X-ray Observations of Eight Young Open Star Clusters: I. Membership and X-ray Luminosity

J. Astrophys. Astr. (2013) 34, 393–429 c© Indian Academy of Sciences

X-ray Observations of Eight Young Open Star Clusters:I. Membership and X-ray Luminosity

Himali Bhatt1, J. C. Pandey2, K. P. Singh3, Ram Sagar2

& Brijesh Kumar2

1Astrophysical Sciences Division, Bhabha Atomic Research Center,Trombay, Mumbai 400 085, India.2Aryabhatta Research Institute of Observational Sciences, Manora Peak,Nainital 263 129, India.3Tata Institute of Fundamental Research, Mumbai 400 005, India.e-mail: [email protected]

Received 17 June 2013; accepted 11 November 2013

Abstract. We present a detailed investigation of X-ray source contentsof eight young open clusters with ages between 4 to 46 Myr using archivalX-ray data from XMM-NEWTON. The probable cluster memberships ofthe X-ray sources have been established on the basis of multi-wavelengtharchival data, and samples of 152 pre-main sequence (PMS) low mass(<2M�), 36 intermediate mass (2–10M�) and 16 massive (>10M�) starshave been generated. X-ray spectral analyses of high mass stars reveal thepresence of high temperature plasma with temperature <2 keV, and meanLX/Lbol of 10−6.9. In the case of PMS low mass stars, the plasma tem-peratures have been found to be in the range of 0.2 keV to 3 keV with amedian value of ∼1.3 keV, with no significant difference in plasma tem-peratures during their evolution from 4 to 46 Myr. The X-ray luminositydistributions of the PMS low mass stars have been found to be similarin the young star clusters under study. This may suggest a nearly uni-form X-ray activity in the PMS low mass stars of ages ∼4–14 Myr. Theseobserved values of LX/Lbol are found to have a mean value of 10−3.6±0.4,which is below the X-ray saturation level. The LX/Lbol values for thePMS low mass stars are well correlated with their bolometric luminosi-ties, that implies its dependence on the internal structure of the low massstars. The difference between the X-ray luminosity distributions of theintermediate mass stars and the PMS low mass stars has not been foundto be statistically significant. Their LX/Lbol values, however have beenfound to be significantly different from each other with a confidence levelgreater than 99.999% and the strength of X-ray activity in the intermediatemass stars is found to be lower compared to the low mass stars. However,the possibility of X-ray emission from the intermediate mass stars due to

Supplementary material pertaining to this article is available on the Journal of Astrophysics &Astronomy website at http://www.ias.ac.in/jaa/dec2013/supp.pdf

393

http://www.ias.ac.in/jaa/dec2013/supp.pdf

394 Himali Bhatt et al.

a low mass star in close proximity of the intermediate mass star can notbe ruled out.

Key words. Open clusters and associations: NGC 663, NGC 869, NGC884, NGC 7380, Berkeley 86, IC 2602, Trumpler 18, Hogg15—stars: pre-main sequence—X-rays: massive stars, intermediate mass stars, low massstars.

1. Introduction

Young open star clusters constitute samples of stars of different masses with approxi-mately the same age, distance and chemical composition, and these are homogeneouswith respect to these properties. These clusters contain massive (>10M�), inter-mediate mass (10–2M�) and PMS low mass (<2M�) stars, and therefore, provideuseful laboratories to study different mechanisms for the generation of X-rays in starswith different masses. In the massive stars, the X-ray emission arises from shocksin radiatively-driven winds (Lucy & White 1980; Owocki & Cohen 1999; Kudritzki& Puls 2000; Crowther 2007), while in the low-mass stars, rotation with convectiveenvelopes drives a magnetic dynamo leading to strong X-ray emission (Vaiana et al.1981; Güdel 2004). Intermediate mass stars, on the other hand, are expected to be X-ray dark because (a) the wind is not strong enough to produce X-rays as in the case ofmassive stars (see Lucy & White 1980; Kudritzki & Puls 2000), and (b) being fullyradiative internal structure, the dynamo action cannot support the X-ray emission.However, the mysterious detection of X-rays from some intermediate mass stars stillremains an open question, and underlying physical mechanisms are not fully known(e.g., Stelzer et al. 2006).

Further, the physical origin of X-ray emission from PMS low mass stars is alsopoorly understood. X-ray studies of low mass PMS stars in young clusters with agesless than 5 Myrs like Orion, IC 348 and NGC 2264 (e.g., see Feigelson et al. 2003;Flaccomio et al. 2003a, b; Stassun et al. 2004; Preibisch et al. 2005), and in olderZero-Age-Main Sequence (ZAMS) clusters like the Pleiades and IC 2391 with agesbetween 30 and 100 Myr (e.g., Micela et al. 1999; Jeffries et al. 2006; Scholz et al.2007) offer strong evidence that X-ray activity of PMS low mass stars originatesdue to coronal activity similar to that present in our Sun. Studies of low mass stellarpopulation with different ages, however, show an evolution of the X-ray activitylevels in the young stages (<5 Myr), the X-ray luminosity (LX) is in the range of1029–1031 erg s−1, compared to much lower activity seen in the older (ZAMS) stars,i.e., LX ∼ 1029 erg s−1. The X-ray activity is found to decay mildly with age duringthe evolution of PMS low mass stars from 0.1 to 10 Myr (Preibisch & Feigelson2005), while it steepens in the Main Sequence (MS) evolution from the ZAMS to afew Gyr age (Feigelson et al. 2004). Thus, the evolution of X-ray activity in the PMSstars is somewhat more complicated than in the MS stars. In addition, the ratio of X-ray luminosities to the bolometric luminosity (LX/Lbol) of PMS low mass stars inyoung clusters is found to be above the saturation level, i.e., LX/Lbol ≈ 10−3, anduncorrelated with the rotation rates, while the low mass stars in ZAMS clusters showLX/Lbol ∼ 10−8–10−4. The stellar X-ray activity deviates from the saturation levelfor low mass stars in between 1 Myr to 100 Myr (Patten & Simon 1996; Güdel 2004;

XMM-Newton View of Eight Young Open Star Clusters 395

Currie et al. 2009). However, it is still not clear at which stage of the PMS evolution,the low mass stars deviate from the X-ray saturation level, and which fundamentalparameters govern their X-ray emission.

X-ray studies of clusters with intermediate age (5 to 30 Myr) have been few andfar between. An extensive study of young open clusters containing a number of starswith a range of masses from massive to PMS low mass stars can address issues spe-cific to the mechanisms producing X-rays in stars with different masses. In addition,the young open clusters with a wide range of ages are also very useful targets forexamining the evolution of X-ray emission with age, especially in low mass stars.Multi-wavelength surveys of young open clusters provide an effective way to iden-tify young cluster members among the huge number of foreground and backgroundstars (a few Gyr) present in the same sky region, as young stars are more luminousin X-rays compared to the older field stars (e.g., Micela et al. 1985, 1988, 1990;Caillault & Helfand 1985; Stern et al. 1981; Preibisch et al. 2005).

The present work deals with characterizing the X-ray source contents of eightyoung open clusters with ages ranging from 4 to 46 Myr. This data sample bridgesthe gap between young clusters like the Orion and the older clusters like the Pleiades,and constrain the evolution of X-ray emission with age for low mass stars. Samplesof massive, intermediate and low mass PMS stars were collected using multi-wavelength archival data. The values of the extinction (E(B–V )), distances and ageof the open clusters studied here are given in Table 1. The data were taken fromXMM-NEWTON pointed observations of the open clusters NGC 663, NGC 869,NGC 884 and IC 2602, whereas for the clusters NGC 7380, Berkeley 86, Hogg 15and Trumpler 18, data have been taken from serendipitous observations targeting themassive stars HD 215835, V444 Cyg, WR 47 and supernova remnant SNR MSH11-62, respectively. X-ray emission characteristics of these eight young open clustershave been investigated here for the first time. However, X-ray emission from a fewmassive stars in the open clusters NGC 7380, Berkeley 86, Hogg 15 and Trumpler 18have been reported earlier (for details, see §6.1). In addition, previous spectral stud-ies of the X-ray sources in the open cluster NGC 869 (h Persei) have been limitedto a region of size ∼15′ (diameter) with CHANDRA (Currie et al. 2009). The presentdata cover the entire NGC 869 cluster region (28′) due to the large field-of-view ofthe XMM-NEWTON. The paper is organized as follows: the details of X-ray obser-vations and data reduction procedure are presented in section 2. We have attempted

Table 1. The sample of the clusters under investigation with their basic parameters.

E(B–V ) NH∗ Distance Age

Cluster name (mag) (1020 cm−2) (pc) (Myr) References

NGC 663 0.80±0.15 40±7.5 2400±120 14±1 Pandey et al. (2005)NGC 869 0.55±0.10 28±5 2300±100 13.5±1.5 Currie et al. (2010)NGC 884 0.52±0.10 26±5 2300±100 14±1 Currie et al. (2010)NGC 7380 0.60±0.10 30±5 2600±400 4±1 Chen et al. (2011)Berkeley 86 0.95±0.10 47.5±5 1585±160 6±1 Bhavya et al. (2007)IC 2602 0.035±0.01 1.75±0.5 150±2 46±5 Dobbie et al. (2010)Hogg 15 1.15±0.1 57.5±5 3000±300 6±2 Sagar et al. (2001)Trumpler 18 0.3±0.04 15±0.2 1300±100 30±15 Delgado et al. (2007)

∗NH is derived using the relation NH = 5 × 1021 × E(B–V ) cm−2 from Vuong et al. (2003).


to ascertain the cluster probable membership of X-ray sources in section 3. X-rayvariability and spectra of the cluster members are presented in sections 4 and 5,respectively. The X-ray properties of cluster members are discussed in section 6 andresults are summarized in section 7.

2. X-ray observations and data reduction

XMM-NEWTON carries three co-aligned X-ray telescopes observing simultane-ously, and covering 30′ × 30′ region of the sky. It consists of three CCD-baseddetectors: the PN CCD (Strüder et al. 2001) and the twin CCD detectors MOS1and MOS2 (Turner et al. 2001). EPIC has moderate spectral resolution ( E

δE ∼ 20–50) and an angular resolution1 of 4.5′′, 6.0′′ and 6.6′′ for PN, MOS1 and MOS2detectors, respectively. It together constitutes the European Photon Imaging Camera(EPIC). We have analysed archival X-ray data from XMM-NEWTON observationsof eight young open clusters and the journal of observations is given in Table 2. Allthree EPIC detectors were active at the time of observations with full frame mode.Data reduction followed the standard procedures using the XMM-NEWTON Sci-ence Analysis System software (SAS version 10.0.0) with updated calibration files.Event files for MOS and PN detectors were generated by using tasks EMCHAIN andEPCHAIN, respectively, which allow calibration, both in energy and astrometry, of theevents registered in each CCD chip and combine them in a single data file. We lim-ited our analysis to the energy band to 0.3–7.5 keV because data below 0.3 keV aremostly unrelated to bona-fide X-rays, while above 7.5 keV only background countsare present, for the kind of sources that we are interested in. Event list files wereextracted using the SAS task EVSELECT. Data from the three cameras were individ-ually screened for high background periods and those time intervals were excludedwhere the total count rate (for single events of energy above 10 keV) in the instru-ments exceed 0.35 and 1.0 counts s−1 for the MOS and PN detectors, respectively.The useful exposure times, i.e., sum of good time intervals, obtained after screen-ing the high background periods for each cluster and corresponding to each detectorused, are given in Table 2.

2.1 Detection of X-ray point sources

Detection of point sources is based on the SAS detection task EDETECT_CHAIN,which is a chain script of various sub tasks (for details, see XMM documentation2).First, the input images were built in two energy ranges, a soft band (0.3–2.0 keV) anda hard band (2.0–7.5 keV) for all three EPIC detectors with a pixel size of 2.′′0, corre-sponding to a bin size of 40 pixels in the event file where each pixel size correspondsto 0.′′05. The task EDETECT_CHAIN was then used simultaneously on these images.This task determined the source parameters (e.g., coordinates, count rates, hard-ness ratios, etc.) by means of simultaneous maximum likelihood psf (point spread

1http://heasarc.gsfc.nasa.gov/docs/xmm/uhb/onaxisxraypsf.html2http://xmm.esac.esa.int/sas/current/documentation/threads/src_find_thread.shtml

http://heasarc.gsfc.nasa.gov/docs/xmm/uhb/onaxisxraypsf.html

http://xmm.esac.esa.int/sas/current/documentation/threads/src_find_thread.shtml


Tabl

e2.

Jour

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)(h

h:m

m:s

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ter

(arc

min

)

NG

C66

302

0116

0101

4191

514

Jan

2004

22:4

0:26

32.5

933

.13

28.6

9M

ediu

m1.

606

NG

C86

902

0116

0201

3950

919

Jan

2004

04:3

9:50

38.5

038

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34.8

8M

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m1.

871

NG

C88

402

0116

0301

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Feb

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833

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26.0

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013

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C73

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0565

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319

Dec

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ley

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0624

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127

Oct

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July

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function) fitting to the source count distribution in the soft and the hard energybands of each EPIC instrument. A combined maximum likelihood value in all threeinstruments was taken to be greater than 10, corresponding to a false detection prob-ability of ≈4.5 × 10−5. The output source lists from the individual EPIC camerasin different energy bands were merged into a common list and the average val-ues for the source positions with count rates were calculated. The final output listwas thus created giving source parameters for the soft and the hard energy bandsalong with the total energy band of 0.3–7.5 keV. Spurious detections due to inter-chip gaps between CCDs, the hot pixels and the surroundings of bright point sourceregions have been removed by visual screening. Finally, the number of X-ray sourcesdetected in NGC 663, NGC 869, NGC 884, NGC 7380, Berkeley 86, IC 2602, Hogg15 and Trumpler 18 were 85, 183, 147, 88, 95, 95, 124 and 208, respectively. Theestimated positions of the all X-ray point sources along with their count rates inthe total energy band of 0.3–7.5 keV are listed in Table 3. Each source has beenascribed a unique Identification Number (ID) which is also given in supplementarymaterial.

The count rate of an X-ray source detected using EDETECT_CHAIN task with 2σ

significance and lying within the cluster radius has been considered as the detectionlimit for each cluster. These detection limits in terms of count rates have been con-verted into flux limits using the Count Conversion Factors (CCFs) used for low massstars (see section 5) and corresponding X-ray luminosities have been tabulated inTable 4 for each cluster.

2.2 Infrared counterparts of X-ray sources

X-ray point sources detected in the clusters were cross-identified with NIR sourceslisted in the Two-Micron All Sky Survey (2MASS) Point Source Catalog (PSC; Cutriet al. 2003). The X-ray counterparts in the 2MASS catalogue were then searched forwithin a radius of 10′′, only those with a ‘read flag’ (representing uncertainties intheir magnitude) value of 1 or 2 were retained. In several cases, multiple counterpartsare possible in the 2MASS PSC corresponding to an X-ray source and the numberof multiple counterparts are given in column 9 (N ) in supplementary material. Insuch cases, NIR sources that are closest to an X-ray source have been adopted ascorresponding counterparts of that source. The J H KS magnitudes, the positions ofthe X-ray sources from the center of the corresponding open cluster (see section 3),and the offsets between the X-ray and the NIR positions of the NIR counterpartsare given in supplementary material. It was thus found that only 70%, 77%, 70%,86%, 94%, 78%, 85% and 93% of the X-ray sources in the open clusters NGC 663,NGC 869, NGC 884, NGC 7380, Berkeley 86, IC 2602, Hogg 15 and Trumpler18, respectively, have 2MASS NIR counterparts. Optical spectroscopic cataloguesof stars from Webda3 and Vizier4 were used for the optical identification of X-raysources (see supplementary material).

3http://www.univie.ac.at/webda/navigation.html4http://vizier.u-strasbg.fr/viz-bin/VizieR

http://www.univie.ac.at/webda/navigation.html

http://vizier.u-strasbg.fr/viz-bin/VizieR


Tabl

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).

http://www.ias.ac.in/jaa/dec2013/supp.pdf


Tabl

e4.

Det

ectio

nfr

actio

nof

X-r

ayso

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sw

ithin

clus

ter

radi

usw

itha

com

pari

son

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and

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ta.

Cen

ter

Det

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Rad

ius

No.

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ces

limits

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ster

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Det

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dId

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g(L

X)

nam

ehh

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hh:m

m:s

s(′ )

(′ )To

tal

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ster

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lust

erer

gs−

1

NG

C66

301

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29+6

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3N

GC

869

02:1

9:00

+57:

08:5

79

1418

318

114

178

30.2

4N

GC

884

02:2

2:04

+57:

08:4

99

1114

711

410

371

30.3

6N

GC

7380

22:4

7:47

+58:

07:1

510

888

4076

3730

.85

Ber

kele

y86

20:2

0:22

+38:

42:0

73

3.5

9611

9011

30.6

2IC

2602

10:4

3:06

−64:

25:3

050

–95

9574

7427

.57

Hog

g15

12:4

3:39

−63:

05:5

83.

57

124

5310

647

30.7

0T

rum

pler

1811

:11:

31−6

0:40

:41

2.5

–20

811

194

1029

.75

1R

adiu

sde

rive

dus

ing

optic

alda

tain

liter

atur

e(D

ias

etal

.200

2).

22σ

dete

ctio

nlim

itsof

obse

rvat

ions

are

deri

ved

from

coun

trat

eco

nver

sion

into

flux

inPN

dete

ctor

for

low

mas

sst

ars

(see

§4).


3. Cluster membership of X-ray sources

The X-ray sources with identifiable counterparts in the NIR band may not necessarilybe members of their respective clusters. It is difficult to decide cluster membershipof an individual X-ray source, because the cluster population is contaminated byforeground and background stellar sources (Pizzolato et al. 2000) and extragalacticsources (Brandt & Hasinger 2005). In order to find which of the X-ray sources actu-ally belong to a cluster, the approach given by Currie et al. (2010) has been adoptedhere. The step by step procedure used is given below.

3.1 Center and radius of the clusters

The stellar population associated with young open clusters is still embedded in parentmolecular clouds, due to which a large variation in extinction is found within youngopen clusters. The young stars embedded within high extinction regions of the clusterand hidden in the optical bands may be visible in the NIR band. Therefore, NIR datafrom the 2MASS PSC (Cutri et al. 2003) were used to estimate the center of theseclusters and their extents rather than the optical data. The center of a cluster was firsttaken to be an eye estimated center of the cluster and then refined as follows. Theaverage RA(J2000) and DEC(J2000) position of 2MASS stars having KS ≤ 14.3 mag(99% completeness limit in KS band) and lying within 1′ radius was computed. Theaverage RA(J2000) and DEC(J2000) were reestimated by using this estimated valueof the center of the cluster. This iterative method was used until it converged to a con-stant value for the center of the open cluster (see Joshi et al. 2008 for details). Typicalerror expected in locating the center by this method is ∼5′′. The estimated values ofthe center of each open cluster are given in Table 4. The positions of the centers esti-mated from the NIR data are consistent within 1′ to that estimated from the opticaldata by Dias et al. (2002). Assuming spherical symmetry for the cluster, a projectedradial stellar density profile of stars was constructed and the radius at which the stel-lar density is at the 3σ level above the field star density was determined. The fieldstar densities were estimated from the 200 arcmin2 region which is nearly more than0.5 degree away from the cluster regions. The estimated values of field star densitiesare 2.21 ± 0.10, 1.95 ± 0.10, 1.95 ± 0.10, 2.86 ± 0.12, 6.91 ± 0.18 and 13.14 ± 0.26stars arcmin−2 for the clusters NGC 663, NGC 869, NGC 884, NGC 7380, Berkeley86 and Hogg 15, respectively. The estimated values of the radii of the open clustersare given in Table 4. We have adopted the estimated radii, reported in Table 4, as ameasure of the extent of the open clusters. Our estimates of the radii of the clustersare larger than that of the values given by Dias et al. (2002) using optical data exceptfor the open cluster Berkeley 86. However, these values are consistent with the val-ues given by Pandey et al. (2005) and Currie et al. (2010) in the case of NGC 663 andNGC 869, respectively. It is not possible to define the cluster extent in the case of IC2602 and Trumpler 18 because the boundary where the stellar densities merge intofield star densities is not clearly marked in the radial density profiles. This may beeither due to the very large size of the cluster in the case of IC 2602, and very smallsize in the case of Trumpler 18, or the stars in the clusters may not be distributed in aspherical symmetry. For further analysis, we used the radii given in Dias et al. (2002)catalogue for these two open clusters. The projected distances of X-ray sourcesfrom the center of the respective clusters are given in supplementary material. The


number of X-ray sources within the radius of cluster are also given in Table 4. Allthe X-ray sources with a counterpart in NIR and falling within the adopted radius ofthe corresponding cluster have been considered for further analysis to check if theyare members of that cluster.

3.2 Color-magnitude diagram of X-ray sources with NIR counterparts

Assigning cluster membership to X-ray sources is a difficult task. It is, however, eas-ier to check if an X-ray source lying within the cluster radius is not a member byusing Color Magnitude Diagrams (CMDs). In Fig. 1, we plot the CMDs using the2MASS J magnitudes and (J − H) colors of the sources selected in section 3.1 andlying within the cluster radii. We define the fiducial locus of cluster members foreach open cluster by the post-main sequence isochrones from Girardi et al. (2002)and PMS isochrones from Siess et al. (2000) according to their ages, distances andmean reddening (see Table 1). These locii have been shown by dashed lines in Fig. 1.Width of each of the cluster locii has been determined by (1) uncertainties in thedetermination of distance and age of the cluster, (2) uncertainties in the photomet-ric 2MASS J and H magnitudes of the sources which is higher at the fainter end,(3) dispersion in the reddening, and (4) binarity. Equal mass binaries may be up to0.75 mag more luminous than single stars. The X-ray sources which are lying out-side this fiducial locus for a given cluster are excluded from being members in thatcluster. The number of stars thus excluded from being members are 10, 22, 11, 9,2, 42, 14 and 2, in the open clusters NGC 663, NGC 869, NGC 884, NGC 7380,Berkeley 86, IC 2602, Hogg 15 and Trumpler 18, respectively.

The remaining X-ray sources were further screened for probable membership ofthe respective open clusters. Each of these X-ray source was investigated on thebasis of information given in the optical spectroscopic catalogues from Vizier ser-vices. This spectroscopic information with references is given in supplementarymaterial. The sources for which the spectroscopic characteristics did not match withtheir photometric location in the CMDs were no longer considered for membershipof the corresponding cluster, and thus removed from the list of probable mem-bers. This method is useful for removing the foreground contamination. However,background contamination is very difficult to separate. Therefore, we have furthercross-identified these selected sources in the all-sky comprehensive catalogue ofradio and X-ray associations by Flesch (2010), in which the probability of a sourcebeing a quasi stellar object (QSO), a galaxy or a star has been given. Those sourcesfor which the probability for being a star is less than 20% were also removed from thelist of selected sources. Using this method, the number of additional X-ray sourcesthat are no longer considered as members of the open clusters NGC 663, NGC 869,NGC 884, NGC 7380, Berkeley 86, IC 2602 and Hogg 15 were 5, 1, 3, 2 , 1, 2 and1, respectively.

The X-ray sources that are no longer considered as members of a cluster aremarked by the symbol of a cross in Fig. 1 and listed as ‘N’ in supplementary material(column 14). The remaining X-ray sources are considered to be the probable mem-bers of their respective clusters. We could thus assign probable cluster membershipfor 21, 70, 34, 25, 8, 10, 30 and 6 X-ray sources in clusters NGC 663, NGC 869,NGC 884, NGC 7380, Berkeley 86, IC 2602, Hogg 15 and Trumpler 18, respectively,


Fig

ure

1.J

vers

us(J

−H

)co

lor

mag

nitu

dedi

agra

m(C

MD

)of

the

X-r

ayso

urce

sw

ithin

clus

ter

radi

us.P

ost-

mai

nse

quen

ceis

ochr

ones

from

Gir

ardi

etal

.(20

02)

and

PMS

isoc

hron

esfr

omSi

ess

etal

.(20

00)

are

show

nby

solid

and

dotte

dlin

es(i

nbl

ue).

The

wid

thof

the

clus

ter

locu

sdu

eto

unce

rtai

ntie

sin

dete

rmin

atio

nof

dist

ance

,age

,mag

nitu

des

and

bina

rity

issh

own

byda

shed

lines

(in

red)

.The

sym

bols

ofst

ar,t

rian

gle

and

dots

repr

esen

tm

assi

ve,

inte

rmed

iate

and

low

mas

sst

ars,

resp

ectiv

ely,

and

thes

ebo

unda

ries

are

mar

ked

byar

row

s.T

heX

-ray

sour

ces

iden

tified

asa

non-

mem

ber

ofth

eir

resp

ectiv

eop

encl

uste

rsar

em

arke

dby

the

sym

bolo

fa

cros

s.


and listed as ‘Y’ in supplementary material (column 14). Further, proper-motionand/or spectroscopic studies are needed to confirm the membership of specific X-ray sources. However, proper motions at a distance of 2.0 kpc are extremely hard todetect. X-ray sources lying outside the cluster radius remain as unclassified.

3.3 Mass estimation of X-ray stars

The masses of X-ray stars, identified as probable members of the clusters, were esti-mated using theoretical isochrones of Girardi et al. (2002) for MS stars and Siesset al. (2000) for PMS stars. The boundaries corresponding to a 10M� (massive) star,10–2M� (intermediate mass) star and a 2M� (low mass) star were derived from theJ magnitudes using model isochrones corrected for distance, age and reddening foreach cluster, and are shown in Fig. 1 by arrows. The estimated mass of each staridentified as a probable member is given in supplementary material. In Fig. 1, themassive stars, intermediate mass stars and low mass stars are marked by the symbolsof star, triangle and dots, respectively.

4. Variability of X-ray sources

X-ray emission of stars is known to be variable, and before estimating their luminos-ity function it is important to first study their variability. Due to its highest sensitivity,data from the PN detector of EPIC were used for variability and spectral analysis.Light curves and spectra for all the probable members were extracted using circu-lar extraction regions centered on the source position provided by EDETECT_CHAIN

task in the energy range of 0.3–7.5 keV. X-ray sources either falling in the inter-chipgaps in the PN detector or having total counts below 40 in the PN detector wereignored for variability and spectral analyses. The wings of the psf for bright sourcesare often largely contaminated by emission from neighboring sources, therefore, theradii of extraction regions were varied between 8′′ and 40′′ depending on the positionof the source in the detector and its angular separation with respect to the neighboringX-ray sources. The background data were taken from several neighboring source-free regions on the detectors. For the timing analysis, we have binned the data with300–5000 s according to the count rate of the sources. Due to poor count statistics,there were several time intervals in which count rates were lesser than 5, therefore,we were not able to perform the χ2-test for variability analysis. Fractional root meansquare (rms) variability amplitude (Fvar) was estimated to quantify the variability inthe X-ray light curves for intermediate and low mass stars. The Fvar and the errorin Fvar (σFvar ) have been defined as follows (Edelson et al. 1990, 2002) and given inTable 5:

Fvar = 1

〈X〉√

S2 − 〈σ 2err〉, (1)

σFvar = 1

Fvar

√1

2N

S2

〈X〉2, (2)

where S2 is the total variance of the light curve, 〈σ 2err〉 is the mean error squared and

〈X〉 is the mean count rate, however, Fvar can not be defined when S2 is lesser than〈σ 2

err〉. Fvar quantifies the amplitude of variability with respect to the mean count rate.


Tabl

e5.

Spec

tral

and

timin

gpr

oper

ties

ofin

term

edia

tean

dlo

wm

ass

star

sin

youn

gop

encl

uste

rs.

Sour

cede

tect

ion

Spec

tral

Tim

ing

Clu

ster

RA

J200

0D

EC

J200

0kT

log(

EM

)lo

g(L

X)†

Bin

size

nam

eID

(deg

)(d

eg)

(keV

)(c

m−3

)(e

rgs

s−1)

log(

LX

Lbo

l)

(s)

Fva

r

Inte

rmed

iate

mas

sst

ars

(2–1

0M

�)N

GC

663

4226

.583

460

61.2

6397

3>

2.78

53.8

1+0.1

4−0

.18

31.0

2−6

.47

NG

C66

348

26.6

0425

061

.208

221

>1.

7953

.85+0

.15

−0.1

930

.98

−6.0

1N

GC

869

1034

.441

750

57.0

8813

930

.70

−5.7

5N

GC

869

1334

.457

584

57.3

1852

731

.26

−4.7

5N

GC

869

1734

.481

415

57.2

1849

830

.98

−6.0

8N

GC

869

2134

.512

001

57.0

7688

930

.51

−4.4

3N

GC

869

6434

.647

793

57.2

1441

730

.81

−4.9

1N

GC

869

6734

.662

083

57.2

2044

40.

54+0

.43

−0.3

353

.42+0

.56

−0.3

230

.48

−4.3

580

00.

89±0

.17

NG

C86

969

34.6

7141

757

.102

196

2.75

+14.

10−1

.45

54.0

3+0.1

6−0

.21

31.1

0−3

.50

NG

C86

988

34.7

5091

657

.154

499

30.5

1−6

.47

NG

C86

910

534

.783

291

57.1

2780

43.

10+2

.50

−1.1

054

.010.

10 −0.1

031

.11

−4.5

5

NG

C86

913

934

.869

793

57.1

5402

61.

95+1

.96

−0.6

254

.06+0

.09

−0.1

131

.09

−6.1

6

NG

C86

914

034

.870

708

57.1

6388

71.

27+0

.43

−0.2

353

.98+0

.13

−0.1

531

.04

−5.4

280

00.

72±0

.14

NG

C86

916

935

.026

749

57.1

2550

030

.70

−6.5

3N

GC

869

173

35.0

3783

457

.020

638

2.80

+2.3

0−0

.80

54.3

40.08 −0.1

031

.52

−5.0

4

NG

C88

423

35.3

2816

757

.146

667

>3.

8054

.12+0

.20

−0.0

931

.33

−4.8

411

000.

13±0

.47

NG

C88

462

35.4

6191

857

.026

669

>3.

2653

.90+0

.22

−0.1

731

.10

−5.0

8N

GC

884

6835

.489

666

57.2

1161

3<

30.5

2<

−6.2

2N

GC

884

7835

.517

044

57.1

4225

030

.81

−5.1

2N

GC

884

103

35.6

1320

957

.054

779

1.30

+0.3

1−0

.23

54.0

0+0.1

1−0

.14

31.0

4−6

.29

NG

C88

412

035

.707

790

57.1

4283

4>

15.8

054

.06+0

.13

−0.2

431

.16

−4.1

8N

GC

7380

5134

1.83

798

58.0

8866

530

.99

−5.4

4N

GC

7380

5934

1.89

688

58.1

2550

02.

38+1

.43

−0.6

554

.59+0

.08

−0.0

931

.65

−4.4

5


Tabl

e5.

(Con

tinu

ed).

Sour

cede

tect

ion

Spec

tral

Tim

ing

Clu

ster

RA

J200

0D

EC

J200

0kT

log(

EM

)lo

g(L

X)†

Bin

size

nam

eID

(deg

)(d

eg)

(keV

)(c

m−3

)(e

rgs

s−1)

log(

LX

Lbo

l)

(s)

Fva

r

Ber

kele

y86

7630

5.04

416

38.6

9513

71.

36+0

.70

−0.3

153

.74+0

.17

−0.2

030

.47

−5.8

0B

erke

ley

8679

305.

0075

538

.680

832

Not

fitte

dw

ithm

odel

Hog

g15

1119

0.79

124

−63.

0030

560.

14+0

.62

−0.0

5>

56.9

532

.57

−4.4

4

Hog

g15

3319

0.89

438

−63.

0147

780.

28+0

.66

−0.1

0>

55.4

531

.73

−4.8

6

Hog

g15

3419

0.98

633

−63.

0814

441.

55+0

.27

−0.2

454

.80+0

.06

−0.0

631

.76

−5.4

6

Hog

g15

4219

0.96

179

−63.

1133

880.

16+0

.12

−0.0

5>

56.1

732

.13

−4.0

9

Hog

g15

4819

0.98

618

−63.

0716

090.

18+0

.05

−0.0

355

.44+0

.38

−0.3

632

.19

−4.1

5H

ogg1

555

191.

0091

2−6

3.06

7471

31.7

1−5

.01

Hog

g15

6519

1.06

125

−63.

1377

490.

69+0

.26

−0.2

254

.20+0

.16

−0.1

731

.30

−4.8

548

000.

69±0

.33

Hog

g15

8219

1.15

871

−63.

0861

93>

6.27

54.3

9+0.1

2−0

.19

31.5

6−4

.87

Tru

mpl

er18

7616

7.88

104

−60.

6646

126.

99+5

0.8

−3.6

153

.53+0

.08

−0.0

930

.70

−4.5

0

Tru

mpl

er18

8016

7.89

404

−60.

6672

785.

63+1

3.8

−2.8

254

.21+0

.08

−0.0

831

.38

−4.7

7

Tru

mpl

er18

9016

7.92

091

−60.

7062

760.

84+0

.12

−0.0

853

.71+0

.05

−0.0

630

.81

−5.7

5

Low

mas

sst

ars

(<2

M�)

NG

C66

31

26.1

7708

461

.141

083

31.2

3−4

.17

NG

C66

34

26.2

8458

261

.317

722

30.9

8−4

.02

NG

C66

37

26.3

1950

061

.133

999

30.7

9−4

.28

NG

C66

38

26.3

2074

961

.281

723

31.3

3−4

.15

NG

C66

312

26.3

5829

261

.181

416

30.9

8−4

.58

NG

C66

315

26.3

7854

261

.141

777

31.3

6−3

.57

NG

C66

317

26.3

8579

261

.345

470

30.9

6−3

.68

NG

C66

318

26.3

9062

561

.064

499

30.9

5−3

.59

NG

C66

341

26.5

8250

061

.229

389

30.9

3−4

.37

NG

C66

352

26.6

3970

861

.056

168

30.8

5−3

.41

NG

C66

356

26.6

6383

461

.150

028

0.27

+0.0

2−0

.02

55.2

9+0.0

9−0

.07

32.1

6(?)

−2.0

1(?)

500

0.10

±0.2

3


Tabl

e5.

(Con

tinu

ed).

Sour

cede

tect

ion

Spec

tral

Tim

ing

Clu

ster

RA

J200

0D

EC

J200

0kT

log(

EM

)lo

g(L

X)†

Bin

size

nam

eID

(deg

)(d

eg)

(keV

)(c

m−3

)(e

rgs

s−1)

log(

LX

Lbo

l)

(s)

Fva

r

NG

C66

361

26.6

8904

161

.107

445

>9.

6654

.35+0

.08

−0.1

831

.46(

?)−1

.84(

?)N

GC

663

6626

.729

250

61.0

4897

331

.14

−3.4

2N

GC

663

8026

.941

376

61.1

8349

8>

9.37

54.6

4+0.1

1−0

.12

31.7

8−3

.69

2000

0.26

±0.1

8N

GC

663

8126

.983

583

61.2

0941

531

.57

−3.2

4N

GC

663

8227

.015

417

61.2

1274

91.

36+0

.79

−0.3

854

.15+0

.16

−0.1

931

.17

−4.1

7

NG

C66

383

27.0

1970

961

.118

973

0.50

+0.2

9−0

.31

54.7

4+0.7

9−0

.26

31.7

6−3

.52

3500

0.54

±0.4

0N

GC

869

134

.341

042

57.1

7475

131

.16

−2.9

6N

GC

869

1934

.498

043

56.9

6566

831

.01

−3.8

5N

GC

869

2034

.511

124

57.1

1072

230

.54

−3.5

1N

GC

869

2234

.518

623

57.3

4547

030

.91

−3.9

2N

GC

869

2334

.522

251

57.0

2436

130

.65

−3.4

0N

GC

869

2534

.530

666

57.2

0755

430

.57

−3.9

6N

GC

869

3334

.557

877

57.2

9772

230

.65

−4.0

3N

GC

869

3434

.559

208

57.0

3061

330

.78

−4.1

8N

GC

869

3634

.571

877

57.1

2541

60.

97+0

.65

−0.1

953

.85+0

.12

−0.1

430

.95

−3.1

6N

GC

869

4034

.584

126

57.2

7314

030

.60

−3.3

6N

GC

869

4134

.585

751

57.1

0994

331

.03

−3.4

6N

GC

869

4234

.586

708

57.1

7411

03.

66+1

4.00

−1.9

054

.25+0

.11

−0.1

231

.38

−3.1

860

00.

47±0

.10

NG

C86

944

34.5

9516

557

.182

278

2.47

+3.8

2−0

.85

53.8

6+0.1

2−0

.14

30.9

3−3

.36

NG

C86

947

34.6

0362

657

.112

083

30.4

7−3

.87

NG

C86

958

34.6

2891

857

.064

999

>2.

2153

.65+0

.28

−0.2

430

.86

−3.0

5N

GC

869

6134

.633

835

57.2

4839

030

.65

−4.5

1N

GC

869

6834

.668

709

57.1

2722

430

.67

−3.5

9N

GC

869

7334

.696

793

56.9

5694

431

.09

−3.0

8N

GC

869

7434

.701

958

57.0

7444

41.

25+0

.77

−0.2

753

.75+0

.15

−0.1

630

.80

−4.5

330

000.

20±0

.55

NG

C86

975

34.7

0308

357

.121

082

2.60

+1.9

0−0

.75

54.1

1+0.0

8−0

.09

31.1

9−3

.85


Tabl

e5.

(Con

tinu

ed).

Sour

cede

tect

ion

Spec

tral

Tim

ing

Clu

ster

RA

J200

0D

EC

J200

0kT

log(

EM

)lo

g(L

X)†

Bin

size

nam

eID

(deg

)(d

eg)

(keV

)(c

m−3

)(e

rgs

s−1)

log(

LX

Lbo

l)

(s)

Fva

r

NG

C86

978

34.7

1787

657

.192

001

3.30

+7.4

5−1

.38

53.8

4+0.1

3−0

.15

29.6

9−4

.46

NG

C86

979

34.7

1916

657

.074

360

30.4

4−3

.86

NG

C86

992

34.7

5483

357

.003

334

>2.

2854

.04+0

.12

−0.1

331

.24

−3.8

2N

GC

869

9334

.754

959

57.0

9769

430

.85

−3.8

3N

GC

869

9634

.759

666

57.1

6072

13.

00+3

.94

−1.3

054

.07+0

.10

−0.1

231

.17

−3.8

220

000.

31±0

.27

NG

C86

999

34.7

6812

457

.160

999

2.42

+2.2

7−0

.90

54.0

7+0.0

9−0

.10

31.1

3−3

.27

1000

0.45

±0.2

1

NG

C86

910

434

.782

333

57.1

0263

8>

1.89

53.8

3+0.1

3−0

.16

30.9

8−3

.23

NG

C86

910

834

.790

833

57.1

5300

030

.35

−4.1

8N

GC

869

111

34.7

9462

457

.125

721

600

0.64

±0.0

9N

GC

869

113

34.8

0196

057

.064

667

0.23

+0.0

5−0

.04

54.4

0+0.2

1−0

.17

31.2

5−3

.67

NG

C86

911

434

.807

625

57.0

3019

31.

71+1

.25

−0.5

2053

.79+0

.14

−0.2

030

.80

−3.2

4N

GC

869

115

34.8

1229

057

.218

582

30.8

0−3

.22

NG

C86

912

134

.824

207

57.1

3497

230

.35

−4.0

3N

GC

869

123

34.8

3091

757

.202

499

30.6

8−3

.93

NG

C86

912

834

.839

500

57.0

3480

530

.90

−3.3

4N

GC

869

129

34.8

4166

757

.232

334

30.5

7−3

.30

NG

C86

913

234

.848

293

57.0

2541

730

.72

−3.7

9N

GC

869

135

34.8

5487

456

.991

890

0.99

+0.7

3−0

.64

54.0

7+0.1

5−0

.17

31.1

7−3

.60

NG

C86

914

334

.883

999

57.1

1399

830

.54

−4.3

2N

GC

869

145

34.8

9058

357

.059

082

0.73

+0.2

8−0

.27

53.6

0+0.1

5−0

.19

30.7

2−3

.24

NG

C86

914

634

.890

751

57.2

5816

730

.51

−4.1

0N

GC

869

149

34.9

1346

57.0

6080

630

.51

−4.2

2N

GC

869

152

34.9

2924

957

.312

527

30.8

3−3

.16

NG

C86

915

534

.947

918

57.1

5119

631

.20

−3.3

8N

GC

869

156

34.9

5266

757

.039

165

4.40

+12.

0−1

.90

54.1

2+0.1

0−0

.12

31.2

8−3

.21

NG

C86

915

834

.960

876

57.2

5127

830

.81

−3.6

4


Tabl

e5.

(Con

tinu

ed).

Sour

cede

tect

ion

Spec

tral

Tim

ing

Clu

ster

RA

J200

0D

EC

J200

0kT

log(

EM

)lo

g(L

X)†

Bin

size

nam

eID

(deg

)(d

eg)

(keV

)(c

m−3

)(e

rgs

s−1)

log(

LX

Lbo

l)

(s)

Fva

r

NG

C86

915

934

.961

082

57.1

9933

32.

14+2

.54

−0.6

954

.09+0

.10

−0.1

231

.12

−3.3

1

NG

C86

916

234

.980

000

57.2

2044

40.

93+0

.30

−0.4

753

.80+0

.14

−0.1

730

.96

−3.8

5N

GC

869

165

34.9

9945

857

.253

887

30.7

4−4

.01

NG

C86

916

635

.011

665

57.3

1863

831

.15

−2.5

2N

GC

869

174

35.0

7104

157

.147

530

30.9

3−2

.95

NG

C86

917

535

.083

668

57.0

1047

1<

30.6

5<

−3.7

8N

GC

884

1635

.268

585

57.1

2913

930

.56

−4.2

6N

GC

884

2835

.357

166

57.0

5022

030

.59

−4.1

6N

GC

884

3035

.369

831

57.1

0380

61.

06+0

.42

−0.3

553

.77+0

.44

−0.3

530

.75

−3.7

1N

GC

884

3935

.409

168

57.0

6830

630

.47

−3.7

6N

GC

884

4935

.434

875

57.2

6447

330

.74

−3.5

3N

GC

884

5335

.442

039

57.0

4449

8>

0.70

<54

.45

30.9

4−3

.37

NG

C88

464

35.4

6654

157

.276

859

2.40

+5.2

0−0

.94

54.0

0+0.1

4−0

.17

31.0

7−2

.84

NG

C88

471

35.5

0395

657

.079

613

>4.

2654

.06+0

.17

−0.0

931

.27

−3.6

7N

GC

884

7535

.512

707

57.3

0852

931

.09

−3.1

3N

GC

884

8035

.524

750

57.1

1869

430

.50

−4.3

9N

GC

884

8835

.552

002

57.1

0075

01.

65+0

.68

−0.3

154

.15+0

.09

−0.1

031

.17

−2.8

122

000.

43±

0.17

NG

C88

489

35.5

5316

557

.176

613

30.9

8−3

.29

NG

C88

490

35.5

6316

857

.140

194

30.7

9−3

.21

NG

C88

491

35.5

6716

957

.094

028

1.83

+1.5

6−0

.47

54.0

1+0.1

0−0

.11

31.0

4−3

.04

2500

0.12

±0.

62N

GC

884

101

35.6

0546

157

.012

085

30.9

2−3

.20

NG

C88

410

235

.611

462

57.0

8974

830

.64

−3.8

7N

GC

884

104

35.6

1454

457

.025

196

30.6

6−4

.33

NG

C88

410

635

.626

041

57.2

4263

830

.59

−3.4

1N

GC

884

110

35.6

4337

557

.191

418

30.4

7−4

.65

NG

C88

411

335

.655

708

57.2

1794

530

.61

−3.3

5N

GC

884

119

35.7

0754

257

.161

026

31.1

1−3

.42


Tabl

e5.

(Con

tinu

ed).

Sour

cede

tect

ion

Spec

tral

Tim

ing

Clu

ster

RA

J200

0D

EC

J200

0kT

log(

EM

)lo

g(L

X)†

Bin

size

nam

eID

(deg

)(d

eg)

(keV

)(c

m−3

)(e

rgs

s−1)

log(

LX

Lbo

l)

(s)

Fva

r

NG

C88

412

135

.722

958

57.0

9880

4<

30.5

3<

−4.4

1N

GC

884

122

35.7

2420

957

.126

083

1.02

+0.2

7−0

.21

53.8

8+0.1

2−0

.14

30.9

6−3

.78

NG

C88

413

035

.747

002

57.1

3741

730

.44

−4.0

8N

GC

884

137

35.8

0370

757

.122

444

30.7

4−3

.42

NG

C88

413

835

.820

457

57.2

2250

031

.27

−3.2

9N

GC

7380

3434

1.70

270

58.1

3319

430

.86

−4.5

9N

GC

7380

3534

1.72

348

58.1

6808

3<

31.2

8<

−2.8

7N

GC

7380

3734

1.72

653

58.1

3472

430

.88

−3.4

1N

GC

7380

4134

1.74

469

58.0

7669

430

.92

−3.1

6N

GC

7380

4634

1.77

658

58.1

1388

831

.09

−2.8

1N

GC

7380

4834

1.80

121

58.1

0941

71.

78+2

.35

−0.6

054

.06+0

.16

−0.2

231

.09

−3.3

0N

GC

7380

4934

1.80

286

58.1

9133

431

.06

−3.2

9N

GC

7380

5034

1.81

659

58.0

6789

01.

37+0

.51

−0.2

254

.18+0

.13

−0.1

531

.20

−3.2

0N

GC

7380

5734

1.89

398

58.1

3805

431

.25

−3.2

7N

GC

7380

6034

1.89

981

58.0

4986

231

.08

−2.7

2N

GC

7380

6234

1.90

005

58.0

7522

22.

72+3

.98

−0.9

254

.36+0

.12

−0.1

431

.45

−3.1

2N

GC

7380

6534

1.91

592

58.0

9999

830

.96

−3.0

0N

GC

7380

6634

1.91

595

58.0

2347

2>

4.68

54.4

8+0.1

8−0

.14

31.6

9−2

.75

1200

0.15

±0.5

3

NG

C73

8067

341.

9224

558

.148

861

1.64

+1.2

5−0

.46

54.4

4+0.1

2−0

.14

31.4

8(?)

−2.3

4(?)

NG

C73

8070

341.

9397

058

.058

472

31.1

3−2

.99

NG

C73

8071

341.

9429

658

.184

776

31.1

2(?)

−2.4

9(?)

NG

C73

8080

341.

9797

758

.056

499

31.1

4−2

.69

NG

C73

8082

342.

0018

358

.044

971

31.0

3−3

.20

NG

C73

8083

342.

0178

558

.068

554

31.2

2−2

.85

NG

C73

8085

342.

0281

758

.075

195

6.50

+4.3

1−2

.22

55.0

3+0.0

5−0

.05

32.2

2−2

.06

NG

C73

8086

342.

1295

558

.147

583

31.4

3−4

.00


Tabl

e5.

(Con

tinu

ed).

Sour

cede

tect

ion

Spec

tral

Tim

ing

Clu

ster

RA

J200

0D

EC

J200

0kT

log(

EM

)lo

g(L

X)†

Bin

size

nam

eID

(deg

)(d

eg)

(keV

)(c

m−3

)(e

rgs

s−1)

log(

LX

Lbo

l)

(s)

Fva

r

Ber

kele

y86

7530

5.04

169

38.7

2100

11.

87+1

.63

−0.6

354

.15+0

.10

−0.1

231

.18

−4.2

425

000.

22±0

.34

Ber

kele

y86

8130

5.08

072

38.7

0072

230

.92(

?)−2

.27(

?)B

erke

ley

8684

305.

0889

938

.687

695

<30

.78

<−4

.51

Ber

kele

y86

8930

5.10

193

38.7

0063

831

.21

−4.5

0B

erke

ley

8693

305.

1244

238

.665

749

30.9

8−2

.75

IC26

026

160.

2498

8−6

4.33

4000

1.31

+0.3

1−0

.09

52.2

1+0.0

5−0

.06

29.2

4−3

.43

400

1.02

±0.0

9

IC26

0216

160.

3618

8−6

4.33

9302

2.33

+4.8

0−0

.74

51.3

4+0.1

3−0

.16

27.8

0−3

.67

IC26

0220

160.

4385

1−6

4.46

7941

1.22

+0.0

5−0

.05

52.6

4+0.0

3−0

.03

29.6

8−3

.18

300

0.19

±0.0

5

IC26

0248

160.

6727

9−6

4.35

1387

0.94

+0.0

1−0

.01

53.2

2+0.0

1−0

.01

30.2

9−3

.41

400

0.25

±0.0

2

IC26

0254

160.

7104

2−6

4.36

4891

0.93

+0.0

9−0

.20

51.5

5+0.0

6−0

.07

28.6

5−3

.06

1500

0.22

±0.1

6

IC26

0266

160.

8151

2−6

4.39

8415

0.62

+0.0

4−0

.04

52.1

3+0.0

3−0

.03

29.2

2−4

.37

800

0.13

±0.0

8

IC26

0267

160.

8449

2−6

4.48

6832

0.80

+0.1

9−0

.08

51.8

8+0.0

5−0

.05

28.9

9−3

.21

1000

0.21

±0.1

4

IC26

0287

161.

0408

3−6

4.24

7528

0.95

+0.0

6−0

.18

52.1

2+0.0

5−0

.05

29.2

2−3

.39

1000

0.19

±0.1

0

IC26

0290

161.

0873

0−6

4.50

2136

0.72

+0.2

4−0

.29

51.7

0+0.1

2−0

.14

28.8

1−3

.08

IC26

0291

161.

0932

2−6

4.25

8446

0.73

+0.0

7−0

.07

52.0

0+0.0

5−0

.05

29.1

1−4

.40

1000

0.28

±0.1

2H

ogg1

510

190.

7904

7−6

3.07

6221

31.0

3−3

.43

Hog

g15

1319

0.80

016

−63.

1468

8931

.15

−3.3

4H

ogg1

514

190.

8017

9−6

3.10

2390

0.96

+0.3

1−0

.27

54.2

1+0.1

4−0

.16

31.3

1−4

.69

Hog

g15

1619

0.82

042

−63.

0842

510.

63+0

.19

−0.4

054

.35+0

.72

−0.1

731

.45

−3.9

6H

ogg1

518

190.

8285

1−6

3.08

3611

31.0

6−4

.98

Hog

g15

2119

0.83

771

−63.

0157

510.

830+0

.14

−0.2

154

.44+0

.13

−0.1

331

.55

−4.1

8H

ogg1

524

190.

8581

2−6

3.10

5526

31.5

6−2

.67

Hog

g15

3019

0.89

046

−63.

0659

4531

.03

−4.7

0H

ogg1

531

190.

8914

5−6

3.07

3444

30.9

2−5

.40


Tabl

e5.

(Con

tinu

ed).

Sour

cede

tect

ion

Spec

tral

Tim

ing

Clu

ster

RA

J200

0D

EC

J200

0kT

log(

EM

)lo

g(L

X)†

Bin

size

nam

eID

(deg

)(d

eg)

(keV

)(c

m−3

)(e

rgs

s−1)

log(

LX

Lbo

l)

(s)

Fva

r

Hog

g15

3219

0.89

250

−63.

0017

7830

.92

−4.6

6H

ogg1

535

190.

9074

7−6

3.15

7776

0.51

0+0.2

4−0

.21

54.4

9+0.3

6−0

.20

31.5

4−3

.83

Hog

g15

3719

0.93

520

−63.

0998

882.

34+0

.81

−0.4

954

.71+0

.06

−0.0

731

.77

−2.5

910

000.

30±0

.15

Hog

g15

5019

0.98

984

−63.

1640

010.

640+0

.19

−0.1

354

.64+0

.11

−0.1

331

.73

−4.0

4H

ogg1

552

190.

9998

3−6

3.00

8915

30.7

6−3

.49

Hog

g15

5619

1.00

992

−63.

1715

0130

.92

−3.4

1H

ogg1

570

191.

0963

0−6

3.06

2611

31.4

2−4

.18

Hog

g15

7119

1.09

637

−63.

1234

440.

860+0

.21

−0.2

754

.12+0

.13

−0.1

631

.33

−4.1

1H

ogg1

572

191.

1033

3−6

3.09

9804

30.8

7−4

.68

Tru

mpl

er18

7916

7.89

359

−60.

6552

241.

92+2

.87

−0.7

553

.36+0

.10

−0.1

230

.38

−5.2

6

Tru

mpl

er18

8416

7.90

800

−60.

6667

79>

26.5

053

.82+0

.07

−0.1

130

.92

−2.7

5

Tru

mpl

er18

8516

7.90

945

−60.

6795

0112

.60+3

0.50

−5.9

753

.86+0

.10

−0.0

631

.05

−3.5

4

Not

es:

Col

umn

1:C

lust

erna

me;

Col

umn

2:id

entifi

catio

nnu

mbe

r(I

D)

inSu

pple

men

tary

Tabl

e3;

Col

umn

3an

d4

repr

esen

tth

epo

sitio

nof

X-r

ayso

urce

;C

olum

n5,

6,7:

estim

ated

valu

esof

coro

nal

tem

pera

ture

s(k

T),

emis

sion

mea

sure

(EM

)an

dX

-ray

lum

inos

ities

(log

(LX))

from

eith

ersp

ectr

alfit

ting

usin

gC

-sta

tistic

sor

deri

ved

from

conv

ersi

onof

coun

tra

tes

into

X-r

ayflu

xes

usin

gC

CFs

†.C

olum

n8:

Tim

ebi

nsi

zein

s;C

olum

n9:

repr

esen

tfr

actio

nal

root

mea

nsq

uare

vari

abili

tyam

plitu

de(F

var)

with

erro

rsan

dno

tdefi

ned

whe

nS2

isle

sser

than

〈σ2 er

r〉(f

orde

tails

see

§4).

†:X

-ray

flux

deri

ved

from

spec

tral

fittin

gar

eco

nver

ted

into

lum

inos

ities

usin

gth

edi

stan

ceto

thei

rco

rres

pond

ing

clus

ters

(see

Tabl

e1)

.The

spec

tral

para

met

ers

are

not

deri

ved

for

the

star

sw

ithpo

orco

unt

stat

istic

san

dth

eir

unab

sorb

edX

-ray

fluxe

sha

vebe

enes

timat

edby

thei

rco

unt

rate

sin

EPI

Cde

tect

orus

ing

CC

Fs(W

ebPI

MM

S),i

.e.,

flux

=C

CF

×co

untr

ates

.The

valu

esof

CC

Fs(i

nun

itsof

erg

s−1

cm2)

are

deri

ved

for

PNan

dM

OS

dete

ctor

s.Fo

rIn

term

edia

tem

ass

star

s:3.

926

×10

−12

and

1.24

7×

10−1

1fo

rN

GC

869

at2.

07ke

V;3

.528

×10

−12

and

1.18

1×

10−1

1fo

rN

GC

884

at1.

30ke

V;4

.799

×10

−12

and

1.43

3×

10−1

1

for

NG

C73

80at

2.38

keV

;1.

835

×10

−11

and

6.74

0×

10−1

1fo

rH

ogg

15at

0.29

keV

;Fo

rlo

wm

ass

star

s:4.

937

×10

−12

and

1.84

7×

10−1

1fo

rN

GC

663

at0.

71ke

V;

3.92

6×

10−1

2an

d1.

292

×10

−11

for

NG

C86

9at

2.05

keV

;3.

610

×10

−12

and

1.20

6×

10−1

1fo

rN

GC

884

at1.

59ke

V;

4.93

1×

10−1

2an

d1.

460

×10

−11

for

NG

C73

80at

2.80

keV

;6.0

54×

10−1

2an

d1.

749

×10

−11

for

Ber

kele

y86

at1.

87ke

V;7

.689

×10

−11

and

2.29

0×

10−1

1fo

rH

ogg

15at

0.97

keV

.


Figure 2. Background subtracted X-ray light curve of the source with ID #20 in the opencluster IC 2602. This source is close to the BY Dra type variable star V554 Car.

The Fvar is found to be more than 3σ of its error for seven sources (see Table 5),therefore, these sources are considered as variable. The light curves of six sourcesshow characteristics of flares and analyses of these flares are presented in Bhatt et al.(2013) (hereafter, Paper II). The background subtracted light curve of one remainingsource with ID #20 in the cluster IC 2602 with ID #20 in the cluster IC 2602 is shownin Fig. 2. This source is very close to V554 Car, which is classified as BY-Dra typevariable (Kazarovets et al. 2001). The X-ray light curve of the source shows ∼20%of variability with respect to its mean count rate during observational time scaleand does not show any flare-like feature. BY Dra type of star may have rotationalmodulation in X-rays (see Patel et al. 2013).

5. X-ray spectra

Spectral characteristics of stars are also required before one can estimate their lumi-nosity functions. X-ray spectra of the sources with counts greater than 40 have beengenerated using the SAS task ESPECGET, which also computed the photon redistri-bution matrix and ancillary matrix. For each source, the background spectrum wasobtained from source-free regions chosen according to the source location (sameregions as used in the generation of light curves). Spectral analysis was performedbased on global fitting using the Astrophysical Plasma Emission Code (APEC) ver-sion 1.10 modelled by Smith et al. (2001) and implemented in the XSPEC version12.3.0. The plasma model APEC calculates both line and continuum emissivities for a


hot, optically thin plasma that is in collisional ionization equilibrium. The absorptiontowards the stars by interstellar medium was accounted for by using a multiplicativemodel PHABS in XSPEC which assumes the photo-electric absorption cross sectionsaccording to Balucinska–Church & McCammon (1992).

The simplest spectral model considered, is that of an isothermal gas which werefer to as the ‘1T APEC’ model. This model is expressed as PHABS × APEC. Weadopted the approach used in Currie et al. (2009) for X-ray spectral fitting and usedan initial temperature kT of 1.5 keV to start the spectral fitting which is a compromisebetween values typical of stars in younger clusters (e.g., M17; Broos et al. 2007) andstars in older clusters (e.g., the Pleiades; Daniel et al. 2002). Elemental abundanceparameter with a value of 0.3 solar is routinely found in fits of stellar X-ray spectra,and was thus fixed to this value in our analysis (Feigelson et al. 2002; Currie et al.2009) for intermediate and low mass stars. For massive stars, however, abundanceparameter of 0.2 solar was fixed for fitting (see Bhatt et al. 2010; Zhekov & Palla2007). The value of absorption column density, NH , was fixed throughout the fittingto the value derived using the relation given by Vuong et al. (2003), NH = 5×1021×E(B − V ) cm−2, and given in Table 1. The temperature, kT and the normalizationwere the free parameters in spectral fitting. We performed C-statistic model fittingtechnique rather than using χ2-minimization technique because of the poor countstatistics. The temperature, normalization and unabsorbed flux values were derivedby this fitting technique. The estimated temperatures, EM and luminosities are givenin Table 6 for the massive stars and in Table 5 for the intermediate and low massstars. A few examples of X-ray spectra of massive, intermediate and low mass starsare shown in Fig. 3 along with the ratios of the X-ray data to the fitted model in thelower panels.

The X-ray fluxes of the probable cluster members having very poor count statistics(counts below 40) or which were lying between the inter-chip gaps between PNCCDs, were derived from their X-ray count rates in the EPIC detectors estimatedfrom the SAS task EDETECT_CHAIN (see section 2.1 and supplementary material).The CCFs to convert count rates into X-ray fluxes were estimated from WebPIMMS5

using 1T APEC plasma model. The value of the model parameter NH was fixedfrom Table 1 for the respective clusters. However, abundance parameter was fixed at0.2 solar for massive stars (Zhekov & Palla 2007; Bhatt et al. 2010), and 0.3 solarfor intermediate and low mass stars (Feigelson et al. 2002; Currie et al. 2009). Theplasma temperature was fixed at 1.0 keV for massive stars (Nazé 2009). However, forintermediate and low mass stars, the plasma temperature was taken as the mean of thetemperatures derived from the spectral fitting of other bright stars in the cluster. Themean values of X-ray temperatures of intermediate mass stars have been found to be2.07, 1.30, 2.38 and 0.29 keV for the open clusters NGC 869, NGC 884, NGC 7380and Hogg 15, respectively. In the case of low mass stars, the mean values of X-raytemperatures have been found to be 0.71, 2.05, 1.59, 2.80, 1.87, 1.06 and 0.97 keVfor the open clusters NGC 663, NGC 869, NGC 884, NGC 7380, Berkeley 86, IC2602 and Hogg 15, respectively. The derived values of conversion factors of countrates into unabsorbed fluxes for massive, intermediate and low mass stars have been

5http://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3pimms/pim_adv

http://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3pimms/pim_adv


Tabl

e6.

X-r

ayte

mpe

ratu

res

and

lum

inos

ities

ofm

assi

vest

ars

with

inyo

ung

clus

ters

.

Off

bkT

avlo

g(L

X)

(Ref

eren

ces)

Clu

ster

IDN

ame

Spec

tral

type

Ba

(′′)

(keV

)(e

rgs−

1)

log(

LX

Lbo

l)

NG

C66

326

BD

+60

329

B1V

2.3

31.1

6(N

azé

2009

)−6

.69

NG

C66

369

V83

1C

asB

1VY

1.1

1.11

+0.0

7−0

.06

34.1

5(L

aPa

lom

bara

&M

ereg

hetti

2006

)–

NG

C86

912

HD

1396

9B

0.5I

1.8

30.5

3††(P

rese

ntst

udy)

−7.3

7N

GC

869

50H

D14

052

B1.

5I4.

30.

44+0

.18

−0.1

131

.06

(Pre

sent

stud

y)−7

.34

NG

C86

910

3[S

HM

2002

]12

0B

1.5V

6.3

30.8

4††(P

rese

ntst

udy)

−7.1

8N

GC

869

109

BD

+56

527

B2I

2.0

1.12

+0.3

7−0

.21

31.0

2(P

rese

ntst

udy)

−7.3

3

NG

C86

911

9[S

HM

2002

]13

8B

5.6

4.64

+21.

04−2

.531

.09

(Pre

sent

stud

y)−6

.43

NG

C88

479

[SH

M20

02]

131

B1.

5III

4.4

30.8

0(N

azé

2009

)−6

.76

NG

C88

493

BD

+56

578

B2I

II2.

230

.89

(Naz

é20

09)

−7.1

4N

GC

7380

36D

HC

epO

5.5V

+O6.

5VY

0.6

0.64

+0.0

2−0

.02

32.4

3(B

hatt

etal

.201

0)−6

.71

NG

C73

8077

LS

III+5

790

O8V

((f)

)Y

1.7

31.5

5(N

azé

2009

)−6

.66

Be8

685

HD

2289

89O

9V+O

9VY

2.7

0.62

+0.1

5−0

.11

31.7

2(P

rese

ntst

udy)

−6.6

2H

ogg1

53

HD

1104

32B

0.5V

epY

2.2

8.0

–11.

035

.00

(Lop

esde

Oliv

eira

etal

.200

7)–

Hog

g15

49M

O1-

78O

B1.

90.

72+0

.23

−0.1

731

.39

(Pre

sent

stud

y)−6

.91

Hog

g15

692.

130

.94††

(Pre

sent

stud

y)−6

.72

Hog

g15

730.

430

.94††

(Pre

sent

stud

y)−7

.02

a ‘Y

’re

pres

ents

the

bina

rity

ofst

ars

from

liter

atur

e.bO

ffse

tbet

wee

nth

epo

sitio

nof

mas

sive

star

sin

2MA

SSca

talo

gue

and

thei

rX

-ray

coun

terp

arts

.††

LX

has

been

deri

ved

from

the

X-r

ayflu

xes

timat

edfr

omth

eco

unt

rate

conv

ersi

onin

PNde

tect

orus

ing

the

Web

PIM

MS,

i.e.,

flux

=C

CF

×co

untr

ate.

The

valu

esof

CC

Fs(i

nun

itsof

erg

s−1

cm2)h

ave

been

deri

ved

form

assi

vest

ars

inPN

are

3.60

6×

10−1

2fo

rNG

C86

9an

d7.

689

×10

−12

forH

ogg

15,r

espe

ctiv

ely.


10−5

10−4

10−3

0.01

norm

aliz

ed c

ount

s s−

1 ke

V−

1

NGC 869 : #50 (HD 14052)

10.5 2

0

5

10ra

tiono

rmal

ized

cou

nts

s−1

keV

−1

ratio

norm

aliz

ed c

ount

s s−

1 ke

V−

1 ra

tio

Energy (keV)

NGC 869 : #139

10.5 520

1

2

3

Energy (keV)

NGC 869 : #156

10.5 52

2

4

Energy (keV)

10−4

10−3

10−4

10−3

(a)

(b)

(c)

Figure 3. A few example of X-ray spectra of (a) massive star, (b) intermediate mass star and(c) low mass star. The ID of the star with the information of its respective cluster is given at thetop of each panel.


given in the footnotes of Tables 5 and 6. For the sources either falling in betweenthe inter-chip gaps of PN CCDs or outside the PN coverage area, the X-ray fluxesin the MOS1 and the MOS2 detectors were estimated from their count rates usingCCFs in the MOS detector. The average value of X-ray flux in the MOS1 and MOS2detectors has been quoted in Table 5. Thus, the X-ray luminosities were estimatedfrom the derived values of the X-ray fluxes and given in Table 6 for massive starsand in Table 5 for intermediate and low mass stars.

X-ray spectrum of star #79 in Berkeley 86 could not be fitted with the model usedfor spectral fitting, therefore, the NH parameter was varied as a free parameter. Thebest-fit value of NH has been found to be 3×1022 cm−2, which is 8 times higherthan that expected in the direction of the open cluster Berkeley 86 (see Table 1).This points to either very high intrinsic extinction in the source or the source doesnot belong to the open cluster Berkeley 86. For the stars showing flares, the valuesof parameters listed in Table 5, were derived from the spectral fitting performed fortheir quiescent state data.

6. X-ray properties of stars in different mass groups

X-ray spectral properties of massive, intermediate and low mass stars were analysedseparately, because the production mechanism of X-rays are different for differenttypes of stars. The bolometric luminosities (Lbol) of the stars were derived from theirbolometric magnitudes (mbol). The absolute J0 magnitudes were estimated fromtheir observed 2MASS J magnitudes using well constrained age, reddening and dis-tance parameters of the corresponding open clusters in literature (see Table 1). Thembol of the stars were derived from their J0 magnitudes by interpolating the mbolbetween J0 magnitude points in theoretical isochrones of Girardi et al. (2002) forMS stars and Siess et al. (2000) for PMS stars, depending upon the age of the opencluster.

6.1 Massive stars

Our sample contains 16 massive stars of which 6 were reported previously (see ref-erences in Table 6). The best fit spectral parameters for 8 stars are given in Table 6.The X-ray fluxes for four massive stars (see Table 6) were derived from their countrates in PN detector by using CCFs. The X-ray temperatures are found to be lessthan 1.2 keV in general. However, X-ray temperatures are found to be higher in thecase of high mass X-ray binary HD 110432 (Lopes de Oliveira et al. 2007) and[SHM202] 138. The LX of massive stars lie in the range of 1031−35 erg s−1. TheLX/Lbol for each massive star is derived and given in Table 6. The average valueof log(LX/Lbol) is found to be −6.92 with standard deviation of 0.31. This valueof LX/Lbol is broadly consistent with the value derived for a sample of nearly 300massive stars by Náze (2009).

6.2 Low mass stars

Although, there is strong evidence that X-ray emission originates from magneticallyconfined coronal plasma in the PMS low mass stars (e.g., Preibisch et al. 2005), the


relationship between rotation and X-ray activity in PMS low mass stars remainedunclear. During the PMS phase, the low mass stars undergo substantial changes intheir internal structure, evolving from fully convective structure to a radiative coreplus convective envelope structure. Consequently, the stellar properties of low massstars—Lbol, magnetic activity and rotation etc., are also changing during the PMSphase. The dependence of Lbol and age upon X-ray emission is examined in thefollowing sections.

6.2.1 X-ray temperatures. Most of these sources have plasma temperaturesbetween 0.2 and 3 keV which are consistent with values derived for PMS stars inyoung clusters e.g., NGC 1333 (Getman et al. 2002), Orion (Feigelson et al. 2002),NGC 1893 (Caramazza et al. 2012) and M16 (Guarcello et al. 2012). The averageplasma temperature of the stars in the open clusters appears to be constant for allstars undergoing PMS evolution from 4 Myr to 46 Myr, and the median value isfound to be ∼1.3 keV.

6.2.2 X-ray luminosity functions and their evolution with age. X-ray Luminos-ity Functions (XLFs) of low mass stars in different clusters have been derivedusing Kaplan Meier (KM) estimator of integral distribution functions and shown inFig. 4(a). No significant difference is observed in the XLFs of low mass stars withages in the range of 4 to 14 Myr. However the XLF of low mass stars in the open clus-ter IC 2602 with an age of 46 Myr appears to be lower than that of others. The meanvalues of log LX with their standard deviations have been found to be 31.26 ± 0.38,30.82 ± 0.31, 30.81 ± 0.26, 31.22 ± 0.31, 31.01 ± 0.18, 29.10 ± 0.65, 31.24 ± 0.32and 30.78 ± 0.35 erg s−1 for the open clusters NGC 663, NGC 869, NGC 884, NGC7380, Berkeley 86, IC 2602, Hogg 15 and Trumpler 18, respectively. The mean val-ues of LX of low mass PMS stars are thus nearly similar in all the open clustersexcept IC 2602.

The evolution of the mean value of log LX with age is shown in Fig. 4(b). Themajority of the low mass stars in our sample have masses greater than 1.4M� as seenin Fig. 1, except for the stars in IC 2602. The stars in the open cluster IC 2602 withmasses above 1.4M� may have LX below 27.57 erg s−1 (detection limits) and arenot detected in the present study. It indicates a sudden decrease in the LX between14 to 46 Myr for the stars with masses above 1.4M�. Thus LX is nearly constantduring the evolution of low mass stars in PMS phase from 4 to 14 Myr and maydecrease thereafter. Scholz et al. (2007) reported that the rotation rates increase inthe first few Myr of their evolution. It is, therefore, possible that an increase in theX-ray surface flux due to an increase in the rotation rate may be compensated by adecrease in the stellar surface area during PMS evolution, between 1 to 10 Myr, asdescribed by Preibisch (1997). Between 10 to 40 Myr, the decrease in LX may belinked with a rapid spin down in the stars, as suggested by Bouvier et al. (1997).However, the faintest cluster members have not been detected here, therefore, thecomplete XLFs of these clusters cannot be derived since the mean luminosities ofthe entire cluster population may be lower than these values. Further, in case of lowmass close binaries the hot winds produced by the coronae in young stars may drivethe evolution of X-rays (Iben and Tutukov 1984).


(a)

(b)

Figure 4. X-ray luminosities of low mass stars. (a) XLFs of low mass stars in differentclusters. (b) Evolution of mean LX of the clusters with age.


(a)

(b)

Figure 5. (a) Relation between LX and Lbol for low mass stars in the sample. Dashed linesin each plot represent the isopleths of log(LX/Lbol) and the values are given above each line.(b) Distribution of log LX/Lbol ratio for all the low mass stars in all the clusters.


6.2.3 X-ray to bolometric luminosity ratios. LX/Lbol provides an estimate of thefraction of total stellar energy that is dissipated through coronal heating, and theestimated values of the LX/Lbol for low mass stars in our sample are listed inTable 5. The mean values of log(LX/Lbol) with standard deviations are foundto be −3.86 ± 0.41, −3.63 ± 0.49, −3.63 ± 0.51, −3.12 ± 0.53, −4.00 ± 0.84,−3.52 ± 0.49, −4.02 ± 0.78 and −3.85 ± 1.28 for the open clusters NGC 663, NGC869, NGC 884, NGC 7380, Berkeley 86, IC 2602, Hogg 15 and Trumpler 18,respectively. These values have been found to be consistent with the values derivedfor the young clusters: the Orion (−3.39 ± 0.63), IC 348 (−3.53 ± 0.43) and NGC2547 (−3.20 ± 0.24) (see Alexander & Preibisch 2012). The derived values of meanlog(LX/Lbol) for each cluster are similar and the mean value log(LX/Lbol) is foundto be −3.6 with a standard deviation of 0.4 for the collective sample of low massstars within these open clusters.

The relation between LX and Lbol along with the isopleths of log(LX/Lbol) =−3.0, −4.0 are shown in Fig. 5(a). It can be seen that most of the sources haveLX/Lbol values below the saturation level. The distribution of log(LX/Lbol) for allthe low mass stars in all the clusters has been shown in Fig. 5(b) which is derivedusing the KM estimator of integral distribution functions. It shows that only 15%of the X-ray sources have LX/Lbol values above the saturation level. There are fivesources with log(LX/Lbol) greater than −2.5; values that are very unlikely to be

Figure 6. Relation between (LX/Lbol) and Lbol for low mass stars in the sample (dots), forstars with age from 4 to 14 Myr (open circles), and for the stars in the cluster IC 2602 withage of 46 Myr (open triangles) derived using least square fitting and shown by continuous line,dashed line and dashed plus dotted line, respectively. The stars with log(LX/Lbol) above −2.5have not been considered while deriving these relations and are marked by the symbol of cross.


found for stellar sources. This is possibly the result of these sources not being mem-bers of the corresponding clusters, therefore, their LX and Lbol may have not beenestimated properly. The LX values of these few sources are marked with a symbolof question mark in Table 5. These sources are marked with the symbol of cross inFig. 6, and were not considered for further analysis. The evolution of log(LX/Lbol)

with age is shown in Fig. 7 and found that the log(LX/Lbol) is nearly constant during4 to 46 Myr.

Among the low mass PMS stars in the Orion, the median LX/Lbol is 2–3 ordersof magnitude greater, i.e., ≈10−3, than that found within ZAMS stars and thereforetheir fractional X-ray luminosities are ‘saturated’. Currie et al. (2009) showed thatthe stars with masses >1.5M� deviate from X-ray saturation by ≈10–15 Myr. Thepresent analysis indicates that most of the low mass PMS stars come out from thesaturation limit earlier than 4–8 Myr, which is quite early as compared to the agedescribed by Currie et al. (2009), i.e., 10–15 Myr. The X-ray emission depends uponthe magnetic dynamo that is the result of a combination of turbulent convectionand rotation within the convection zone. As a low mass star contracts onto the MS,its internal structure changes and its outer convective zones shrinks. Therefore, theevolution of fractional X-ray luminosity with age might be due to either the changein the internal structure of a star or spin-down rotation of a star during the PMSphase, or both. Alexander & Preibisch (2012) showed that there was no correlationbetween the LX/Lbol and the rotation period. They also found some rather slowlyrotating stars (period > 10 days) with very strong X-ray activity, and suggested thatinstead of rotation it is the change in the internal structure of PMS stars during the

Figure 7. Evolution of mean LX/Lbol of the clusters with age.


evolution which is likely to be responsible for the generation of magnetic dynamoand consequently the X-ray emission.

The dependence of LX/Lbol on Lbol is shown in Fig. 6. The correlation coeffi-cients between L X/Lbol and Lbol have been derived using Pearson product-momenttest and Kendall tau rank test, and their values are found to be −0.65 and −0.58,respectively for all the low mass stars in the sample. Thus, the probability of nocorrelation between LX/Lbol and Lbol, i.e., the null hypothesis, is estimated to be2.2 × 10−16 from both the tests.

The linear regressions have been calculated using the least-squares Marquardt–Levenberg algorithm (Press et al. 1992) corresponding to the following relation forall the low mass stars and shown by continuous line in Fig. 6,

log(LX/Lbol) = −0.48(±0.05) × log(Lbol) + 12.95(±1.63). (3)

For the low mass stars with ages between 4 to 14 Myr and shown by dashed linein Fig. 6,


For the low mass stars in the open cluster IC 2602 with ages 46 Myr and shownby dashed and dotted line in Fig. 6,


Equation (3) shows a power-law dependence of the fractional X-ray luminosityon Lbol during 4 to 46 Myr. The power-law indices are found to be different for starswith age of 4–14 Myr and the stars in the cluster IC 2602 with age ∼46 Myr fromeq. (4) and eq. (5), respectively. Prebisch et al. (2005) showed LX ∝ Lbol for thestars in Orion which implies that LX/Lbol is nearly constant at 1 Myr. For NEXXUSsample of nearby field stars (Schmitt & Liefke 2004), Prebisch et al. (2005) foundLX ∝ L0.42

bol , which implies that (LX/Lbol) ∝ L−0.58bol . The values of power-law

indices of LX and LX/Lbol relation for the open clusters NGC 663, NGC 869, NGC884, NGC 7380, Berkeley 86, IC 2602 and Hogg 15 are derived to be −0.7 ± 0.2,−0.8 ± 0.1, −1.1 ± 0.1, −1.0 ± 0.1, −0.9 ± 0.1, −0.4 ± 0.2, −1.0 ± 0.1, respec-tively. It implies that the (L X/Lbol) depends upon Lbol during 4 to 46 Myr and thisdependence upon Lbol may be started earlier than 4 Myr. As low mass stars evolveto MS, their effective temperatures eventually increase and the depth of their con-vective envelopes reduce, therefore their Lbol changes. During 4 Myr to 46 Myr, theLbol increases nearly three times (Siess et al. 2000) for low mass star with masses inthe range of 1.4–2.0M�. This increase in Lbol can produce a decrease of nearly one-third in (LX/Lbol) which can give a decrease of nearly 0.5 dex in logarithmic scale.Such a variation cannot be distinguished using present data because the standarddeviation in log(LX/Lbol) is comparable with the decrease of 0.5 dex.

6.3 Intermediate mass stars

A convincing and unique explanation for the generation of X-ray emission fromintermediate mass stars has not been forthcoming, despite abundant speculationsabout the possible mechanisms. The presence of magnetic field of the order of a fewhundred Gauss (Donati et al. 1997; Hubrig et al. 2004; Wade et al. 2005) has been


detected in these stars, that can support a shear dynamo which may be responsiblefor X-ray emission from intermediate mass stars as in the T-Tauri stars. At the sametime, the option of unresolved companions is also considered because the interme-diate mass stars are more likely to be found in binaries, i.e., companion hypothesis(Baines et al. 2006; Stelzer et al. 2006 and references therein).

The detection limits are in the range of 1027.6–1030.8 erg s−1 for different clus-ters as they are located at different distances from the Earth. For making a sample ofintermediate mass stars from different clusters, a highest detection limit of log LX ≈1030.8 erg s−1 among all clusters (see also Table 4) was used, which shows that astar with log LX > 30.8 erg s−1 could be detected in any of the clusters. In this way,a total of 27 intermediate-mass stars were identified and examined further. The XLFsof the low mass and intermediate mass stars having log LX > 30.8 erg s−1 werederived using the KM estimator of integral distribution functions. A comparison ofthe X-ray luminosities and fractional X-ray luminosities of low mass stars and inter-mediate mass stars in the present sample is shown in Fig. 8. The results of twosample tests are given in Table 7. The results of the Wilcoxon Rank Sum, Logrank,Peto and Peto Generalized Wilcoxon and Kolmogorov–Smirnov (KS) statistical testsshow that the X-ray luminosity distribution of intermediate mass stars is differentfrom that of low mass stars with confidence of 93%, 98%, 92% and 77%, respec-tively. Therefore, the X-ray luminosities of both types of stars above this limit oflog LX(>30.8) erg s−1 are not significantly different from each other. Further, the

(a) (b)

Figure 8. Comparison of the X-ray activity of low mass and intermediate mass stars havinglog LX > 30.8 erg s−1 based on the Kaplan Meier estimator. (a) Distribution of LX and(b) distribution of LX/Lbol ratio.


Table 7. Results of two sample tests.

Statistics of objects in two groups

log(LX) ( erg s−1) > 30.8

Number of stars (low mass) 100Number of stars (intermediate mass) 27

Probability of having a common parent LX distribution

Wilcoxon rank sum test 0.07Logrank test 0.02Peto and Peto generalized Wilcoxon test 0.08Kolmogorov–Smirnov test 0.23

log(LX/Lbol)

Probability of having a common parent LX/Lbol distribution

Wilcoxon rank sum test 8.1×10−13

Logrank test 0.0Peto and Peto generalized Wilcoxon test 0.0Kolmogorov–Smirnov test 9.6×10−12

L X/Lbol ratio of intermediate mass stars and low mass stars are different, with aconfidence of greater than 99.999% using these statistical tests.

Recently, Balona (2013) suggested the light variation due to rotation modulationcaused by star-spots in nearly 875 A-type stars using Kepler’s data. If A-type starshave spots, then it is natural to expect a magnetic field, and therefore X-ray activity inintermediate mass stars. The median values of log(LX/Lbol) are found to be −5.06and −3.41 for intermediate mass stars and low mass stars, respectively. It impliesthat if the intermediate mass stars themselves produce X-rays, the strength of theX-ray activity is possibly weaker as compared to the low mass stars. However, thepossibility of the X-ray emission from a nearby low mass star cannot be ruled outhere due to the poor spatial resolution data of XMM-NEWTON.

7. Summary and conclusions

We have described the X-ray source contents of eight young open clusters using theXMM-NEWTON data. These clusters have ages ranging from 4 Myr to 46 Myr andthus provide a link between the X-ray properties of young clusters like the Orionand older clusters like the Pleiades. The association and membership of these X-raysources with stars has been deduced using optical and NIR data. Overall 152 X-raysources have been identified with low mass PMS stars, 36 with intermediate massstars and 16 with massive stars. The main results are summarized below:

(1) The X-ray temperatures, luminosities and fractional X-ray luminosities of mas-sive stars are consistent with the values reported previously in the literature for othermassive stars.(2) The plasma temperatures are found to be in the range of 0.2 keV to 3 keV witha median value of 1.3 keV for all low mass stars irrespective of their ages.


(3) The observed XLFs of low mass stars in the open clusters with ages from 4 to14 Myr appear to be similar, which implies that LX is nearly constant during PMSevolution from 4 to 14 Myr. Therefore, the decrease in LX of low mass stars mayoccur during 14 to 100 Myr. Non-detection of X-rays from the stars above 1.4M�in the open cluster IC 2602 may give an indication of a sudden decrease in their LXduring 14 to 46 Myr.(4) The log(LX/Lbol) of most of the low mass stars are below the saturation limitsand the mean value has been found to be −3.6 with a standard deviation of 0.4. Thisvalue is consistent with the values derived for other young clusters the Orion, IC 348and NGC 2547. Thus, a deviation of low mass stars with masses greater than 1.4M�from X-ray saturation may occur before the age of 4–8 Myr, earlier than the agederived by Currie et al. (2009), i.e., 10–15 Myr.(5) The (LX/Lbol) of low mass stars correlate well with their Lbol, suggesting itsdependence on the internal structure of stars.(6) No statistically significant difference in LX from the intermediate mass andthe low mass PMS stars has been detected. But the observed LX/Lbol for interme-diate mass stars have been found to be significantly lower than that of low massstars. It possibly indicates that the strength of X-ray activity in intermediate massstars is weaker than in the low mass stars. Another possibility is that the origin ofX-ray emission from intermediate mass stars might be the result of X-ray emissioncoming from an unresolved nearby low mass PMS star. Deeper and higher spatialresolution data with CHANDRA is needed to check for this possibility and to estimatethe complete XLFs of these clusters.

Acknowledgements

The authors would like to thank the anonymous referee for his/her constructive com-ments. This publication makes use of data from the Two-Micron All-Sky Survey,which is a joint project of the University of Massachusetts and the Infrared Process-ing and Analysis Center/California Institute of Technology, funded by the NationalAeronautics and Space Administration and the National Science Foundation, anddata products from XMM-Newton archives using the high energy astrophysicsscience archive research center which is established at Goddard by NASA. Weacknowledge XMM-Newton Help Desk for their remarkable support in X-ray dataanalysis. Data from Simbad, VizieR catalogue access tool, CDS, Strasbourg, Francehave also been used. HB is thankful for the financial support for this work throughthe INSPIRE Faculty Fellowship granted by the Department of Science & Technol-ogy India. They also acknowledge R. C. Rannot, Nilesh Chouhan and R. Koul fortheir support to complete this work.

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X-ray Observations of Eight Young Open Star Clusters: I. Membership and X-ray Luminosity

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Transcript of X-ray Observations of Eight Young Open Star Clusters: I. Membership and X-ray Luminosity