al. 2000; Hughes & Hartigan 1992) showed some evidence that the mass spectrum in Cha II might be biased towards very-low mass objects. It is thus interesting to investigate what our data suggest on this matter.

Assuming that all the 20 candidates are true cloud members, we end up with 6 candidates with $T_{\rm eff} \leq2900$ K, i.e. below the substellar limit for 5 Myr old objects. The remaining 14 objects turn out to be PMS star candidates (see Table 6). Taking the BD ISO-CHA II 13 and the 33 confirmed PMS stars into account, the resulting fraction of sub-stellar objects ( $0.02 ~M_{\odot} < M < 0.08~ M_{\odot}$) relative to the PMS stars ( $0.08 ~M_{\odot} < M < 10 ~M_{\odot}$), $R_{ss}=\frac{N(0.02{-}0.08)}{N(0.08-10)}$, in Cha II would be $\sim$15% (19% if we consider the two planetary-mass objects reported by Jayawardhana & Ivanov 2006). This last value would imply a peculiar mass spectrum for the population of Cha II relative to other T associations (where Rss=12-14%, López-Martí et al. 2004; Briceño et al. 2002), in particular in the sub-stellar domain. If we assume that only the candidates having H$\alpha $ emission are true cloud members, the Rss value would drop to $\sim$10% (15% if we consider the two planetary-mass objects reported by Jayawardhana & Ivanov 2006), i.e. more similar to that measured in other T associations and lower than measured in OB associations ($\sim$26%, Barrado y Navascués et al. 2001; Boulanger et al. 1998; Hillenbrand & Carpenter 2000; Muench et al. 2002; Briceño et al. 2002). If Rss in Cha II is similar to other T associations, we expect that at least 50% of the candidates reported in Table 6 will be genuine Cha II members.

Assuming that the 7 IRAS sources with H$\alpha $ emission discussed in Appendix C are PMS stars, the fraction of substellar objects would drop to about 8% or 13% depending on whether the two planetary-mass objects are counted among the sub-stellar ones.

By comparing the position of the candidates in the HR diagram with the Baraffe et al. (1998) and Chabrier et al. (2000) evolutionary tracks, their age distribution would peak between 2 and 3 Myr, if they were genuine Cha II members. This would be fairly consistent with the average age of 3.6 Myr determined by Cieza et al. (2005) for a sample of T Tauri stars in Cha II. Taking the candidates into account, the mean mass in Cha II would be $\sim$0.5 $M_{\odot}$, i.e. slightly lower than the mean mass of the confirmed members alone (0.6 $M_{\odot}$) and comparable to the mean mass for the Cha I population ($\sim$0.45 $M_{\odot}$[*]).

6 Conclusions

The optical survey presented here is one of the deepest and more extensive conducted so far in Cha II, and it constitutes the optical ancillary data for the c2d Spitzer Legacy survey in this cloud. The photometric selection based on our optical imaging data, combined with data in the near-IR from the 2MASS, allowed us to recover basically all the previously known members of the cloud, including 10 candidates from previous IR surveys, with a high membership probability. Furthermore, we provided 10 new likely member that have not been detected by previous surveys, thereby increasing the total number of candidates to 20. Should all these objects be spectroscopically confirmed as PMS stars and young BDs, the population of Cha II will increase to 56 members.

Up to 50% of the sample may be contaminated by field dwarfs. According to our characterisation criteria of the candidates, we estimate that at least about 50% of them will result in true Cha II members and, based on our temperature estimates, several of these objects are expected to be sub-stellar. Under these assumptions, we conclude that the fraction of substellar objects relative to the PMS stars in Cha II is $\sim$19%, i.e. more than reported for other T associations like Taurus. In the most conservative hypothesis in which only the candidates showing H$\alpha $ emission would result in true members of the cloud, this fraction would drop to $\sim$15%, i.e. comparable to what is measured in other T associations.

From the completeness of our survey, both in space and flux, we conclude that the optical population of Cha II members discovered so far is nearly complete. We do not expect to find many more PMS stars based on optical observations. More clues to these issues will be addressed in a forthcoming paper by combining the optical data presented here with those from the MIPS@Spitzer and IRAC@Spitzer observations in Cha II (Alcalá et al., in preparation).

This work was partially financed by the Istituto Nazionale di Astrofisica (INAF) and the Italian Ministero dell'Istruzione, Universitá e Ricerca (MIUR). We thank the anonymous referee for his/her constructive comments and suggestions. We also thank M. Radovich for many explanations of the use of ASTROMETRIX. We thank N. Evans, PI of the c2d Spitzer Legacy Program, and L. Cieza for many useful comments and suggestions on an earlier version of the paper. We also thank F. Comerón and E. Marilli for discussions. L. Spezzi acknowledges financial support from COFIN-MIUR-2004 (The X-Shooter spectrograph for the VLT) and PRIN-INAF-2005 (Stellar clusters: a benchmark for star formation and stellar evolution) and a grant by the Europen Southern Observatory (ESO) for a two-month stay in Garching. Support for this work, part of the Spitzer Legacy Science Program, was provided by NASA through contract 1224608 issued by the Jet Propulsion Laboratory, California Institute of Technology, under NASA contract 1407. We also thank the Lorentz Center in Leiden for hosting several meetings that contributed to this paper. This research made use of the SIMBAD database, operated at the CDS (Strasbourg, France) and the data products from the Two-Micron All-Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center at the California Institute of Technology, funded by NASA and the National Science Foundation.



7 Online Material

Table 7: Optical photometry of objects in the Cha II dark cloud.
Designation RA J2000 Dec J2000 $R_{\rm c}$ $I_{\rm c}$ z $\rm H\alpha _7$ H $\alpha _{12}$ m856 m914 Main Ref. Note
  (hh:mm:ss) (dd:mm:ss)                  

IRAS 12416-7703

12:45:06.43 -77:20:13.52 S S S 11.86 $\pm$ 0.09 12.10 $\pm$ 0.13 S S i $\bullet$
IRAS 12448-7650 12:48:25.70 -77:06:36.72 S S S 13.15 $\pm$ 0.09 13.55 $\pm$ 0.13 S S i $\bullet$
IRAS F12488-7658 12:52:30.49 -77:15:12.92 14.84 $\pm$ 0.01 12.21 $\pm$ 0.02$^\ast$ S 15.48 $\pm$ 0.10 15.60 $\pm$ 0.13 12.27 $\pm$ 0.09 S f, i $\bullet$
IRAS 12496-7650 12:53:17.17 -77:07:10.63 16.43 $\pm$ 0.01 13.92 $\pm$ 0.01 12.82 $\pm$ 0.02 16.12 $\pm$ 0.10 16.84 $\pm$ 0.14 13.52 $\pm$ 0.10 13.41 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
WFI J12533662-7706393 12:53:36.62 -77:06:39.31 22.22 $\pm$ 0.14 19.35 $\pm$ 0.04 18.00 $\pm$ 0.04 ND ND 19.26 $\pm$ 0.20 18.58 $\pm$ 0.12 m $\leftthumbsup$NEW
C 17 12:53:38.84 -77:15:53.21 17.08 $\pm$ 0.01 15.09 $\pm$ 0.01 14.23 $\pm$ 0.01 17.67 $\pm$ 0.10 17.36 $\pm$ 0.13 15.10 $\pm$ 0.09 14.79 $\pm$ 0.08 f $\bullet$ $\leftthumbsup$
IRAS 12500-7658 12:53:42.79 -77:15:11.59 21.54 $\pm$ 0.08 19.40 $\pm$ 0.01 18.44 $\pm$ 0.06 20.81 $\pm$ 0.38 21.24 $\pm$ 0.46 19.43 $\pm$ 0.23 19.06 $\pm$ 0.15 i $\bigstar$ $\ding{74}$
ChaII 304 12:55:16.00 -76:46:21.83 21.85 $\pm$ 0.09 19.28 $\pm$ 0.02 18.27 $\pm$ 0.05 ND ND 19.26 $\pm$ 0.17 18.65 $\pm$ 0.11 j $\bullet$ $\ding{68}$ $\leftthumbsdown$
ChaII 305 12:55:16.48 -76:46:20.89 21.82 $\pm$ 0.09 19.30 $\pm$ 0.01 18.29 $\pm$ 0.05 ND ND 19.20 $\pm$ 0.17 18.68 $\pm$ 0.12 j $\bullet$ $\ding{68}$ $\leftthumbsdown$
C 33 12:55:25.72 -77:00:46.62 16.86 $\pm$ 0.01 15.16 $\pm$ 0.01 14.66 $\pm$ 0.01 17.15 $\pm$ 0.22 16.93 $\pm$ 0.13 15.22 $\pm$ 0.09 14.79 $\pm$ 0.08 f $\bullet$ $\leftthumbsup$
2MASS12560549-7654106 12:56:05.43 -76:54:10.69 15.95 $\pm$ 0.01 14.54 $\pm$ 0.01 14.04 $\pm$ 0.02 16.17 $\pm$ 0.10 16.15 $\pm$ 0.13 14.58 $\pm$ 0.09 14.42 $\pm$ 0.08 i $\bullet$ $\leftthumbsdown$
Sz 46NW 12:56:32.85 -76:45:44.78 18.45 $\pm$ 0.01 17.69 $\pm$ 0.17 17.56 $\pm$ 0.04 18.72 $\pm$ 0.13 18.70 $\pm$ 0.14 17.96 $\pm$ 0.10 18.07 $\pm$ 0.09 a $\bullet$ $\ding{68}$ $\leftthumbsdown$
Sz 46N 12:56:33.59 -76:45:45.18 14.61 $\pm$ 0.01 13.16 $\pm$ 0.01 12.68 $\pm$ 0.01 14.80 $\pm$ 0.10 15.07 $\pm$ 0.13 13.24 $\pm$ 0.09 13.08 $\pm$ 0.08 a $\bigstar$ $\ding{68}$ $\leftthumbsup$
Sz 46S 12:56:33.64 -76:45:49.54 16.81 $\pm$ 0.01 16.05 $\pm$ 0.02 15.93 $\pm$ 0.02 17.07 $\pm$ 0.10 17.03 $\pm$ 0.13 16.32 $\pm$ 0.09 16.38 $\pm$ 0.08 a $\bullet$ $\ding{68}$ $\leftthumbsdown$
Sz 47 12:56:58.63 -76:47:06.72 13.86 $\pm$ 0.01$^\ast$ 13.38 $\pm$ 0.01$^\ast$ 13.45 $\pm$ 0.01 13.68 $\pm$ 0.10 14.07 $\pm$ 0.13 13.72 $\pm$ 0.09 13.86 $\pm$ 0.08 a $\bigstar$ $\ding{80}$
IRAS 12533-7632 12:57:00.49 -76:48:35.10 21.19 $\pm$ 0.12 19.76 $\pm$ 0.01 ND ND 20.90 $\pm$ 0.38 19.92 $\pm$ 0.33 19.83 $\pm$ 0.29 e $\bullet$
IRAS 12535-7623 12:57:11.64 -76:40:11.14 13.61 $\pm$ 0.01 12.13 $\pm$ 0.01 11.53 $\pm$ 0.02 13.69 $\pm$ 0.10 13.73 $\pm$ 0.13 12.22 $\pm$ 0.09 12.03 $\pm$ 0.08 e, i $\bullet$ $\leftthumbsup$
ISO-CHA II 13 12:58:06.67 -77:09:09.22 22.51 $\pm$ 0.27 19.38 $\pm$ 0.01 17.89 $\pm$ 0.04 ND ND 19.44 $\pm$ 0.19 18.33 $\pm$ 0.11 g, k $\bigstar$ $\leftthumbsup$
WFI J12583675-7704065 12:58:36.75 -77:04:06.53 22.45 $\pm$ 0.25 19.32 $\pm$ 0.01 17.85 $\pm$ 0.04 ND ND 19.51 $\pm$ 0.20 18.27 $\pm$ 0.10 m $\leftthumbsup$NEW
WFI J12585611-7630105 12:58:56.11 -76:30:10.48 16.39 $\pm$ 0.01 14.38 $\pm$ 0.01 13.54 $\pm$ 0.02 16.08 $\pm$ 0.09 16.60 $\pm$ 0.13 14.34 $\pm$ 0.09 13.81 $\pm$ 0.08 m $\leftthumbsup$NEW
C 41 12:59:09.86 -76:51:03.49 18.51 $\pm$ 0.01 17.03 $\pm$ 0.01 16.45 $\pm$ 0.02 18.06 $\pm$ 0.10 18.88 $\pm$ 0.14 17.31 $\pm$ 0.09 17.07 $\pm$ 0.08 f,h $\bigstar$ $\ding{80}$
ISO-CHA II 29 12:59:10.19 -77:12:13.72 16.19 $\pm$ 0.02 14.23 $\pm$ 0.01 13.36 $\pm$ 0.01 16.41 $\pm$ 0.19 16.32 $\pm$ 0.13 14.37 $\pm$ 0.09 13.91 $\pm$ 0.08 g,i $\bullet$ $\leftthumbsup$
WFI J12591881-7704419 12:59:18.81 -77:04:41.92 15.69 $\pm$ 0.01 14.09 $\pm$ 0.01 13.34 $\pm$ 0.01 15.86 $\pm$ 0.10 15.88 $\pm$ 0.17 14.19 $\pm$ 0.09 13.86 $\pm$ 0.09 m $\leftthumbsup$NEW
WFI J12592348-7726589 12:59:23.48 -77:26:58.96 14.41 $\pm$ 0.01 12.99 $\pm$ 0.01 12.39 $\pm$ 0.01 14.78 $\pm$ 0.12 14.62 $\pm$ 0.12 13.13 $\pm$ 0.09 12.75 $\pm$ 0.08 m $\leftthumbsup$NEW
IRAS 12556-7731 12:59:26.50 -77:47:08.70 S S S 12.34 $\pm$ 0.09 12.66 $\pm$ 0.13 S S i $\bullet$
WFI J13005297-7709478 13:00:52.97 -77:09:47.77 22.63 $\pm$ 0.30 19.64 $\pm$ 0.07 18.25 $\pm$ 0.05 ND ND 19.85 $\pm$ 0.25 18.66 $\pm$ 0.12 m $\leftthumbsup$NEW
Sz 48NE 13:00:53.15 -77:09:09.18 16.08 $\pm$ 0.01 14.32 $\pm$ 0.04 13.55 $\pm$ 0.01 16.11 $\pm$ 0.09 16.24 $\pm$ 0.13 14.44 $\pm$ 0.09 14.02 $\pm$ 0.08 a $\bigstar$ $\ding{68}$ $\leftthumbsup$
Sz 49 13:00:53.26 -76:54:14.98 14.86 $\pm$ 0.01 13.53 $\pm$ 0.03 13.01 $\pm$ 0.01 13.90 $\pm$ 0.09 14.93 $\pm$ 0.13 13.62 $\pm$ 0.09 13.35 $\pm$ 0.08 a $\bigstar$ $\ding{80}$
Sz 48SW 13:00:53.56 -77:09:08.28 16.26 $\pm$ 0.01 14.50 $\pm$ 0.02 13.72 $\pm$ 0.01 16.13 $\pm$ 0.09 16.57 $\pm$ 0.13 14.51 $\pm$ 0.09 14.19 $\pm$ 0.08 a $\bigstar$ $\ding{68}$ $\leftthumbsup$
Sz 50 13:00:55.28 -77:10:22.01 14.42 $\pm$ 0.01 12.50 $\pm$ 0.01$^\ast$ 11.96 $\pm$ 0.02 14.34 $\pm$ 0.10 14.74 $\pm$ 0.13 12.79 $\pm$ 0.09 12.43 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
WFI J13005531-7708295 13:00:55.31 -77:08:29.54 14.93 $\pm$ 0.01 13.29 $\pm$ 0.01 12.64 $\pm$ 0.01 15.26 $\pm$ 0.14 15.24 $\pm$ 0.13 13.39 $\pm$ 0.09 13.15 $\pm$ 0.09 m $\leftthumbsup$NEW
RXJ1301.0-7654a 13:00:56.22 -76:54:01.76 11.59 $\pm$ 0.01$^\ast$ 10.54 $\pm$ 0.01$^\ast$ 10.17 $\pm$ 0.01 11.63 $\pm$ 0.09 11.71 $\pm$ 0.14 10.61 $\pm$ 0.09 10.66 $\pm$ 0.09 d $\bigstar$ $\leftthumbsup$
IRAS F12571-7657 13:00:59.21 -77:14:02.80 18.62 $\pm$ 0.01 16.16 $\pm$ 0.01 15.00 $\pm$ 0.01 19.26 $\pm$ 0.14 19.23 $\pm$ 0.16 16.48 $\pm$ 0.23 15.79 $\pm$ 0.14 e, g $\bullet$ $\leftthumbsup$
ISO-CHA II 73 13:01:46.03 -77:16:02.89 16.31 $\pm$ 0.02 15.12 $\pm$ 0.01 14.68 $\pm$ 0.01 16.71 $\pm$ 0.22 16.55 $\pm$ 0.13 15.38 $\pm$ 0.09 15.13 $\pm$ 0.08 g $\bullet$ $\leftthumbsdown$
WFI J13014752-7631023 13:01:47.52 -76:31:02.32 16.83 $\pm$ 0.05 15.23 $\pm$ 0.01 14.65 $\pm$ 0.06 16.84 $\pm$ 0.20 17.05 $\pm$ 0.21 15.77 $\pm$ 0.17 15.18 $\pm$ 0.12 m $\leftthumbsup$NEW
Sz 51 13:01:58.94 -77:51:21.74 13.47 $\pm$ 0.02$^\ast$ 12.38 $\pm$ 0.01$^\ast$ 11.80 $\pm$ 0.02 12.46 $\pm$ 0.09 13.37 $\pm$ 0.13 12.55 $\pm$ 0.09 12.29 $\pm$ 0.08 a $\bigstar$
CM Cha 13:02:13.49 -76:37:57.68 12.93 $\pm$ 0.01 11.84 $\pm$ 0.01 11.18 $\pm$ 0.01 12.15 $\pm$ 0.10 12.57 $\pm$ 0.13 11.55 $\pm$ 0.09 11.62 $\pm$ 0.08 b $\bigstar$ $\leftthumbsup$
C 50 13:02:22.82 -77:34:49.51 17.78 $\pm$ 0.01 15.49 $\pm$ 0.02 14.53 $\pm$ 0.01 17.78 $\pm$ 0.10 18.19 $\pm$ 0.20 15.48 $\pm$ 0.09 14.93 $\pm$ 0.08 f, g $\bullet$ $\leftthumbsup$
IRAS 12589-7646 13:02:47.73 -77:02:46.32 S S S 12.72 $\pm$ 0.09 12.92 $\pm$ 0.13 S S i $\bullet$
RXJ1303.1-7706 13:03:04.46 -77:07:02.75 12.27 $\pm$ 0.01$^\ast$ 11.17 $\pm$ 0.01$^\ast$ 10.80 $\pm$ 0.01 12.34 $\pm$ 0.09 12.08 $\pm$ 0.13 11.29 $\pm$ 0.09 11.16 $\pm$ 0.08 c $\bigstar$ $\leftthumbsup$
C 51 13:03:09.04 -77:55:59.52 16.88 $\pm$ 0.01 14.47 $\pm$ 0.01 13.46 $\pm$ 0.02 16.82 $\pm$ 0.10 17.16 $\pm$ 0.13 14.39 $\pm$ 0.09 14.03 $\pm$ 0.10 f $\bullet$ $\leftthumbsup$
ChaII 376 13:03:12.45 -76:50:50.82 16.31 $\pm$ 0.01 15.45 $\pm$ 0.01 15.35 $\pm$ 0.05 16.40 $\pm$ 0.10 16.42 $\pm$ 0.13 NO NO j $\bullet$ $\leftthumbsdown$
WFI J13031615-7629381 13:03:16.15 -76:29:38.15 16.42 $\pm$ 0.01 14.41 $\pm$ 0.03 13.68 $\pm$ 0.03 16.30 $\pm$ 0.10 16.67 $\pm$ 0.15 NO NO m $\leftthumbsup$NEW
ISO-CHA II 98a 13:03:25.85 -77:01:48.36 17.09 $\pm$ 0.02 16.42 $\pm$ 0.04 16.24 $\pm$ 0.03 17.53 $\pm$ 0.10 17.27 $\pm$ 0.11 16.83 $\pm$ 0.09 17.17 $\pm$ 0.24 g $\bullet$ $\ding{68}$ $\leftthumbsdown$
ISO-CHA II 98b 13:03:26.10 -77:01:48.40 17.09 $\pm$ 0.02 16.37 $\pm$ 0.01 16.18 $\pm$ 0.02 17.47 $\pm$ 0.10 17.24 $\pm$ 0.11 17.68 $\pm$ 0.43 16.99 $\pm$ 0.22 g $\bullet$ $\ding{68}$ $\leftthumbsdown$
ISO-CHA II 110 13:04:18.96 -76:53:59.96 14.52 $\pm$ 0.01 14.03 $\pm$ 0.01 14.03 $\pm$ 0.01 14.59 $\pm$ 0.09 14.58 $\pm$ 0.13 NO NO g $\bullet$ $\leftthumbsdown$
Hn 22 13:04:22.78 -76:50:05.86 13.51 $\pm$ 0.01$^\ast$ 12.68 $\pm$ 0.01$^\ast$ 12.17 $\pm$ 0.02 13.23 $\pm$ 0.09 13.85 $\pm$ 0.13 NO NO b $\bigstar$ $\ding{68}$
Hn 23 13:04:24.06 -76:50:01.39 12.23 $\pm$ 0.01$^\ast$ 11.33 $\pm$ 0.08$^\ast$ S 12.11 $\pm$ 0.09 12.34 $\pm$ 0.13 NO NO b $\bigstar$ $\ding{68}$ $\leftthumbsup$
Sz 52 13:04:24.84 -77:52:30.11 16.89 $\pm$ 0.01 14.94 $\pm$ 0.01 14.12 $\pm$ 0.01 17.00 $\pm$ 0.10 17.35 $\pm$ 0.13 14.84 $\pm$ 0.09 14.43 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
Hn 24 13:04:55.74 -77:39:49.21 13.30 $\pm$ 0.06$^\ast$ 11.95 $\pm$ 0.02$^\ast$ 11.50 $\pm$ 0.02 13.49 $\pm$ 0.09 13.43 $\pm$ 0.14 12.19 $\pm$ 0.09 12.01 $\pm$ 0.08 b $\bigstar$ $\leftthumbsup$
Hn 25 13:05:08.51 -77:33:42.66 16.11 $\pm$ 0.01 14.17 $\pm$ 0.02 13.36 $\pm$ 0.01 16.45 $\pm$ 0.12 16.61 $\pm$ 0.13 14.40 $\pm$ 0.09 14.01 $\pm$ 0.09 b $\bigstar$ $\leftthumbsup$
Sz 53 13:05:12.66 -77:30:52.56 15.58 $\pm$ 0.02 13.86 $\pm$ 0.02 13.15 $\pm$ 0.01 15.58 $\pm$ 0.09 15.90 $\pm$ 0.18 15.16 $\pm$ 0.12 14.47 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
Sz 54 13:05:20.80 -77:39:01.48 11.58 $\pm$ 0.02$^\ast$ 10.61 $\pm$ 0.06$^\ast$ 10.47 $\pm$ 0.43 11.28 $\pm$ 0.10 11.57 $\pm$ 0.13 10.80 $\pm$ 0.09 11.05 $\pm$ 0.15 a $\bigstar$ $\leftthumbsup$
Sz 55 13:06:30.49 -77:34:00.12 16.37 $\pm$ 0.01 14.70 $\pm$ 0.06 14.06 $\pm$ 0.47 15.65 $\pm$ 0.09 16.96 $\pm$ 0.13 14.61 $\pm$ 0.09 14.53 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
Sz 56 13:06:38.70 -77:30:35.39 15.41 $\pm$ 0.01$^\ast$ 13.47 $\pm$ 0.02$^\ast$ 12.71 $\pm$ 0.44 15.41 $\pm$ 0.09 15.75 $\pm$ 0.13 13.67 $\pm$ 0.09 13.15 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
Sz 57 13:06:56.56 -77:23:09.46 15.97 $\pm$ 0.01 13.65 $\pm$ 0.01 12.59 $\pm$ 0.41 15.71 $\pm$ 0.09 16.23 $\pm$ 0.13 13.74 $\pm$ 0.09 13.07 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
Sz 58 13:06:57.35 -77:23:41.46 14.63 $\pm$ 0.01 13.16 $\pm$ 0.01 12.53 $\pm$ 0.43 14.44 $\pm$ 0.09 14.58 $\pm$ 0.13 12.98 $\pm$ 0.09 12.78 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
Sz 59 13:07:09.23 -77:30:30.24 13.42 $\pm$ 0.02$^\ast$ 12.08 $\pm$ 0.01$^\ast$ 11.51 $\pm$ 0.43 12.96 $\pm$ 0.10 13.39 $\pm$ 0.13 12.57 $\pm$ 0.09 12.50 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
C 62 13:07:18.04 -77:40:53.00 18.60 $\pm$ 0.01 16.10 $\pm$ 0.01 15.00 $\pm$ 0.44 18.42 $\pm$ 0.11 18.97 $\pm$ 0.17 16.23 $\pm$ 0.09 15.60 $\pm$ 0.08 f, l $\bullet$ $\leftthumbsup$
WFI J13071960-7655476 13:07:19.60 -76:55:47.64 18.33 $\pm$ 0.01 16.49 $\pm$ 0.02 15.77 $\pm$ 0.02 18.12 $\pm$ 0.10 18.57 $\pm$ 0.14 NO NO m $\leftthumbsup$NEW
Sz 60W 13:07:22.30 -77:37:22.62 14.88 $\pm$ 0.02$^\ast$ 13.45 $\pm$ 0.02$^\ast$ 13.02 $\pm$ 0.46 14.60 $\pm$ 0.09 15.01 $\pm$ 0.13 13.42 $\pm$ 0.09 13.23 $\pm$ 0.07 a $\bigstar$ $\ding{68}$ $\leftthumbsup$
Sz 60E 13:07:23.33 -77:37:23.20 15.32 $\pm$ 0.02$^\ast$ 13.60 $\pm$ 0.01$^\ast$ 12.94 $\pm$ 0.46 15.01 $\pm$ 0.09 15.69 $\pm$ 0.13 13.98 $\pm$ 0.09 13.61 $\pm$ 0.08 a $\bigstar$ $\ding{68}$
Hn 26 13:07:48.50 -77:41:21.73 16.07 $\pm$ 0.01 14.31 $\pm$ 0.01 13.54 $\pm$ 0.46 15.89 $\pm$ 0.11 16.25 $\pm$ 0.13 14.41 $\pm$ 0.09 14.13 $\pm$ 0.08 b $\bigstar$ $\leftthumbsup$
Sz 61 13:08:06.33 -77:55:05.05 13.69 $\pm$ 0.02$^\ast$ 12.38 $\pm$ 0.01$^\ast$ 11.48 $\pm$ 0.42 12.69 $\pm$ 0.09 13.47 $\pm$ 0.13 12.27 $\pm$ 0.09 12.14 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
C 66 13:08:27.19 -77:43:23.41 18.91 $\pm$ 0.01 16.51 $\pm$ 0.01 15.44 $\pm$ 0.44 19.24 $\pm$ 0.14 19.72 $\pm$ 0.22 16.76 $\pm$ 0.09 16.03 $\pm$ 0.08 f, l $\bullet$ $\leftthumbsup$
IRAS F13052-7653NW 13:09:09.81 -77:09:43.52 15.37 $\pm$ 0.01 14.04 $\pm$ 0.02 13.52 $\pm$ 0.44 15.59 $\pm$ 0.09 16.11 $\pm$ 0.13 14.02 $\pm$ 0.09 13.84 $\pm$ 0.08 e $\bullet$ $\ding{68}$
IRAS F13052-7653S 13:09:10.67 -77:09:46.84 13.49 $\pm$ 0.01 12.57 $\pm$ 0.01 12.38 $\pm$ 0.47 13.09 $\pm$ 0.09 13.04 $\pm$ 0.13 12.69 $\pm$ 0.09 12.77 $\pm$ 0.08 e $\bullet$ $\ding{68}$
IRAS F13052-7653N 13:09:10.98 -77:09:44.14 13.47 $\pm$ 0.01 12.42 $\pm$ 0.02 12.07 $\pm$ 0.45 13.51 $\pm$ 0.09 13.74 $\pm$ 0.13 12.54 $\pm$ 0.09 12.47 $\pm$ 0.08 e $\bullet$ $\ding{68}$
Sz 62 13:09:50.44 -77:57:23.94 14.03 $\pm$ 0.02$^\ast$ 12.56 $\pm$ 0.02$^\ast$ 11.66 $\pm$ 0.43 NO NO NO NO a $\bigstar$ $\leftthumbsup$
Sz 63 13:10:04.12 -77:10:44.62 14.66 $\pm$ 0.01 13.21 $\pm$ 0.01 12.66 $\pm$ 0.44 14.40 $\pm$ 0.09 15.09 $\pm$ 0.13 13.25 $\pm$ 0.09 13.03 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
2MASS13102531-7729082 13:10:25.26 -77:29:08.70 16.49 $\pm$ 0.01 15.36 $\pm$ 0.05 14.95 $\pm$ 0.48 16.49 $\pm$ 0.10 16.49 $\pm$ 0.13 15.54 $\pm$ 0.09 15.43 $\pm$ 0.08 i $\bullet$ $\leftthumbsdown$
2MASS13110329-7653330 13:11:03.27 -76:53:32.89 14.50 $\pm$ 0.01 13.85 $\pm$ 0.00 13.68 $\pm$ 0.01 14.62 $\pm$ 0.09 14.57 $\pm$ 0.13 NO NO i $\bullet$ $\leftthumbsdown$
2MASS13125238-7739182 13:12:52.37 -77:39:18.58 15.36 $\pm$ 0.01 13.62 $\pm$ 0.02 12.97 $\pm$ 0.47 15.57 $\pm$ 0.09 15.78 $\pm$ 0.13 13.83 $\pm$ 0.09 13.51 $\pm$ 0.08 i $\bullet$ $\leftthumbsup$
Sz 64 13:14:03.83 -77:53:07.48 15.82 $\pm$ 0.01 14.01 $\pm$ 0.01 13.16 $\pm$ 0.45 15.25 $\pm$ 0.09 16.29 $\pm$ 0.13 14.40 $\pm$ 0.09 13.65 $\pm$ 0.08 a $\bigstar$ $\leftthumbsup$
Notes: The positions for most objects are from the R-band images. For saturated objects in the R-band the positions are from the H $\alpha _{12}$ images; the coordinates of Sz 62, not observed in H$\alpha $, are from the z-band image. Magnitudes marked with an asterisk are taken from the the literature (see the relative references) and correspond to objects saturated in our images. Comments about some of these objects are given in Appendix C. Symbols and labels are as follows:
  • $\bigstar$ = Confirmed member
  • $\bullet$ = Previously known candidate
  • $\leftthumbsup$ = Object recovered by our selection criteria
  • $\leftthumbsup$NEW = Member candidate firstly identified in this work

  • $\leftthumbsdown$ = Object rejected on the basis of our selection criteria
  • $\ding{80}$ = Veiled object
  • $\ding{74}$ = Embedded object
  • $\ding{68}$ = Component of a visual binary or multiple system
  • S = Saturated object
  • ND = Object not detected
  • NO = Object not observed

Main references: a Hughes & Hartigan (1992); b Hartigan (1993); c Alcalá et al. (1995); d Covino et al. (1997); e Alcalá et al. (2000); f Vuong et al. (2001); g Persi et al. (2003); h Barrado y Navascués & Jayawardhana (2004); i Young et al. (2005); j López-Martí et al. (2005); k Alcalá et al. (2006); l Allers et al. (2006); m This work.

\end{figure} Figure 14: IC vs. (R-I)C and (J-H) vs. (H-K) diagrams for the PMS populations in Taurus and IC348. The isochrones, shifted by the corresponding distance modulus, are as in Fig. 9. The selection limits in the (J-H) vs. (H-K) diagrams are as in Fig. 10.

Appendix A: Flux at z = 0

In order to derive the flux at zero magnitude in the z band we used the following procedure:

The spectrophotometry of SA 98 653 is available from Gutiérrez Moreno et al. (1988) in the form of monochromatic instrumental magnitudes, covering a wavelength range from 3200 to 8000 Å, in steps of 40 Å for $\lambda \le 5400$ Å  and 80 Å  for $\lambda > 5400 $ Å. Since the z band is centred at 9648 Å  (FWHM = 616 Å), we need the spectrophotometry in a wider wavelength range. We thus extrapolate the SED of SA 98 653, up to 16 000 Å, by using the Kurucz grid of stellar atmosphere models (Kurucz 1979) in the following way. We first deredden the SED of SA 98 653 using the canonical absorption-to-reddening ratio $R_{\rm V}=3.2$, as suggested by Chen et al. (2001) for the SA 98 field. Assuming that SA 98 653 is a main sequence star, as indicated by its spectrophotometry and broad-band photometric indexes, we derive its distance by comparing its observed V magnitude with the expected absolute visual magnitude (MV=0.7, Gray 1992) for an A0-type main sequence star. Assuming an error of two spectral sub-classes on SA 98 653 spectral type, we constrained its distance between 300 and 450 pc, i.e. $0.30~{\rm mag}< A_{\rm V} <0.45$ mag. Note that the extinction of the nearby star SA 98-667 (spectral type B9V) is $A_{\rm V}=0.33$ calculated from its Hipparcos parallax (Chen et al. 2001), which is consistent with the range of values allowed for SA 98 653. The interstellar extinction coefficients at each wavelength were derived by using the extinction curve by Cardelli et al. (1989). We then performed a two-parameters least-square fit to the observed spectrophotometry of SA 98 653 by taking as free parameters the extinction (in the range $0.30~{\rm mag}< A_{\rm V} <0.45$ mag in steps of 0.01 mag) and the effective temperature of the model spectrum. Both the Kurucz models and the dereddened SED of SA 98 653 were normalised to the flux at 5556 Å. A minimum $\chi^2$ for a model with $T_{\rm eff}=11~000$ K, $\log g=4$ and $A_{\rm V}=0.33$ was derived (Fig. A.1). We then used this synthetic spectrum reddened with $A_{\rm V}=0.33$ to extrapolate the spectrophotometry of SA 98 653 in the near-IR.

\end{figure} Figure A.1: Observed (dots) and dereddened (open triangles) spectral energy distribution of SA 98 653 normalised to the flux at 5556 Å. The best fitting model by Kurucz (1979) is over-plotted (continuous line).

The absolute Earth flux at 5556 Å  of a star with magnitude V* calibrated with Vega, the primary photometric standard, is given by (Gray 1992)

 \begin{displaymath}\log(F^*_{5556})=-0.400 \cdot V^*-8.451
\end{displaymath} (A.1)

in units of erg $\cdot$ cm-2 $\cdot$ s-1 $\cdot$  Å-1.

When this expression is used for stars whose effective temperature is different from Vega ( $T_{\rm eff}=9790$ K), as in the case of SA 98 653, a small systematic error is introduced because the effective wavelength of the V band is 5480 Å, i.e. 76 Å  shorter than the reference wavelength (5556 Å). An adequate correction for this is given by the following colour term (Gray 1992)

 \begin{displaymath}\log \left (\frac{F^*_{5556}}{F^*_{5480}} \right )= -0.006+0.018 \cdot (B-V)^* .
\end{displaymath} (A.2)

Considering the dereddened visual magnitude of SA 98 653 and its intrinsic (B-V), and applying Eqs. (A.1) and (A.2), respectively, we obtain

\begin{displaymath}V^{\rm corr} _{\rm SA~98~653} = 9.24


\begin{displaymath}\log[F^{5556} _{\rm SA~98~653}]= -12.1465.

We now can calculate the flux of SA 98 653 at any wavelength with respect to the flux at 5556 Å  as

\begin{displaymath}F^{\rm synt} _{\rm SA~98~653} (\lambda) =F^{5556} _{\rm SA~98... 10^{-0.4 \cdot \rm mag^{\rm synt}_{\rm SA~98~653}(\lambda)}
\end{displaymath} (A.3)

where mag $^{\rm synt} _{\rm SA~98~653}$ is the synthetic magnitude of SA 98 653 at each wavelength normalised to the magnitude at 5556 Å, determined in the previous step.

We finally determined the synthetic flux of SA 98 653 in the z band by integrating its synthetic spectrum $F^{\rm synt} _{\rm SA~98~653} (\lambda)$ in the z filter transmission curve T $_{\rm\lambda}$, i.e. the product of the detector quantum efficiency and the filter transmission curve:

\begin{displaymath}F^z _{\rm SA~98~653}=\frac{1}{A_z} \cdot \int_{\lambda_1} ^{\...
..._{\rm SA~98~653}(\lambda)\cdot T_{\lambda} \cdot \rm d\lambda}
\end{displaymath} (A.4)


\begin{displaymath}A_z=\int_{\lambda_1} ^{\lambda_2} {T_{\lambda} \cdot \rm d\lambda}.
\end{displaymath} (A.5)

We obtained

\begin{displaymath}F^z _{\rm SA~98~653}=1.62 \times 10^{-13} \hspace{0.5cm} \rm erg \cdot cm^{-2} \cdot s^{-1} \cdot \AA^{-1}

where $F^{z} _{\rm SA~98~653}$ is the flux in the z filter of an A0 type star with visual magnitude V=9.539 and z=IC=9.522. What we need, in order to obtain an absolute flux calibration for the z band, is the flux of an A0-type star with V=0 and z=0 ( Fz A0V). Since SA 98 653 is an A0-type star and we have defined its intrinsic colour as (IC-z)=0, we have

\begin{displaymath}z_{\rm SA~98~653} \stackrel{{\rm def}}{=} I_{\rm SA~98~653} = 9.522.
\end{displaymath} (A.6)

Thus, we derive Fz A0V as

\begin{displaymath}F^z _{A0V} = F^z _{\rm SA~98~653} \times 10^{(0.4 \cdot z_{\rm SA~98~653})}.
\end{displaymath} (A.7)

By considering the uncertainties on magnitude and interstellar extinction measurements, we obtain
                         Fz A0V = $\displaystyle (8.4 \pm 0.1) \times 10^{-10} \hspace{0.25cm} {\rm erg} \hspace{0.1cm} {\rm cm}^{-2}~ {\rm s}^{-1} {\rm\AA}^{-1}$  
  = $\displaystyle (2608 \pm 31)\ {\rm Jy}.$  

A slightly smaller absolute flux calibration constant for the WFI z band has been estimated by F. Comerón (private communication) using the absolute flux calibration of Vega obtained by Bohlin & Gilliland (2004) with the Space Telescope Imaging Spectrograph (STIS). Using these data we derive a value of $7.2 \times 10^{-10}$ erg cm-2 s-1 Å-1. However, for the sake of homogeneity, we prefer to use the value we derive above.

Appendix B: Isochrones

In order to transform the colours and magnitudes of the isochrones reported by Baraffe et al. (1998) and Chabrier et al. (2000) into the WFI-Cousins system we determined first the WFI synthetic magnitudes, $m_{\Delta \lambda}$, as

 \begin{displaymath}m_{\Delta \lambda} = -2.5 \cdot \log_{10} (f_{\Delta \lambda}) + C_{\Delta \lambda}
\end{displaymath} (B.1)

where $C_{\Delta \lambda}$ is the absolute calibration constant of the photometric system, tied to the Earth flux of an A0-type star with magnitude V=0 (see Table 5), and $f_{\Delta \lambda}$ is the observed flux.

In order to compute the expected flux of a PMS star or a young BD in the pass-band $\Delta \lambda$ of a given filter we proceeded as follows. We first determined the transmission curve, $T_{\lambda}$, i.e. the product of the detector quantum efficiency and the given filter transmission curve[*]. We then compute

\begin{displaymath}f_{\Delta \lambda} = \frac{1}{A_{\Delta \lambda}}
\cdot \int...
... ^{\lambda_2} F_{\lambda} \cdot T_{\lambda} \cdot \rm d\lambda
\end{displaymath} (B.2)


\begin{displaymath}A_{\Delta \lambda} = \int_{\lambda_1} ^{\lambda_2} T_{\lambda} \cdot \rm d\lambda
\end{displaymath} (B.3)

and $F_{\lambda}$ is the absolute flux of a given PMS star or BD. For this purpose we used the synthetic low-resolution spectra for low-mass stars by Hauschildt et al. (1999) calculated with their NextGen model-atmosphere code. For simulating very cool objects (i.e. $T_{\rm eff}<3000$ K) we used the StarDusty and BD-Dusty atmosphere models by Allard et al. (2000), which include dust opacity.

\end{figure} Figure B.1: The R, I, z (upper panel), $\rm H\alpha _7$, $\rm H\alpha _{12}$ (central panel), 856 and 914 nm (lower panel) WFI+ESO2.2m transmission bands. Examples of two normalised NextGen spectra (Hauschildt et al. 1999) with $T_{\rm eff}=2500$ K and $T_{\rm eff}=6000$ K are over-plotted.

\end{figure} Figure B.2: The Cousins colours as a function of the WFI instrumental colours for the standard stars in the Landolt's field SA 98. The continuous lines display the best fit of the data while the dashed lines have slope equal to 1 (i.e. no correction).

As an example of our procedure we show in Fig. B.1 (upper panel) two model spectra and the R, I and z WFI+ESO2.2m transmission curves over-plotted. However, the flux $f_{\Delta \lambda}$ that we obtained in this way is the one at the surface of the star. The observed flux at any distance (d) is given by

\begin{displaymath}f_{\Delta \lambda} (d) = f_{\Delta \lambda} \cdot \left (\frac{R^\star}{d} \right )^2
\end{displaymath} (B.4)

where $R^\star$ is the object radius. The values for $R^\star$ can be computed from theoretical PMS evolutionary tracks, which in our case are those by Baraffe et al. (1998) for low-mass stars and those by Chabrier et al. (2000) for sub-stellar objects (i.e. $M\la 0.1~ M_{\odot}$). Finally, setting d=10 pc and using Eq. (B.1), one can derive the absolute magnitudes for a given $T_{\rm eff}$, age and $R^\star$ and hence, determine the isochrones in the CMDs.

Note that, consistently with the RIz photometric calibration of the sources in Cha II, the absolute calibration constants used in Eq. (B.1) are those of the WFI-Cousins photometric system (see Table 5). However, a non-negligible difference between the I-WFI and the I-Cousins pass-bands arises from the fact that the red cut-off of the I-WFI filter is defined by the CCD response and it is centred on a redder wavelength (0.85 $\mu $m) with respect to the standard Cousins I filter (0.79 $\mu $m). In order to take into account this difference in the determination of the theoretical isochrones, we have exploited the relations between the WFI-Cousins (R-I)C and (IC-z) standard colours and the WFI instrumental colours for the standard stars in the Landolt's fields. By a least-square linear fitting (Fig. B.2), we derive the coefficients which allow us to correct the colours of the computed isochrones

\begin{displaymath}(R-I)_{\rm synt} = 0.75 \cdot (R-I)_{\rm instr}
\end{displaymath} (B.5)

\begin{displaymath}(I-z)_{\rm synt} = 1.10 \cdot (I-z)_{\rm instr}.
\end{displaymath} (B.6)

In this way we were able to report in the WFI-Cousins photometric system the isochrones originally calculated by Baraffe et al. (1998) and Chabrier et al. (2000) in the Cousins system of Bessel (Bessel 1990). This allows us to use the isochrones in a photometrically homogeneous way.

In Fig. B.3, the 1, 5 and 10 Myr isochrones by Chabrier et al. (2000) are compared with those derived by us in the WFI-Cousins photometric system. Note the large difference, in particular for the coolest objects, i.e. (R-I)C >1.7, where the use of the isochrones in the Bessel-Cousins system would have produced many spurious member candidates.

In order to check our procedure, we used the standard filters of Bessel (1990) and we were successful in reproducing the isochrones presented by Chabrier et al. (2000) (see the smaller panel in Fig. B.3).

\end{figure} Figure B.3: Theoretical isochrones by Baraffe et al. (1998) and Chabrier et al. (2000) for the Cousins photometric system of Bessel (continuous line) and for the WFI-Cousins system (dashed lines, this work). In the smaller panel the 5 Myr isochrone by Chabrier et al. (2000) (continuous line) is compared with that reproduced by our algorithm when using the same Bessel filters (dashed line).

Appendix C: Comments on individual objects

IRAS 12416-7703, IRAS 12448-7650, IRAS F12488-7658, IRAS 12556-7731, IRAS 12589-7646

These objects were selected as member candidates of Cha II by Young et al. (2005) on the basis of MIPS@Spitzer observations. IRAS F12488-7658/C 13 was also previously proposed as a young object by (Vuong et al. 2001). We cannot conclude whether our selection criteria are satisfied for these objects because they are saturated in at least one of the three broad-bands RIz. However, their ( $\rm H\alpha _{12}-H\alpha _7$) indices suggest H$\alpha $ emission (see Table 7) at $\sim$2$\sigma $ level and their location on the (J-H) vs. (H-K) diagram is consistent with the PMS stars locus (see Sect. 4 and Fig. 10).

ChaII 304, ChaII 305, ChaII 376

ChaII 376 was the only H$\alpha $ emitter detected by López-Martí et al. (2005) in their WFI pointing. They also proposed that their candidates ChaII 304 and ChaII 305 are two low-mass BDs or planetary-mass objects close to the Deuterium burning limit. These two objects, however, do not satisfy our selection criteria. The early classification of these objects as possible Cha II members by López-Martí et al. (2005), though based on WFI images taken with the same set of filters as in our survey, is probably not accurate because the isochrones in the Cousins system of Bessel (Chabrier et al. 2000) used for the selection are inadequate for the WFI-Cousins photometric system (see Appendix B). From their dereddened ( m856-m914) colour index we estimate an effective temperature of 2800 K and 3000 K for ChaII 304 and ChaII 305, respectively. Assuming these temperature values, we find that the location of the two objects on the HR diagram would be inconsistent with membership to Cha II. We found that the optical colours of ChaII 376 are inconsistent with membership, as also asserted by López-Martí et al. (2005), and we do not detect any significant H$\alpha $ emission for this object. Moreover, its R-band magnitude (16.31 mag) is remarkably brighter than the one reported by López-Martí et al. (2005) (18.38 mag). Since the magnitude of this object in the other bands are in fair agreement with ours, we think that the R-band magnitude reported in López-Martí et al. (2005) may be incorrect. This may explain why ChaII 376 was selected as their only H$\alpha $ emitter.

Sz 46

This is one of the 19 PMS stars originally identified by Schwartz (1977) in Cha II. Subsequent photometric observations by Hughes & Hartigan (1992) showed that the object is a visual triple system. The H$\alpha $ emission star originally discovered by Schwartz (1977), Sz 46N, has two companions, namely Sz 46NW and Sz 46S; both are bluer than Sz 46N and are probably unrelated to it (Hughes & Hartigan 1992). Based on our analysis, both Sz 46NW and Sz 46S have optical colours inconsistent with membership to the Cha II cloud, their temperature is well above 4000 K according to their ( m856-m914) index and the ( $\rm H\alpha _{12}-H\alpha _7$) colour suggests no H$\alpha $ emission. Thus, Sz 46NW and Sz 46S are most probably not PMS stars and hence, not physically associated with Sz 46N.

IRAS 12533-7632, IRAS F13052-7653

These objects were detected as ROSAT X-rays sources (Alcalá et al. 2000) and proposed as possible members of Cha II by different authors (Larson et al. 1998; Prusti et al. 1992; Young et al. 2005).

IRAS 12533-7632 appears on our images as a point-like object associated with a nebulosity, consistently with what reported in Prusti et al. (1992). However, it was not detected in our z and H$\alpha $ images; its RI magnitudes and 2MASS near-IR colours are inconsistent with those expected for a PMS star.

In Young et al. (2005) IRAS F13052-7653 is reported as CHIIXR 60 which is a non-existent identifier. This must be a typo because IRAS F13052-7653 is identified with the ROSAT source CHIIXR 37 (Alcalá et al. 2000). The optical counterpart of IRAS F13052-7653 appears as a visual triple system (Larson et al. 1998) whose components are designated as IRAS F13052-7653N, IRAS F13052-7653NW and IRAS F13052-7653S in Table 7. None of the three components was recovered by our selection criteria, most likely because their colours, both in optical and and near-IR, may be contaminated. However, two of them, namely IRAS F13052-7653N and IRAS F13052-7653NW, fall just below the 10 Myr isochrone and their ( $\rm H\alpha _{12}-H\alpha _7$) indices suggest H$\alpha $ emission above 3$\sigma $ level (see Table 7); moreover, their near-IR colours are also consistent with those expected for PMS stars. IRAS F13052-7653S is not detected in 2MASS; its optical colours are inconsistent with membership to Cha II and its H$\alpha $ colour does not indicate significant emission.

IRAS F12571-7657

This object was detected in X-rays with ROSAT (Alcalá et al. 2000). Its optical and near IR colours are consistent with those of a PMS star with strong IR excess (see also Young et al. 2005). Previous authors (Allers et al. 2006; Vuong et al. 2001) report an extinction value significantly higher than the one determined by us (Table 6). Given the strong IR excess, our procedure for the temperature estimate provides unreliable results. Therefore, it may be a deeply embedded source or a highly reddened PMS star.

C 62, C 66

These sources were first detected by DENIS and proposed as young stellar object candidates of Cha II by Vuong et al. (2001) based on their IR colours and, more recently, by Young et al. (2005) and Allers et al. (2006) on the basis of IRAC@Spitzer and MIPS@Spitzer observations, respectively. From the effective temperature and $A_{\rm V}$ estimates provided in Table 6, and calculation of their luminosities as explained in Sect. 5, we estimate a mass of 0.08 and 0.06 $M_{\odot}$ for C 62 and C 66, respectively. The age we derive for both objects using the Baraffe et al. (1998) and Chabrier et al. (2000) tracks is around 2 Myr. These estimates of mass and age are fairly consistent with those derived by Allers et al. (2006). From their H$\alpha $ index we also notice strong line emission ($\sim$3$\sigma $ level) in both objects (see Table 6).


This source has been detected in the near and mid-IR by Persi et al. (2003), but it was first revealed as an X-ray emitting source, CHIIXR3, by Alcalá et al. (2000). The association of this object with IRAS 12551-7657 is still unclear. Based on this evidence and the effective temperature estimated by us (Table 6), ISO-CHA II 29 is proposed to be a PMS star in Cha II with a mass of about 1 $M_{\odot}$.


None of these sources, originally identified by Persi et al. (2003), was selected with MIPS@Spitzer by Young et al. (2005). These sources do not satisfy our criteria and their ( $\rm H\alpha _{12}-H\alpha _7$) index suggests no emission in H$\alpha $.

In our optical images, ISO-CHA II 98 appears as a visual binary, but the colours and magnitudes of its components are inconsistent with our selection criteria. We stress that ISO-CHA II 98 was wrongly identified in SIMBAD with the ROSAT X-ray source CHIIXR 24 by Alcalá et al. (2000). The optical spectrum of CHIIXR 24 shows the Li I 6708 Å  absorption line, but the H$\alpha $ line is well in absorption (Alcalá et al. 2000). Being classified as a G7-type star, CHIIXR 24 is most likely a field object unrelated to Cha II.

2MASS 13102531-7729085, 2MASS 13110329-7653330, 2MASS 13125238-7739182

These sources are reported in the list of candidates by Young et al. (2005). While 2MASS 13102531-7729085 and 2MASS 13110329-7653330 are rejected by our selection criteria, 2MASS 13125238-7739182 is selected as a PMS star candidate. The first two sources have H$\alpha $ indices indicating no H$\alpha $ emission, while the H$\alpha $ index of the last suggests H$\alpha $ in emission at 1$\sigma $ level. Thus, it is likely that 2MASS 13102531-7729085 and 2MASS 13110329-7653330 are field stars unrelated to Cha II.

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