3 NGC 2264

Stellar activity in NGC 2264 has been studied most recently by Flaccomio et al. (2000) through 6 different ROSAT HRI observations covering, in two different pointings, a large fraction of the star forming cloud. One hundred sixty nine distinct sources were detected, ${\sim} 95\%$ of which are estimated to be associated with members of the association. One of the main problem at the time of this work was the lack of good optical characterization of the members, so that, for example, lacking individual measurements, extinction toward sources had to be assumed uniform and placement of counterparts in the HR diagram was performed solely from optical photometric data, an error-prone procedure for PMS stars. Moreover the distinction between CTTS and WTTS was not well established, as indications on the NIR excesses were not available and H $_{\rm\alpha }$ measurements were in most cases qualitative and non-uniform. Since that work new improved optical data have been recently published by Rebull et al. (2002). We therefore updated the previous analysis according to the general principles stated in the introduction.

3.1 The reference sample

We adopted optical data from Tables 1 and 3 of Rebull et al. (2002). Out of the full list of 687 photometrically selected candidate members (i.e. the in cloud, in locus population defined by Rebull et al. 2002) we selected the 202 stars for which reliable spectral types and extinction (A_V) estimates (through R and I photometry + spectral types) were available. This latter is a subset of the full spectroscopic sample studied by Rebull et al. (2002), initially selected primarily from a list of I-band variable stars, with the addition of previously known candidate members based on their X-ray or H $_{\rm\alpha }$ emission or on their proper motion. Our reference sample is therefore the intersection (logical and) of the photometrically selected member sample and of the spectroscopic sample. While the former is arguably free from selection biases in favor of faint CTTS, the degree of representativeness of the latter in this respect is less clear: H $_{\rm\alpha }$ is however only a secondary selection criterion and I-band variability (periodic in >50% of the stars), although maybe more frequent in CTTS, does not obviously favor the inclusion of optically faint stars. Moreover the disk (CTTS) fractions Rebull et al. (2002) derive for the spectroscopic and the photometric sample (cf. their Table 6) are remarkably similar, suggesting that the former is not strongly biased toward CTTS.

Contamination of our reference sample from field stars may on the other hand be non negligible: according to preliminary proper motion data Rebull et al. (2002) report that ${\sim} 50\%$ of their photometric candidates are actually non-members. The spectroscopic sample, selected on the basis of PMS stellar characteristics, is expected to be less contaminated, although an estimate based on proper motion data is not provided. We recall (cf. Sect. 1) that field star contamination is expected to artificially lower the activity levels of WTTS.

We place stars in our reference sample in the HR diagram. Effective temperatures and bolometric corrections are estimated from spectral types and Kenyon & Hartmann (1995) conversions. We then evaluate bolometric luminosities from I band magnitudes. Out of the 202 spectrally characterized candidate members, 193 fall within the Siess et al. (2000) evolutionary model grid and have therefore been assigned a mass and an age.

The reference sample used for our following analysis comprises the 178 stars, out of these 193 candidate members characterized in terms of mass and age, that fall in the field of view of the X-ray observations described by Flaccomio et al. (2000).

3.2 X-ray data

We matched the photometric catalog of Rebull et al. (2002), out of which our reference member list is drawn, with the list of 169 X-ray sources published by Flaccomio et al. (2000). The identifications were carried out as described in Flaccomio et al. (2000), i.e. assuming as identification radii the the off-axis dependent X-ray source position error summed in quadrature to 1 $^{\prime\prime}$, i.e. a conservative estimate of the optical position error. Before performing the final identification we first registered the coordinate systems of the optical and X-ray lists by comparing the positions of 125 uniquely identified pairs (RA $_{\rm X}-{\rm RA}_{\rm opt}=-1.7^{\prime\prime}$, Dec $_{\rm X}-{\rm Dec}_{\rm opt}=1.9^{\prime\prime}$). Sixty seven stars in our reference sample were identified with an X-ray source, 56 of which uniquely, while the the remaining 11 fell in the identification circle of an X-ray source along with other objects in the photometric catalog. To each of the 67 candidate members with X-ray counterparts we then assigned a Maximum Likelihood (ML) X-ray count rate: these are values computed by Flaccomio et al. (2000) in order to define a mean source brightness among 6 different observations. Due to source variability, five of these mean values are actually upper limits. The count rates of the 11 ambiguous identifications were also treated as upper limits and upper limits, computed as in Flaccomio et al. (2000), were also assigned to 111 X-ray undetected candidate members lying within the FOV of the HRI observations.

Finally we converted count-rates, measured and upper limits, to X-ray luminosities in the 0.1-4.0 keV band. We assumed a thermal emission spectrum with kT=2.16 keV, close to recent estimates for PMS stars (e.g. Flaccomio et al. 2002b; Imanishi et al. 2001; Getman et al. 2002). The hydrogen column density was assumed proportional to the optical extinction measured individually for each star: $N_{\rm H}=2\times 10^{21}A_{V}$. The distance was assumed to be 760 pc like in Flaccomio et al. (2000). Figure 4 compares the $L_{\rm X}$ derived by Flaccomio et al. (2000) to those derived here from the same count-rates. Our new estimates are ${\sim} 0.3$ dex higher respect to the old ones, with the main differences due to the assumed kT (2.16 vs. 0.75 keV) and the increased absorption (mean $A_{V} \sim 0.45$ vs. a constant A_V=0.19) and a small difference, $\lesssim$0.1 dex, due to the different spectral band in which $L_{\rm X}$ is computed.

Figure 4: Comparison of X-ray luminosities computed by Flaccomio et al. (2000) for NGC 2264 stars (see text) and those recomputed from the same data in this work. No distinction is made here between detections and upper limits. The solid line indicates the locus of equal values; the dotted lines indicate the relation $L_{\rm X}$(This work) $=2 \cdot L_{\rm X}$ (Flaccomio et al. 2000). The scatter of points about the mean relation is due to the adoption of individual extinction corrections.

3.3 Activity vs. circumstellar environment

We will now try to establish whether the activity levels of our sample of NGC 2264 candidate members depend on circumstellar properties, evidenced either by accretion or by IR disk indicators. We will use two indicators provided by Rebull et al. (2002): H $_{\rm\alpha }$ equivalent width, $EW_{\rm H\alpha}$, and excess in the H-K near IR color, $\Delta(H-K)$. In the former case we will consider as accreting those stars with $EW_{\rm H\alpha} > 5$ and non-accreting those with $EW_{\rm H\alpha} < 5$. In the latter case we take as threshold $\Delta(H-K)=0.15$ with stars showing larger excesses considered as surrounded by circumstellar disks.

Figure 5 shows the scatter plots between $L_{\rm X}$ and stellar mass and $L_{\rm X}/L_{\rm bol}$ and mass, for the accreting and non-accreting stars. Similarly to what seen in other SFRs, $L_{\rm X}$ is seen to correlate with stellar mass, although the relation appears somewhat fuzzier in this case respect to, e.g., the ONC or the Chamaeleon I cases (see below). A systematic difference between the two classes is not readily apparent. However, given the large number of upper limits a more quantitative analysis is needed. Figure 6 shows XLFs and $L_{\rm X}/L_{\rm bol}$ distributions for the two H $_{\rm\alpha }$ separated stellar classes, in two different mass ranges, marked in Fig. 5 by vertical lines, and for the whole sample. In all cases the distributions of accreting stars appear to lie below those for non-accreting ones. Statistical test (in the ASURV package) confirm such differences with varying degree of confidence (results of the tests are given in the figure, along with the number of detection and upper limits that enter in the distributions; the same information is repeated in the first part of Table 1, where the results of other tests described below are also reported). Particularly significant is the difference in $L_{\rm X}$ for the whole sample ( ${>} 3.9\sigma$). Such a difference might however result from a larger fraction of low mass, and low $L_{\rm X}$, accreting stars (see Fig. 5). In the two narrower mass ranges, the result is however retrieved, with greater than $3\sigma$ confidence in the $0.5{-}1.0~M_\odot$ mass range. The difference is also observed in $L_{\rm X}/L_{\rm bol}$. Note that, as pointed out above, our reference sample may be significantly contaminated by field stars, and this would tend to depress the non-accreting stars distributions, thus lowering the significance of the result. If we repeat the same analysis including only stars confirmed as members by their IR excess ( $\Delta(I-K) > 0.3$ or $\Delta(H-K) > 0.15$), H $_{\rm\alpha }$ emission (EW > 5) and X-ray detection, we indeed find even more significant differences. Particularly interesting are the differences in $L_{\rm X}/L_{\rm bol}$, because they are less likely to be influenced by selection effects, a concern for this latter restricted sample. Table 1 reports the results of these tests.

Table 1 also reports the results of the comparisons between the stars with and without near IR excess, for the same mass ranges and the two stellar samples described above. The same results is retrieved: stars showing a $\Delta(H-K)$ excess, indicating the presence of a disk, have lower activity levels respect to the complementary sample.

 

 

Table 1: Significance of difference between CTTS and WTTS in NGC 2264.

Mass $[M_\odot]$	$N_{\rm d}(C)$	$N_{\rm u}(C)$	$N_{\rm d}(W)$	$N_{\rm u}(W)$	$Sign.(L_{\rm X})$	$Sign.(L_{\rm X}/L_{\rm bol})$
$EW(\rm H_{\alpha})$ - Whole sample
0.5-1.0	5	16	18	14	3.0/3.4	2.5/2.8
1.0-2.0	5	14	13	10	1.4/1.9	1.1/2.1
All	18	79	33	44	>3.9	1.5/2.1
$EW(\rm H_{\alpha})$ - Confirmed members
0.5-1.0	5	16	18	2	>3.9	3.2/>3.9
1.0-2.0	5	14	13	1	2.1/2.7	1.9/ 3.1
All	18	79	33	7	>3.9	3.3/>3.9
$\Delta(H-K)$ - Whole sample
0.5-1.0	1	9	18	20	1.3/1.8	2.2/3.2
1.0-2.0	1	10	14	11	1.9/2.2	2.1/2.5
All	4	30	40	71	2.2/2.4	2.6/3.2
$\Delta(H-K)$ - Confirmed members
0.5-1.0	1	9	18	9	1.8/2.3	2.6/ 3.5
1.0-2.0	1	10	14	2	2.4/2.8	2.9/ 3.4
All	4	30	40	44	2.8/ 3.0	3.4/>3.9

Note - Description of columns: (1) Mass range; (2) Number of detected CTTS; (3) Number of CTTS with upper limits; (4) Number of detected WTTS; (5) Number of WTTS with upper limits; (6) Range of significance (expressed in $\sigma$ equivalent) for the difference between the $L_{\rm X}$ distributions of CTTS and WTTS according to the tests in ASURV. Results greater than $3\sigma$ are in boldface; (7) like (6), but for $L_{\rm X}/L_{\rm bol}$.

Figure 5: $L_{\rm X}$ and $L_{\rm X}/L_{\rm bol}$ vs. mass for NGC 2264 stars with high and low H $_{\rm\alpha }$ equivalent width (black and and gray symbols, respectively). Filled circles represent detections, down-pointing arrows upper-limits.

Figure 6: Distributions of $L_{\rm X}$ and $L_{\rm X}/L_{\rm bol}$ (left and right columns) for NGC 2264 stars with high and low H $_{\rm\alpha }$ equivalent width (solid and dashed lines, respectively). Each row refers to a different mass range as indicated. Legends inside panels as in Fig. 2.