NASA/CR.- / ,_ c/_,., 208130

NATIONAL

OPTICAL

ASTRONOMY

OBSERVATORIES

/N-d __- E_

Preprint Series

NOAO Preprint No. 775

X-ray Activity in the Open Cluster IC 4665

Mark S. Giampapa

Charles F. Prosser

Thomas A. Fleming

Accepted for publication in: The Astrophysical Journal

December 1997

Operatedfor ttleNationalScienceFounda_onby IheAssociationof Universitiesfor ResearchinAstronomy,Inc.

X-ray Activity in the Open Cluster IC 4665

Mark S. Giampapa 1, Charles F. Prosser 2, and Thomas A. Fleming 3

Accepted for publication in The Astrophysical Journal

1 National Optical Astronomy Observatories, National Solar Observatory, P.O. Box

26732, Tucson, AZ 85726-6732. The National Optical Astronomy Observatories are operated for the National Science Foundation by the Association of Universities for Research

in Astronomy. E-mail: giampapa@noao.edu.

2 Harvard-Smithsonian Center for Astrophysics, 60 Garden St. MS-66, Cambridge, MA

02138. Current address: NOAO, P. O. Box 26732, Tucson, AZ 85726-6732. E-mail:

cprosser@noao.edu.

3 Steward Observatory, University of Arizona, Tucson, AZ 85721. E-mail: tfleming_as.arizona.edu.

ABSTRACT

We present the results of a joint ROSAT High Resolution Imager (HRI) and optical

investigation of the open cluster IC 4665. The ROSAT data contains detections for 28

stellar sources in the field, including 22 cluster members and candidate members spanning

the color range -0.18 <_ (B - V)o <_ +1.63 (--_ B3 - M3). Upper limits are given for the

remaining members (or candidate members) in the HRI field. Keck HIRES spectra have

been obtained that yield radial and rotational velocity measures, respectively, for faint,

low mass candidate members located within the field of the ROSAT HRI observation. In

addition, photometry of possible optical counterparts to previously uncatalogued X-ray

sources in the HRI field is presented.

The trends in X-ray properties with (B - V) color in IC 4665 are found to be quite

similar to that for other, more nearby young clusters such as the Pleiades and a Persei. In

particular, a maximum in normalized X-ray luminosity of log (Lx/Lbot) _- -3 is observed,

beginning in the color range of (B - V)o = 0.7 - 0.8. This is similar to the corresponding

color range among Pleiades members, in agreement with the earlier estimate by Prosser

& Giampapa (1994) that the age of IC 4665 is similar to the age of the Pleiades. The

correlation of rotation and X-ray emission levels is consistent with that in other young

clusters.

Among the high mass stars in IC 4665, five B stars are detected as X-ray sources. Of

these, one is a spectroscopic binary while the remaining objects are apparently single stars.

The level of intrinsic X-ray emission observed in the rapidly rotating (v sini > 200 km

s -1), single B stars is consistent with an origin due to shock heating of the ambient medium

by radiatively driven, rotationally enhanced winds. On the basis of these observations and

the results for other clusters, we argue that observed levels of X-ray emission in high mass

stars of log (Lx/Lbol) > -6.0 are likely due to intrinsic emission associated with the B

star itself rather than emission from a late-type secondary. Finally, our results suggest

that the initial angular momentum distributions and the subsequent evolution of rotation,

Lx, and (L_/Lbot) occur in practically identical fashions in clusters with ages _ 100 Myr.

1. INTRODUCTION

The open cluster IC 4665 is among those nearby clusters for which there is significant

membership information for both the high-mass and solar-type stars. The most recent

studies include a proper motion membership survey by Prosser (1993), in which earlier

work is summarized, followed by a radial velocity survey by Prosser & Giampapa (1994).

A summary of the properties of IC 4665 is given in Table 1; additional cluster names for

IC 4665 and positional information are from Alter et al. (1970).

With an age similar to the Pleiades (Prosser & Giampapa 1994), IC 4665 is distinguished from the other relatively nearby clusters surveyed by ROSAT by its high galactic

latitude, corresponding to a distance from the Galactic plane of order 100 pc. The location of IC 4665 relative to the galactic plane is suggestive of an environment for star

formation that was distinct from that in other open clusters. The kinematics of the IC

4665 cluster to reach large distances from the Galactic plane can indicate, for example,

that its chemical composition may be intrinsically different from other clusters in the solar

vicinity. The relationship between chemical composition and X-ray activity in solar-type

stars has been recently discussed in an investigation of NGC2516 (Jeffries et al. 1997). A

study of additional young clusters, such as IC 4665, represents an important comparison to

other clusters with extensive optical data bases that have also been the objects of intensive

investigations with ROSAT. In order to study the X-ray properties of the members of IC

4665, we have obtained a deep ROSAT observation utilizing the High Resolution Imager

(HRI). We also use the X-ray observations presented in this investigation to assess the

effectiveness of membership identification based on previous proper motion surveys. In

particular, we utilize the HRI data to determine if any additional cluster members would

be identified in X-rays which had been previously overlooked in optical studies.

2. OPTICAL DATABASE

Our optical database for IC 4665 is based upon the membership analysis in Prosser

(1993), along with the radial velocity survey by Prosser & Giampapa (1994). Rotation

periods from photometric monitoring of solar-type IC 4665 members have been reported in

Prosser (1993), Allain et al. (1996), and Prosser & Grantdn (1997). To further investigate

the low-mass membership of IC 4665 and to supplement the X-ray observations, new

observations of IC 4665 were conducted and are described below. The results of these

additional observations have been combined with our previous results to form our optical

database.

3
2.1 Spectral Type Determinations

MK spectral classifications have been previously obtained for the brighter members

of IC 4665 for the range B3-A5 (see Table 6 of Prosser 1993 and references therein). As

the recent proper motion and radial velocity surveys successfully restricted the fist of

candidate members to only the most likely members, we wish to continue the spectral

classifications begun by Abt and Levato (1975). Spectral types are useful for confirming

duster membership and estimating other stellar parameters (i.e. color excess, temperature,

approximate radius and mass, etc.) which are utilized in detailed abundance studies (e.g.,

lithium) or in interpreting additional data for these stars (e.g., X-ray fluxes, rotation

periods, etc.). A queue observing program utilizing the FAST CCD spectrograph at the

Whipple Observatory 60-inch telescope was carried out during 1994 and 1996. Spectra were

obtained using a 600 l mm -1 grating to provide ,,,2000/_ coverage at 1.49/_ resolution at

a central wavelength of ,-,4400 _. A series of standard stars classified on the MK system

in the A5-K7V range were observed and are listed in Table 2. The spectral types are from

Morgan et al. (1943), Johnson & Morgan (1953), Keenan & McNeil (1976), Keenan & Pitts

(1980), and Jacoby et al. (1984). A series of AT-K2III standards were observed to identify

luminosity criteria. IC 4665 stars observed include those Kopff stars previously classified

only on objective prism spectra (Table 6 in Prosser 1993); the radial velocity members

identified in Prosser & Giampapa (1994); and some of the brighter radial velocity (Vrad)

nonmembers. Since some of the vr_l nonmembers of Prosser _: Giampapa (1994) could

conceivably be binaries which are cluster members, spectral type information was obtained

for some of these cases in order to ascertain if their spectral types might be consistent with

duster membership and therefore make them more likely targets for further radial velocity

monitoring.

MK spectral classifications were obtained using a set of line strength ratios calibrated

by the standard star observations: H6/Ca4226, H6/Fe4045, HT/Fe4325 , and HT/G-band.

Earlier than F0, only the H6/Ca4226 ratio was considered, though none of our target stars

were found to fall in this range. Measures of the HT/Fe4325 ratio were not attempted for

stars earlier than _- F6. For a particular target, the spectral types resulting from each index

ratio were found to be consistent within a few subclasses in most cases, and an average, final

MK spectral type for each star was estimated from the results of the measured indices and

appearance of the spectrum. The MK spectral types for the stars in IC 4665 are given in

Table 3 along with corrresponding indicators of membership status based on the spectral

classification (i.e., "Y/Y?" = member/likely member, "N" = nonmember, and "?" =

uncertain membership). All stars are found to be dwarfs, except for the nonmember P106

4
which is classified as a giant. The spectral types for K65, K88, and P107 are derived solely

from the H$/Ca4226 ratio, and might therefore be considered slightly less well determined.

Among those Kopff stars with previous spectral type estimates (Table 6, Prosser 1993),

the new FAST spectral types in Table 3 are more consistent with the positions of these

stars in the color-magnitude diagram for the assumed cluster distance and reddening, thus

resolving a previous discrepancy in this cluster.

Once considered as a cluster candidate, P12 can now be identified as a nonmember on

the basis of its spectral type which is inconsistent with cluster membership for its observed

color and assumed cluster reddening. P12 was noted as a rapid rotator (v sin i __ 70 km

s -1) by Mart_n &Montes (1997); the original Hamilton echeUe observation reported in

Prosser (1993) is considered to be in error. P12 appears therefore to be a rapidly rotating

late-type field star which, as will be seen, is strongly detected in X-rays. Among observed

radial velocity nonmembers, it is interesting to note that those proper motion candidates

with the highest proper motion membership probabilities (P1, P33, P44, and P 159) all have

spectral types acceptable for membership - suggesting these stars as strong candidates for

further study. On the other hand, the radial velocity nonmembers with low proper motion

membership probabilities (K38=P6, P51, and P106), and the photometrically selected

star P177, have spectral types which are either not consistent with cluster membership

or are uncertain. Although photometry indicated that P20, P30, and P108 are likely not

members, spectral types were nevertheless ascertained for these stars since, as will be seen,

they were detected in X-rays. The resulting spectral types for these stars are not consistent

with membership in IC 4665.

2.2 EcheIle Observations

High-dispersion spectra of some Ha emission M dwarfs lying within the ROSAT HRI

observation were obtained during 1994 Sept. 23-24 UT with the HIRES spectrograph (Vogt

1992) at Keck Observatory, Mauna Kea. The echelle spectrograph provided several orders

including the Li _6707 and Ha regions. A 0.861 arcsecond wide slit was used, providing a

spectral resolution of//,-, 45,000. To determine radial and rotational velocities, an order

in the 6405-6490/_ range containing several strong Fe & Ca lines was chosen. As noted

in Jones, Fischer, & Stauffer (1996), the HIRES spectrograph is very stable; wavelength

calibration was derived from thorium-argon comparison lamp exposures obtained at the

beginning and end of each night. The flatfielded, extracted spectra were wavelength calibrated using the onedspec package in IRAF; radial velocity measures were subsequently

obtained using the xcsao task (version 1.6) in the IRAF rvsao package.

5
To serve as a check, the IC 4665 member P39 was also observedon each night that

observations were obtained for the fainter M dwarf candidates. Radial velocity analysis

followed the sameproceduresas in Prosser& Giampapa (1994). Due to the differencesin

spectral type between P39 and the fainter M dwarfs, standard stars appropriate for the

target star(s) wereadopted: for P39the standard vr_ stars were GL33, ttD3765, HD29587,

and HD154417, while for the fainter M dwarfs the standards used were GL176, GL654,

GL873. The HD stars were taken from Latham & Stefanik (1991), using the CfA radial

velocities, while the Gliese stars are from Marcy & Benitz (1989). Rotation velocities were

derived using the same procedure as in Prosser & Giampapa (1994); due to the differences

in spectral types, two calibration template spectra were used - a spectrum of HR 996

(GSV, Soderblom et al. 1989) for P39 and a spectrum of GL654 (M3.5V) for the Ha

emission M dwarfs.

Equivalent width measures from the echelle spectra for the M dwarfs that exhibit Ha

emission are listed in Table 4. We estimate an uncertainty of about 10_ in the equivalent

width measures of the Ha emission line stength, as based primarily on the accuracy of the

fitting of the local continuum. This method of error estimate follows the approach outlined

by Wordea, Schneeberger & Giampapa (1981) in their measurements of the strengths of

chromospheric lines in M dwarfs. Other factors that can affect the uncertainty include the

v sini of the star and intrinsic variability arising from stellar activity.

A summary of the vr,,d/V sini measures from the Keck HIRES observations is also

given in Table 4. Members of IC 4665 are known to have radial velocities of ,-, -13 km s-1

(Prosser & Giampapa 1994, Crampton et al. 1976, MorreU & Abt 1991). The Vrad measures

for P39 are consistent with previous measures (Prosser & Giampapa 1994), though a

smaller v sin i value is found from these higher resolution, higher signal-to-noise CCD

spectra. Among the I-In emission M dwarfs observed, most have Vrad measures consistent

with membership, with nonmembership indicated only for P261. From comparison of the

radial velocities derived for a given target star using different radial velocity standards

(which were observed throughout each night), the accuracy of the v_d measures for the IC

4665 stars is on the order of +0.3 km s -1, though the error increases to -4-0.6 km s -1 for

P257 - presumably as a result of its moderate rotation. Following Prosser & Giampapa

(1994), we estimate the accuracy of our v sin i measurements to be about 2-3 km s -1 in

the range 10 < v sin i < 30 km s -1, and increasing to ,-, 5 kin s-1 for higher values of v

sin i.

In addition to the Keck echelle observations, we also list in Table 4 the v,:_.d/V sin i

measures obtained for P97 and Pl17 using the KPNO 4m echelle and UV Camera with

79-63 ° grating to yield .-, 0.38_ (2 pix.) resolution spectra. For P97, the Vr_.a/vsini

6
measures were derived from a region of spectrum around 6400]t, while for Pl17 the ?3ra d

measurement is derived from a portion of spectrum containing the Ha absorption feature

- no reliable v sin/ measurement could be obtained for Pl17 due to low signal-to-noise.

As with the Keck observations, five/six radial velocity standard star observations each

night were utilized to obtain an average l)ra d measure for the IC 4665 stars. The resulting

radial velocities for P97 & P117 clearly indicate nonmembership; from comparison using

different radial velocity standards, the accuracy of the KPNO Vrad measure for P97 is :t:2

km/s, while that for Pl17 using the Ha feature is -t-6 km/s. Both P97 and Pl17 show

I-Ia in absorption. The v sin i estimate for P97 was determined from comparison with a

calibration template spectrum of the radial velocity standard HD3765 (KS).

The Ha profiles of the M dwarfs observed by Keck are illustrated in Fig. 1. The

moderate rotation observed for P257 is reflected in its broadened Ha emission profile. The

strong Ha emission observed in P230 exhibits a 'pedestal' or broadened wing component

in its profile that is not evident in the other narrow-lined stars. This phenomenon has

been observed in similar high-dispersion echelle observations of young stars in the Pleiades

and IC2391/IC2602 clusters (Stauffer et al. 1997, Jones, Fischer & Stauffer 1996). The

likely origin of this feature has been attributed to flare or microflare events. In fact, the

Ha equivalent width from the Keck echelle observation is twice as large as the previous

measure from low-dispersion spectra (see Prosser 1993; his Table 5), in accordance with

the flare hypothesis. However, it is also possible that the additional pedestal seen in the

echelle data may not have been resolved from the surrounding spectrum in the earlier

lower dispersion spectrum, resulting in the lower equivalent measurement given by Prosser

(1993).

The Keck echelle spectra were also examined for the presence of the Li ,k6707 feature.

For the G0V star P39, an obvious Li absorption line is seen, with an equivalent width on

the order of 170m_. Among the dMe stars, no significant feature is seen and an upper

limit of 1150 m./t to the EW is estimated, with a higher upper limit of 300 m]k assigned to

P261 due, in part, to the lower S/N level for this star. The lack of detectable lithium and

the estimated upper limits are consistent with the results obtained for similar low mass M

dwarf members of the Pleiades (Jones, Fischer & Stauffer 1996).
3. X-RAY OBSERVATIONS

3.1. ROSAT HRI Data

Our present HRI observation of IC 4665 consists of a 76.9 ksec exposure obtained at

a position near the cluster center; details of the observation are given in Table 5. The

location of the HRI field was chosen so as to include IC 4665 members extending over a

range of mass, particularly solar-type members from Prosser & Giampapa (1994) and a

selection of the Ha emission M dwarfs found in the cluster region (see Table 5 and Fig. 5

in Prosser 1993). The data were analyzed using EXSAS routines running under MIDAS;

the two separate observations obtained in 1995 and 1996 were combined directly into one

observation since no significant offset in X-ray positions was found between the separate

A & B observations.

LOCAL & MAP detections in EXSAS were run to create a provisional source list.

This list was used as an input list to the maximum likelihood (ML) detection algorithm

(Cruddace et al. 1988) to obtain a list of sources having ML>5. The source detections

provided by EXSAS were visually examined in order to establish the reality of weaker

sources. While most sources are identifiable at the ML>7 level, some sources with 5 < ML

< 7 were retained, particularly in cases where a star in our optical catalogue was known

to be at the X-ray position.

In Table 6 we list those catalogued stars which were detected in X-rays in our observation. The table lists star name, available optical information, the (a,$) offsets between

X-ray and optical positions (in the sense X-ray - optical), followed by the corresponding

X-ray information, luminosities and additional notes. The X-ray and optical positions

all agree to within 15 arcsec or better. There are 253 catalogued stars in our IC 4665

database which fall within the HRI observation. For the ,,_ 43 X-ray sources detected,

one may estimate the number of chance alignments or matches between X-ray sources

and catalogued stars by multiplying the number of X-ray sources by the fraction of catalogued stars falling within a given distance from the X-ray position (following Randich et

al. 1995). For matches within 15 arcsec from the X-ray position, the number of chance

alignments is on the order of ,,_ 1.5. Most of the stars in Table 6 lie within 10 arcsec from

the X-ray position; the corresponding number of chance alignments for this smaller radius

is estimated to be less than one.

In order to convert the count-rates to observed X-ray fluxes and, eventually, X-ray

8
luminosities, we utilized the cluster parametersin Table 1 and a conversionfactor (CF)

corresponding to a hydrogen column density of log NH (cm -2) = 21.0 (Bohlin et al. 1978)

along with the assumption of a Raymond-Smith plasma of temperature kT = 1.0 keV.

The corresponding CF, based on the HRI calibration data given in Table 11.4 of the RMD

(ROSAT Mission Description), is CF = 7.63 × 10 -11 erg cm -2 cnt -1. The assumption

of kT = 1.0 keV is based on the results of Gagn6, Caillault & Stauffer (1995) who found

that kT = 1.0 keV was an accurate representation of the composite PSPC spectra they

constructed from their observations of young, active dwarf stars in the Pleiades. In particular, Gagn6 et al. (1995) find that the emission measure of the hot (,-, 107 K) component

exceeds the emission measure of the soft component in most cases. This hot component is

characterized by kT ,,, 1.0 keV in the composite spectra of G-M dwarfs constructed by

Gagn6 et aI., and slightly less than this value for F dwarfs. In addition, the higher extinction (,,- 5 times higher than that of the Pleiades) associated with the greater distance of

IC 4665 will tend to weight the present X-ray observations toward the harder component.

In the particular case of the B-star detections, we considered the possibility of UV

contamination arising from the UV sensitivity of the ROSAT HRI. We adopted the scaling

given by Schmitt et al. (1993) to find, at its maximum, a UV count rate of ,,_ 10 -4 ct

s-1 for K62 and less for the other detected high mass members. In particular, the ratio

between the observed HRI count rate and the calculated UV rate ranges between ,,_ 3 - 62.

We therefore conclude that the observed HRI counts in the detected high mass members

are dominated by X-ray emission rather than photospheric UV radiation.

In order to calculate the parameter (Lx/Lbot), bolomet tic corrections must be applied

for the early-type and the later-type stars. Bolometric corrections (BCv) were obtained

from Harris (1963, his Table 7) for stars with (B-V) < 1.20 (approximately spectral type

K5 and earlier), while BCI values were obtained from Monet et al. (1992, eqn. 6) for latertype stars. Solar values of log Lbot =33.583 and Mbol =4.76 (Allen 1976) were assumed.

For nonmember stars in Table 6, the Lx, Lbol, and B-Vo values are provided only for

comparison purposes as if they were really cluster members. Finally, we assume uniform

extinction across our HRI field.

For those catalogued stars considered to be members and candidate members (Y/Y?)

or undecided (?), and which were not detected in X-rays (i.e., ML < 7), upper limits in the

flux corresponding to a confidence level of 99.7% were determined; these values are listed

in Table 7. For those stars in Table 6 with ML < 7, we provide upper limits in Table 7

should these sources later prove to be spurious. For stars with no measured B-V color, the

B-Vo value in Table 7 has been estimated using the observed V-IK color and the B-V

vs. V--IK relation in Prosser et al. (1996). Because of the increasing measurement errors
in the HRI field for off-axis positions greater than 18 arcminutes, we regard the upper

limits in Table 7 for such stars as less reliable. Therefore, in the discussion of these results,

we will utilize only those upper limits determined for objects with off-axis positions of less

than 18 arcminutes.

4. RESULTS AND DISCUSSION

We exhibit in Fig. 2 log L, vs intrinsic (B-V) color for detections of members and

candidate members (i.e., "Y',"Y?" & "?') from Table 6, and upper limits from Table 7

for members (Y) and candidate members (Y?). Inspection of Fig. 2 reveals stellar X-ray

sources spanning the color range -0.17 < (B - Y)o _< +1.63 (,,_ B3- M3). The general

features of Fig. 2 also include a rise in Lx beginning near F0 which is widely interpreted

as evidence for the onset of outer convection zones and dynamo activity, analogous to that

seen in the Sun. The observed X-ray emission in the cool stars reaches a maximum at ,,_

GO and then declines toward later spectral types. The decline in Lx toward lower masses

is a general feature of the X-ray properties of late-type stars, both in the field (Fleming

et al. 1993) and in young clusters (Caillault 1996 and references therein). Fleming (1988)

and Fleming et al. (1993) attribute this trend to the decline in stellar surface area, at least

in the case of the most X-ray active stars where the surfaces ir/ay be entirely covered, or

"saturated", with magnetic regions that are the sources of X-ray emission.

In Fig. 3 we show log (Lx/Lbot) vs (B - V)o color for the same IC 4665 stars given

in Fig. 2. The parameter (Lx/Lbot) is a normalized quantity that permits intercomparison

among different spectral types. It is also interpreted as representative of the efficiency of

coronal heating in late-type stars (Mullan 1984).

The general features in Fig. 3 include a rise from (L,/Lbot) "" 10 -7 among the hot

stars to a maximum level of (Lx/Lbot) _ 10 -3, possibly as early as (B - V)0 "-_ 0.7 -

0.8, mad certainly by (B - V)0 ,,- 0.9. The latter has been referred to as the so-called

'saturation level' by a number of authors in studies of the X-ray properties of cluster

members (Stauffer et al. 1994 [Pleiades]; Randich et al. 1996 [a Per]; Stauffer et al. 1997

[IC 2602/IC 2391]). We note that prior studies of active, short-period binary systems

in the field revealed saturation in the fluxes of far UV transition region lines (Vilhu &

Rucinski 1983; also see Vilhu 1984). The level of saturation is not due to radiative transfer

effects; rather, it is considered a real property that is a manifestation of the underlying

dynamo mechanism. The structures that give rise to the X-ray emission are likely compact

regions (e.g., Giampapa et al. 1996; Schmitt & Kiirster 1993) since extended structures

would be characterized by a volume filling factor that could yield values of (Lx/Lbot) that

substantially exceed the observed maximum.

10
4.1. The High Mass Stars

Our HRI field contains several high-mass members of IC 4665, including several detected B stars. A summary of the observed properties of the bright cluster members is

given by Prosser (1993; see his Table 6). Using the membership information from Prosser

(1993), we summarize in Table 8 the spectral type and X-ray data for the B and A star

members and candidate members. Those stars not contained within the present HRI observation are noted in Table 8 with the term 'beyond field'. The observed X-ray emission

among IC 4665 stars in Table 8 is comparable to that seen in similar high-mass stars in

other clusters. While many chemically peculiar (CP) stars may be expected to have detectable levels of X-ray emission (see the discussion in Jeffries et al. 1997), the Hg-Mn star

K76A is a non-magnetic CP star and is not necessarily expected to be an intrinsic X-ray

source. Additional A-type members (Prosser 1993, his Table 6) not listed in Table 8 are

generally outside of the field of the present HRI observation, though an upper limit has

been determined for K66 (see Table 7).

A dispersion of ,,_ 1.5 dex in log (Lx/Lbot) is seen for stars hotter than A0. Previous Xray studies of young clusters revealed similar -_ 1.5 dex dispersions in the X-ray properties

among the high mass members (e.g, Randich et al. 1995 [IC2602]; Patten & Simon 1993 [IC

2391]). We note, however, that the lower sensitivity limit of log (Lx/Lbot) "_ -7.6 for the

high mass stars in the (more nearby) cluster IC 2391 would suggest that the actual range

in log Lx among high mass members in IC 4665 is greater than indicated by our present

data-set. This is in contrast to prior investigations that have quantitatively characterized

the X-ray emission in young B stars by Lx _ 10 -v Lbot (Schmitt et al. 1993). In the field,

however, the X-ray properties of B stars are quite different. In particular, Cassinelli et al.

(1994) find that (Lx/Lbot) declines sharply from ,-_ 10 -7 to ,-_ 10 -9 in the nearby, mainsequence B stars. Furthermore, the detected X-ray emission is characterized by relatively

low source temperatures as inferred from the soft PSPC spectra. As Cassinelli et al. point

out, B stars such as these would be difficult to detect at hydrogen column densities of

log Nh > 20 cm -2. Hence, the B star sources in IC 4665 (where log Nh = 21 cm -2)

must have relatively harder spectra in addition to the greater range in X-ray emission as

compared to their counterparts in the field.

The X-ray emission observed in B stars is generally attributed to either intrinsic emission arising from shock heating of the surrounding medium by a high velocity, radiatively

driven wind or to the presence of an active, late-type companion. For example, Meurs

et al. (1992) suggest that the high L,/Lbol ratios deduced for some of the late-B stars in

the ROSAT All-Sky Survey are due to the presence of late-type, binary companions.

11
We see from inspection of Fig. 2, combined with the published X-ray data-sets for

the Pleiades and a Per (see Fig. 4), that the maximum, non-flare X-ray luminosities of

the late-type members of young clusters is log L, ._ 30.5. Adopting this value leads to an

estimate of log (Lx/Lbot) _ -6.0 for mid-B stars if the X-ray emission is due entirely to

a late-type companion. It would therefore seem unlikely that the observed X-ray emission

in high mass stars at levels exceeding log Lx/Lbot ": -6.0 is due to a late-type companion.

Moreover, the Morrell & Abt (1991) survey of binaries in IC 4665 shows that the mass

ratio of binary systems among the high mass stars is of order unity. Therefore, it would

seem unlikely that the presence of an active, late-type companion is the origin of strong

X-ray emission in the high mass stars.

According to Morrell & Abt (1991), the detected B stars in Table 8 exhibit constant

radial velocities, with the exception of K49 which is a SB1 system. Since the secondary

is not visible in the spectrum, it is presumably a low mass star. We can utilize the more

complete X-ray data-set available for the Pleiades (since IC 4665 and the Pleiades are

approximately the same age) and the observed value of log L, = 29.28 for K49 to infer

that the companion is a late F - early K dwarf (see Fig. 12 in Stauffer et aJ. 1994).

The remaining, ostensibly single, high mass X-ray sources include two objects, K73

and K58, with log (Lx/Lbot) > -6.0. These stars are rapid rotators with v sini > 200 km

s -x. Thus, rotationally enhanced winds, such as those discussed by Poe & Friend (1986)

and Marlborough (1987), may play a role in the ultimate origin of the relatively strong

X-ray emission exhibited by these objects. The relatively slowly rotating B star, K62, is

characterized by a low value of log (Lx/Lbol) = -7.05. However, K64 with a similarly high

value of v sin/exhibits a more intermediate level of X-ray emission among the high mass

stars of IC 4665 with log (Lx/Lt, ot) = -6.78. Interestingly, the field star sample considered

by Cassinelli et al. includes rapid rotators (v sin/ > 200 km s-1) with quite low values

of relative X-ray emission, i.e., log (L,/Lbot) _ -7 - -9. Apparently, there is a qualitative

difference in the X-ray properties between the B stars in young clusters and B stars in

the field. In brief summary, we attribute the strong X-ray emission at levels exceeding

,.- -6.0 in log (L_:/Lbot) to an intrinsic property of the star and its interaction with the

surrounding medium, perhaps through a rotationally enhanced wind.

The relatively lower number of detections in the A0 - F0 range is consistent with

the claimed absence of X-ray emission from stars without either outer convection zones or

high velocity winds. In particular, the A stars appear to be X-ray quiet (Schmitt 1997);

X-ray emission that has been detected in early-to mid-A stars is generally attributed to

the presence of a late-type companion (e.g., Schmitt et al. 1985; Schmitt &: Kiirster 1993).

12
We note that while a radial velocity survey of the B star members of IC 4665 has been

performed (Morrell _ Abt 1991), a comparable survey of the A-type members has not

yet been conducted. Hence, our understanding of binarity among the cluster A stars is

limited. Apparently single, late-A stars in the field have been detected as X-ray sources

(Simon et al. 1995). These objects likely represent the demarcation on the main sequence

for the onset of outer convection (e.g., Wolff et al. 1986).

4.2. The Low Mass Stars

An intercomparison with other young clusters can yield further insight on the X-ray

properties of the cool stars in IC 4665. In Fig. 4 we plot log Lx vs B-Vo for detected

sources that are members of the Pleiades (Stauffer et al. 1994) and a Per (Randich et al.

1996), respectively. It is evident from Fig. 4 that the detected sources in IC 4665 follow

the same trends with color as the other two clusters. A notable exception occurs at

(B - V)o "_ 1.5 where several objects in IC 4665 - P232, P257, P270, and P315 - exhibit

levels of emission that are greater than that seen in either the Pleiades or a Per at this

color. These particular objects may therefore not be members of IC 4665. Rather, they

may be foreground dMe stars with true values of Lx that are lower than would be deduced

by assuming these objects are at the distance of IC 4665. In fact, P232, P257, P270

and P315 all lie ,,_ 1 mag above what may be considered the single star sequence in the

V vs (V - Ig) color-magnitude diagram for IC 4665 (see Prosser 1993; his Fig. 5). This

serves as an example of how X-ray observations can reveal possible nonmembers among

photometrically selected candidates. Further observations are needed to finally confirm

the membership status of these specific stars.

We have also compared the distribution of X-ray emission with color between IC 4665

and that for the even younger clusters IC 2391/2602 [Patten & Simon 1993 (IC 2391);

Ra_dich et al. 1995 (IC 2602)], each with an age ,,_ 35 Myr. As in the case of the

comparison with the Pleiades and a Per, the observed trends of log Lx with color are

similar for both clusters. Again, near (B - V)o "_ 1.5, the suspected non-members in the

IC 4665 field appear to exhibit higher levels of X-ray emission than the late-type members

of IC 2391/2602.

In Fig. 5 we include for comparative purposes the detected sources in IC 4665 along

with the detections in a Per and the Pleiades, respectively. The values of log (L_/Lboi)

for a Per are taken from Randich et al. (1996) while those for the Pleiades are given by

Stauffer et al. (1994). 4 While the IC 4665 observations axe limited due to fewer stars,

* We note in passing that the published upper limit of log(Lx/Lbol) < --4.91 for HII

13
the distributions indicate that the log Lx/Lbol --_ --3 saturation level is reached at a

similar color and spectral type in both the Pleiades and IC 4665. In IC 4665, the onset

of saturation occurs in the color range (B - V)o _- 0.7 - 0.8. which, in turn, is similar

to the Pleiades where saturation begins at (B - V)o --- 0.75. Previous investigations have

shown that the color at which this maximum first occurs increases with increasing age of

the cluster. For comparison we list in Table 9 the approximate colors where the onset

of the saturation level of log (Lx/Lbol) ,,_ -3 has been observed in other clusters. The

increasingly redder color (i.e., lower mass) at which saturation is attained in progressively

older clusters is regarded as a manifestation of the mass-dependent rate for rotational

spin down. In particular, the time scale for spin down via magnetic braking appears to

increase with decreasing mass. Thus, the observed X-ray activity in IC 4665 provides

further evidence that its age is similar to that of the Pleiades, as postulated by Prosser &

Giampapa (1994).

Inspection of Fig 5 reveals that the observed X-ray activity levels of the dwarf M stars

in IC 4665, the Pleiades and a Per extend above the log(Lx/Lbol) ,'_ -3 saturation level

seen for G and K dwarfs. The higher, relative levels of coronal activity in the M dwarfs

could be due to: 1) ROSAT is simply detecting those M dwarfs which are flare-active at

the time of observation, or 2) the -3 saturation level may not apply to the M dwarfs.

Currently, no specific calculation has indicated why the activity among G and K dwarfs

reaches a plateau at log Lx/Lbol -- --3, and it may be that the changing physical structure

in lower mass stars can lead to higher levels of coronal activity.

4.2.1. The Correlation Of Coronal Activity With Rotation

We illustrate in Fig. 6(a) the correlation between X-ray activity and rotation (v sin/)

in the late-type stars. In Fig. 6(a) we see that log L_ exhibits an increasing trend with

increasing rotational velocity. We also include for comparative purposes the relation for

L_ vs rotation from Pallavicini et al. (1981) inferred for F - M field stars, namely, log L_

= 27 + 2 log i v sin i). The empirical relationship between the normalized quantity log

(L_/Lboz) and rotational velocity is illustrated in Fig. 6(b). In view of the limited data,

we do not attempt a functional fit. We note that the saturation plateau at (L_/Lbot) "_

10 -3 begins to appear in stars with projected rotational velocities as low as 10 km s-1 , and

2462 given by Stauffer et al. is incorrect and apparently results from an erroneous estimate

of Lbol- A value of log(Lx/Lbol) < -4.35 for HII 2462 has been used here. This correction

becomes important for the accurate assessment of both the range and the lower limit of

X-ray activity observed in the solar-type stars in the Pleiades.

14
definitively by v sini ,,_ 20 km s -1. The so-called saturation itself is presumably related

to a maximum efficiency in dynamo action that results in a filling factor of X-ray emitting

structures on the stellar surface of order unity. In contrast to the absolute level of emission

(L,), no further increase in (Lx/Lboi) occurs at rotational velocities exceeding 20 km s -1.

We note, parenthetically, that the cluster member P27 (B - Vo = 0.59) is not detected at

a relatively high upper limit of log L, = 30.00 or log (Lx/Lbot) = -3.73, although it has a

high rotational velocity of v sin i "_ 33 km s -1. The measured X-ray flux limit may have

been affected by its large off-axis position and/or otherwise influenced in some unknown

manner by the fact that P27 is considered to be a SB1 system (Prosser & Giampapa 1994).

We further examine the relationship between coronal X-ray emission in IC 4665 and

dynamo parameters in Fig. 7(a) by utilizing the Rossby number (R0). This hydrodynamic

parameter incorporates both rotation and convection, the key processes in dynamo action,

in the form of a ratio between rotation period (P,.ot) and convective turnover time (vc). We

compute the Rossby number for the G-M type cluster members for which we have rotation

periods by adopting the relation between convective turnover time and (B - V) color given

by Noyes et al. (1984). We supplement these data with upper limits in rotation periods

for those late-type stars in IC 4665 for which we have values of v sin i. Stellar radii were

estimated from (B - V)o colors and data given by Allen (1976). We include in Fig. 7(a)

the linear relation between log (L,/Lbot) and R0 found for stars in a Per by Randich et al.

(1996). Our data appear consistent with a declining trend in normalized X-ray emission

toward larger Rossby numbers. However, the few data combined with the upper limits in

log (Lx/Lbol) at the greatest extent of our range in Ro precludes any further quantitative

characterization. Since the Rossby number is a model-dependent quantity, we illustrate in

Fig. 7(b) the relation between observed quantities, namely, log (Lx/Lbot) and the rotation

period. The data for IC 4665 are supplemented with estimates of rotation periods for stars

with v sin i values only, as described above for Fig. 7(a). We include for comparative

purposes the results given by Prosser & Grankin (1997 and references therein) for G/Ktype members of a Per. As observed by Randich et al., (L,/Lbot) appears to decline at

rotation periods less than ,.- 10-12 hrs in a Per. Currently, there are insufficient data to

know whether this same trend is also present in IC 4665. We note that the early F dwarf,

K53, appears underluminous in (L,/Lt, ot) in Fig. 7(b). By contrast, the location of K53 in

Fig. 7(a) is somewhat more consistent with previous findings that the X-ray emission levels

in early F stars correlate well with Rossby number while no correlation is seen between

rotation alone, and Lx or (Lx/Lbot) for these objects with thin convection zones (Schmitt

et al. 1985; Randich et al. 1996)

15
4.2.2. Chromospheric and coronal emission

Finally, we illustrate in Fig. 8 the observed correlation between Ha equivalent width

and coronal X-ray emission for the late-type stars in IC 4665 and the Pleiades. The Ha

equivalent widths are taken from Table 4 and Prosser (1993) for IC 4665 while the Ha

equivalent widths for Pleiades members are given by Stauffer & Hartmann (1987). The

basic form of this empirical relationship is the same in both clusters, i.e., the Ha absorption

equivalent width fills in with (presumably) chromospheric emission, becoming weaker with

increasing coronal activity. At log (Lx/Lbot) "_ -3.5, Ha appears in emission in the spectra

of those IC 4665 stars for which we have both Ha and X-ray data. At the saturation level

of log (Lx/Lbot) "_ -3.0, there is a large scatter in the Ha emission equivalent width.

The origin of the scatter in Ha strength may be of interest in the context of coronal

and chromospheric structure. In particular, the saturation in X-ray emission is likely

related to the compact nature of the X-ray structures, as suggested for the M dwarfs by

Giampapa et al. (1996) and earlier, active (G) dwarfs by the results of Schmitt _ Kiirster

(1993). Hence, the upper limit or saturation in X-ray emission would be a function of

stellar surface area, as pointed out by Fleming (1988). By contrast, the large scatter in

chromospheric activity at the saturation level near log (L,/Lbot) "-" -3 suggests that the

region of Ha emission is not similarly constrained. That is, the structures which are the

principal contributors to Ha emission may be more geometrically extended, in contrast to

the relatively compact X-ray emitting regions. Flare activity in Ha would contribute to

the scatter as well.

4.3. X-ray Variability

The two ROSAT observations of this HRI field, IC4665-A/IC4665-B, enables us to

investigate the possibility of long-term (,-_ 1 year) X-ray variability in sources that are

detected in at least one of the two observations. We identified potentially variable Xray sources by comparing the difference in X-ray counts between objects detected in both

epochs to the quadratic sum of the error of the individual measurements, following Fleming

et al. (1995). In the case of sources detected in one epoch but not the other, we compared

the difference between the counts in the detection and the upper limit with the error in

the detected source. Only one source, K56, was detected in both observations and was

found to have a difference at the 3a level; the count rates, rate errors, and ML values are

listed in Table 10.

We also list in Table 10 sources that are detected only in one observation where the

difference between the detected count rate and the upper limit from the second observation

16
exceeded the combined error. The 'R' X-ray sources in Table 10 will be described in the

next section. We then constructed X-ray light curves for these objects with an adopted bin

size of,-- 900 s. Inspection of the X-ray light curve for the 3a source variable, K56 (A0; nonmember), did not show evidence of short-term variability within each separate observation.

Moreover, we did not find evidence for variability similar to a flare-like outburst within

the light curves of the remaining candidate variable X-ray sources given in Table 10. In

particular, the count-rate appeared uniform to within the errors throughout the observing

period in each epoch.

In the cases of K43, P230, P237, and R6, which were detected in IC4665-A, the count

rates are less than the sensitivity limit we roughly estimate of ,,_ 7 x 10 -4 cnts s-1 for

the shorter-exposure, IC4665-B observation. Hence we would not expect to detect these

objects in IC4665-B. Conversely, P248 and R12 should have been detected in IC4665-A.

Similarly, the M dwarf P232 is strongly detected in the first epoch but not the second. In

these cases, as with K56, we can only conclude that long-term variability on time scales

of ,--, 1 year can occur. We note that the dwarf M stars that may be non-members (§4.2)-

P257, P270, and P315-in addition to P232, are characterized by relatively uniform X-ray

light curves in our observations. However, the example of P232 suggests that these stars

can exhibit long-term excursions in their X-ray emission between "high" and "low" states.

Finally, we examined the light curves for the two B stars with the highest levels

of X-ray emission, namely, K58 and K73. Again, the light curves were unremarkable

and characterized by a uniform count-rate (to within the errors) during the course of

observation in each epoch.

4.4 Additional )(-ray Sources

Besides those X-ray sources associated with catalogued stars (Table 6), there were

some additional X-ray sources identified at the ML > 5 significance level in the HRI

observation. These additional X-ray sources which were not matched to a catalogued star

are given in Table 11. The sources are listed by a running 'R' number, followed by the

X-ray (a,_) coordinates, ML value, off-axis distance, count rate, and error in the count

rate. Although these X-ray sources are not necessarily associated with the IC 4665 cluster,

for comparison purposes the corresponding X-ray luminosity for the distance of IC 4665 is

given.

To identify possible optical counterparts to these X-ray sources, BVI CCD photometry

was obtained during Nov. 1994 and Oct. 1995 using the 48 inch telescope of Whipple

Observatory at Mt Hopkins, Arizona. For each night's observations, a set of extinction

17
stars were monitored, along with standard stars from Landolt (1992), and somea Per

members with photometry from Stauffer et al. (1985, 1989) and Prosser (1992). Each

night's instrumental magnitudes were transformed onto a standard BV (Johnson) and I

(Kron) system. Standards with Cousins-system V-I colors were transformed to the Kron

V-I system using the transformation equations of Bessell & Weis (1987). The standards

were also used in determining the nightly aperture corrections applied to the photometry.

In Fig. 9, we provide finding charts at the X-ray position for the 'R' X-ray sources in

Table 11. Charts are approximately 3 x 3 arcmin, North up and East to the left, and are

extracted from a digitized scan of the Palomar Schmidt XE plate using the GASP software

at STScI. The circles are centered at the position of the X-ray detection and have a radius

of ,_ 30 arcsec. For each X-ray source, typically several stars within ,,- 30 arcsec of the X-ray

position were measured for photometry. As seen from Table 6, most optical counterparts

- particularly for the stronger X-ray sources - may be expected to be found within ,,_ 15

arcsec from the X-ray position. Additional counterparts out to ,-_ 30 arcsec radius in Fig.

9 were measured as either secondary targets or for comparison purposes. Following the

identifications in Fig. 9, the BVI photometry for the possible optical counterparts to the

'R' X-ray sources is given in Table 12. No photometry is available for stars in the region of

sources R2 and R5. VI CCD images obtained during non-photometric conditions for us by

R. Schild indicate that a faint red object is located near or at the X-ray position of R5. For

the strong X-ray source R14, we note the presence of a faint red companion to the bright

star in Table 11 with a separation of only 3 arcsec. An MK spectral type of K5-K7 V has

been derived for the brighter star from a spectrum obtained with the FAST spectrograph

(section 2.1). Without additional data however, we cannot precisely determine whether

the bright star or its faint red companion is the source of X-ray emission.

5. SUMMARY

Our joint X-ray and optical investigation of IC 4665 represents a further contribution

to the recent ROSAT studies of young clusters. At a galactic latitude of +17 degrees, the

location of IC 4665 suggests that its formation history in the galactic environment may be

uniquely different from other low galactic latitude clusters. Our ROSAT investigation of

IC 4665 is complemented by a significant and expanding optical data base which has been

supplemented herein with new Keck HIRES spectra of several faint, low mass members. In

addition, photometry of possible optical counterparts to previously uncatalogued sources

is given in this investigation.

The comparison of the detected sources and upper limits for members and candidate

18
members of IC 4665 with that of other clusters, including the Pleiades, a Per, and IC

2391/2602, reveals considerable overlap and similarities in the distribution of X-ray properties with (B - V)o color. In all clusters, a maximum or "saturation" level in normalized

X-ray luminosity is attained at log (Lz/Lbot) " -3 in the low mass stars.

More specifically, the detected sources in IC 4665 follow the same trends in Lx and

(Lx/Lbot) with (B - V) color as seen in other young clusters such as a Per, the Pleiades,

and IC 2391/2602. These trends include a rise in both Lx and (Lx/Lb,,t) in the late-type

stars beginning near ,,_ F0. This trend is associated with the onset of solar-like dynamo

activity. A maximum or "saturation limit" of log (L,/Lbot) "- -3 is attained in IC 4665

in the color range (B - V)0 "_ 0.70 - 0.8 which, in turn, is similar to that found in

the Pleiades (i.e., B-Vo "_ 0.75). We view this as further evidence that the ages of IC

4665 and the Pleiades are approximately the same, as concluded by Prosser &: Giampapa

(1994).

The X-ray observations suggest that several photometrically selected M dwarfs with

Ha emission are actually non-members. However, the relatively enhanced X-ray emission

in these stars, compared to those in the Pleiades and a Per at the same color, could also

result from flare activity--though no evidence for flare-like outbursts was evident in an

inspection of the X-ray light curves for both epochs of observation. Nevertheless, the

results for the specific case of P232 imply that long-term (,,- 1 year) variability that is

characterized as a change in activity from a "high" to a "low" state can occur in these

stars (independent of the question of membership).

The empirical correlation of rotation and X-ray emission in IC 4665 is similar to that

of other young clusters, at least as seen within the limitations of the available data. The

saturation level is certainly attained at rotational velocities of v sini = 20 km s-1 and, in

IC 4665, possibly by a rotational velocity as low as ,-_ 10 km s -1. Finally, we note that

in the emission line stars there is significant scatter in the Ha emission equivalent widths

at log (L_/Lboz) "_ -3. Apparently, Ha is not similarly "saturated" but may instead

be dominated by emission from geometrically extended regions while the X-ray emission

arises primarily from surface regions that are relatively compact (Giampapa et al. 1996).

Conversely, the Keck HIRES observations of Ha emission line profiles in objects such as

P230 imply that flaring, which is more prominent in Ha, could be the source of the large

scatter in the strength of this feature at high normalized X-ray luminosities.

Among the 5 high mass B stars that are X-ray sources, one source is a spectroscopic

binary (K49). The level of emission for K49 is consistent with that of a late F - early K

secondary though intrinsic X-ray emission in the B star primary is by no means excluded.

The remaining sources are not radial velocity variables and are presumed to be single B

19
stars. Moreover, we argue that for values of (Lr/Lbot) > 10 -6, it is unlikely that the

X-ray emission is due to a late-type companion. Instead, X-ray activity in high mass stars

is attributed to shock heating of the ambient medium by a radiatively driven wind that

may also be enhanced by rotation.

The similarity of X-ray properties between IC 4665 and other young (_ 30 - 100 Myr)

clusters suggests that formation environment or formation history of open clusters is not

reflected in the nature of stellar coronae, at least for clusters with solar-like metallicities.

We therefore infer that, in clusters with ages < 100 Myr, the distributions of rotational

velocities with color for the low mass stars later than about (B - V)o = 0.6 are similar.

This study was supported by NASA Grant Nos. NAG5-3238 (TAF); NAGW-2698

and NAG5-2203 (CFP); and an interagency transfer from NASA to the NOAO (MSG).

The FAST spectra in IC 4665 were obtained by Perry Berlind during service observing

and we gratefully acknowledge his diligence in obtaining the requisite observations for this

program. The W.M. Keck Observatory is operated as a scientific partnership between the

California Institute of Technology and the University of California. It was made possible

by the generous financial support of the W.M. Keck Foundation.

2O
References

Allain, S., Bouvier, J., Prosser, C., Marschall, L.A., & Laaksonen, B.D., 1996, A&A, 305,

498

Allen, C.W. 1976, Astrophysical Quantities, 3rd ed. (Athlone, London)

Alter, G., BalAzs, B., & Ruprecht, J. 1970, Catalogue of Star Clusters and Associations,

2nd Ed. (Akad6miai Kiad6, Budapest)

Abt, H.A., & Levato, H. 1975, PASP, 87, 849

Bessell, M.S., & Weis, E.W. 1987, PASP, 99, 642

Bohlin, R. C., Savage, B. D., & Drake, J. F. 1978, ApJ 224, 132

CMUault, J.-P. 1996 in Cool Stars, Stellar Systems and the Sun, eds. R. Pallavicini & A.

K. Dupree, ASP Conf. Series vol. 109, p. 325

Cassinelli, J. P., Cohen, D. H., MacFarlane, J. J., Sanders, W. T., & Welsh, B. Y. 1994,

ApJ, 421, 705

Crampton, D., Hill, G., & Fisher, W.A. 1976, ApJ, 204, 502

Cruddace, R., Hasinger, G., & Schmitt, J.H.M.M. 1988, in Proceedings of ESO Conference

on Large Databases, edited by F. Murtagh (ESO, Garching), p 177

Fleming, T. A. 1988, Ph.D. thesis, University of Arizona, Tucson.

Fleming, T. A., Giampapa, M. S., Schmitt, J.H.M.M., & Bookbinder, J. A. 1993, ApJ,

410, 387

Fleming, T. A., Molendi, S., Maccacaro, T., & Wolter, A. 1995, ApJS, 99, 701

Gagn6, M., Caillault, J.-P., & Stauffer, J. R. 1995, ApJ, 450, 217

Giampapa, M. S., Rosner, R., Kashyap, V., Fleming, T. A., Schmitt, J. H. M. M., &:

Bookbinder, J. A. 1996, ApJ, 463, 707

Harris, D.L. 1963, in Basic Astronomical Data, edited by Strand. (University of Chicago,

Chicago), p. 263

Jacoby, G.H., Hunter, D.A., & Christian, C.A. 1984, ApJS, 56, 257

Jeffries, R.D., Thurston, M.R., &: Pye, J.P. 1997, MNRAS, in press

Johnson, H.L., & Morgan, W.W. 1953, ApJ, 117, 313

Jones, B.F., Fischer, D.A., & Stauffer, J.R. 1996, A J, 112, 1562

Keenan, P.C., & McNeil, R.C. 1976, An Atlas of the Spectra of the Cooler Stars (Ohio

State University Press, Columbus)

Keenan, P.C., &: Pitts, R.E. 1980, ApJS 42, 541

Kopff, E. 1943, ANac, 274, 69

Landolt, A. 1992, A J, 104, 340

Latham, D.W., 8z Stefanik, R.P. 1991, in Reports in Astronomy, IAU Transactions XXIB

(Kluwer, Dordrecht), p. 269 (CfA preprint No. 3316)

21
Marcy, G.W. & Benitz, K.J. 1989,ApJ, 344,441

Marlborough, J. W. 1987,in Physicsof Be Stars, eds. A. Slettebak & T. P. Snow, p. 316

(Cambridge: Cambridge University Press)

Martltn, E.L. &: Montes, D. 1997 A&A 318, 805

Meurs, E. J. A. et al. 1992, A&A, 265, L41

Monet, D.G., Dab_n, C.C., Vrba, F.J., Harris, H.C., Peir, J.R., Luginbuhl, C.B, & Ables

H.D. 1992, AJ 103, 638

Morgan, W.W., Keenan, P.C., and Kellman, E. 1943, An Atlas of Stellar Spectra (University of Chicago Press, Chicago)

Morrell, N., & Abt, H.A., 1991, ApJ, 378, 157

Mullah, D. J. 1984, ApJ, 282, 603

Noyes, R. W., Hartmann, L. W., Baliunas, S. L., Duncan, D. K., &: Vaughan, A. H. 1984,

ApJ, 279, 763

Pallavicini, R., Golub, L., Rosner, R., Vaiana, G. S., Ayres, T., & Linsky, J. L. 1981, ApJ,

248, 279

Patten, B. M., & Simon, T. 1993, ApJL 415, L123

Poe, C. H., ,_ Friend, D. B. 1986, ApJ, 311,317

Prosser, C.F. 1992, A J, 103, 488

Prosser, C.F. 1993, AJ 105, 1441

Prosser, C.F., & Giampapa, M.S., 1994, AJ 108, 964

Prosser, C.F., K: Grankin, K. 1997, Center for Astrophysics Preprint No. 4539.

Prosser, C.F., Stauffer, J.R., Caillault, J.-P., Balachandran, S., Stern, R.A., & Randich,

S. 1995, AJ 110, 1229

Prosser, C.F., Randich, S.R., Stauffer, J.R., Schmitt, J.H.M.M., 8z Simon, T. 1996, AJ

112, 1570

Randich, S., &Schmitt, J.H.M.M. 1995, A&A 298, 115

Randich, S., Schmitt, J.H.M.M., Prosser, C., & Stauffer, J.R. 1995, A&A 300, 134

Randich, S., Schmitt, J.H.M.M., Prosser, C., & Stauffer, J.R. 1996, A&A 305, 785

Schmitt, J.H.M.M., Golub, L., Harnden, F. R., Jr., Maxson, C. W., Rosner, R., & Vaiana,

G. S. 1985, ApJ, 290, 307

Schmitt, J.H.M.M., Zinnecker, H., Cruddace, R., & Harnden, F. R., Jr. 1993 ApJL, 402,

L13

Schmitt, J.H.M.M. _ Kfirster, M. 1993, Science, 262, 215

Schmitt, J.H.M.M. 1997, A&A, 318, 215

Simon, T., Drake, S. A., & Kim, P. D. 1995, PASP, 107, 1034

Soderblom, D.R., Pendleton, J., & Pallavicini, R. 1989, A J, 97, 539

22
Stauffer, J.R., Hartmann, L.W., Burnham, J.N., and Jones,B.F. 1985,ApJ, 289,247

Stauffer, J. R., & Hartmann, L. W. 1987,ApJ, 318, 337

Stauffer, J.R., Hartmann, L.W., & Jones,B.F. 1989,ApJ, 346, 160

Stauffer, J.R., Caillault, J.-P., Gagn_,M., Prosser,C.F., & Hartmann, L.W. 1994,ApJS

91,625

Stauffer, J.R., Hartmann, L.W., Prosser,C.F., Randich, S., Balachandran,S., Patten, B.,

Simon, T., & Giampapa, M. 1997,ApJ, 479, 776

Stern, R.A., Schmitt, J.H.M.M., & Kahabka, P.T. 1995, ApJ 448, 683

Vogt, S.S. 1992, in ESO Workshop on High Resolution Spectroscopy with the VLT (ESO,

Garching), p. 223

Vilhu, O. & Rucinski, S. 1983, A&A 127, 5

Vilhu, O. 1984, A£:A, 133,117

Wolff, S. C., Boesgaard, A. M., & Simon, T. 1986, ApJ, 310, 360

Worden, S, P., Schneeberger, T. J., &: Giampapa, M. S. 1981, ApJS, 46, 159

23

Figure Captions

Fig. 1. Ha line profiles in the spectra of the dMe stars observed with the Keck

HIRES echelle.

Fig. 2. The distribution of X-ray luminosity with color for members and candidate

members in IC 4665. The solid circles represent detections while the arrows denote upper

limits. Approximate spectral types are indicated.

Fig. 3. The distribution of the normalized quantity log (Lx/Lbot) with B-Vo color

for IC 4665 stars displayed in Fig. 2. A dashed line is drawn at the saturation level log

(L_/Lbot) -_ -3 that has been observed in other young clusters.

Fig. 4. Comparison of the distribution of Lz with (B - V) color of stars in IC 4665

with members of (a) the Pleiades and (b) a Per.

Fig. 5. Comparison of the distributions in Lx/Lbol among IC 4665 stars with (a)

the Pleiades and (b) a Per. The figure indicates that the onset of the saturation level in

solar-type stars occurs at a similar color for IC 4665 and the Pleiades.

Fig. 6. (a) The log L_ vs log (v sin/) distribution for dwarf G/K/M stars in IC 4665

with available v sin i measures. Stars with upper limits in Lx and/or v sin i are indicated

with arrows. The linear relation is discussed in the text. (b) The correlation between

log(Lx/Lbol) and rotation in IC 4665. The saturation level at log(Lx/Lbol) "_ -3.0 is

indicated by the dashed line.

Fig. 7. (a) log (Lx/Lbot) vs. the logarithm of the aossby number (Prot/Vc) for

IC 4665 stars. Upper limits in rotation periods have been estimated for those stars with

measured vsin/ values. Objects with upper limits in either rotation period or normalized X-ray luminosity are indicated with arrows. The dashed line is the approximately

inverse relation inferred by Randich et al. (1996) for a Per stars in the regime where log

(Prot/rc) >_ -1.2. (b) Late-type stars in IC 4665 with either observed rotation periods

or upper limits estimated from v sin i measurements and estimates of the stellar radius .

Stars in a Per with observed rotation periods are also displayed for comparison.

Fig. 8. The correlation between Ha equivalent width and X-ray emission for stars

in IC 4665 and the Pleiades, respectively.

Fig. 9. Finding charts for the additional X-ray sources in Table 11. Fields shown are

3 × 3 arcmin, north up and east to the left. Catalogued stars (K/P) and stars measured

for photometry are marked.

24

TABLE 1. IC4665: General Information

IC4665: (Mel 179, Cr 349)

(a, 5) = (17h 43.8m, +05°44')(1950) .........

(l,b) = (30.61 °, 17.08 °) .....................

age _ 70 Myr (,,_ Pleiades age) ...............

m-Mv __ 8.2, E(B-V)= 0.18, E(V--I)K_ 0.28

d=340pc, z_ 100pc .......................
TABLE2. MKSpectralTypeStandards

Star Spt Star Spt

80 UMa A5V e Boo G8V

0 Cas A7V ¢r Dra K0V

e Cep F0V SAOl1483 K4V

_r Boo F2V 61 CygA K5V

X Dra F6V 61 CygB K7V

Peg F7V 3' Boo A7III

SAOl1547 F8V 7 Her A9III

/_ Com G0V /3 Del F5III

A Ser G0V 84 Her G2III

16 CygA G3V _¢Oph K2III

16 CygB G5V
TABLE3. IC4665MKSpectralTypes

Star Spt Mem Star Spt Mem

K34 F5V Y P44 G0V Y

K37 F2V Y P51 F8V

K38 F3/FSV Y? P60 G8V Y

K48 F3V Y PT1 K0V Y

K53 F2V Y P75 G8V Y

K57 F1V Y P94 >KOV Y

K65 F0:V Y P100 K0V Y

K88 F0:V Y P106 KOIII N

P1 F5V Y P107 G2:V Y

P12 G8V N P108 G8V N

P19 F8V Y P146 K0V Y

P20 F6V N P150 G5V Y

P27 G0V Y P155 K0V Y

P30 G3V N P159 F6V Y

P33 G0V Y P165 K0V Y

P38 G0V Y P166 K0V Y

P39 G0V Y P177 G0V N
TABLE4. Keck/KPNO Echelle Observations

Star V V--IK Mem? Vrad vsini Ha EW UT Date JD -

(km/s) (km/s) (/_) 2449600.0

Keck:

P39 12.93 0.69 Y -14.6 12 (abs)

............ - 13.7 12 (abs)

P230 17.57 2.60 Y -14.6 12 8.25

P248 17.92 2.49 Y -13.8 <10 1.84

P257 17.14 2.60 Y -14.6 52 4.85

P261 19.31 3.18 N -36.9 <10 2.33

P265 17.20 2.29 Y? -10.8 <10 0.54

P270 17.27 2.57 Y? -18.4 <10 2.47

P336 17.37 2.29 Y -14.4 10 1.61

KPNO:

P97 14.93 1.33 N +22 <10 (aDs)

Pll7 14.89 1.05 N +52 ... (aDS)

Sep 23, 1994 18.75217

Sep 24, 1994 19.76371

Sep 23, 1994 18.82164

Sep 23, 1994 18.85182

Sep 24, 1994 19.83039

Sep 24,1994 19.78315

Sep 23, 1994 18.77098

Sep 24, 1994 19.80974

Sep 23, 1994 18.79859

Nov 24, 1994 80.57465

Nov 21, 1994 77.56568
TABLE5. ROSAT HRI Observations of IC4665.

Name ROSAT ID Exp. Time Dates of Observ.

IC4665-A 201827H 43,874 sec

IC4665-B 201827H-1 33,102 sec

MAR. 23-29, 1995

MAR. 18-28, 1996

Nominal pointing position:

(ct,$)--(17h 45m 43.2s, +05°38r24 '') (2000) .............

(l,b)= (30028',÷17010') ...............................
TABLE 6. IC4665 Stars Detected in X-rays

Name Mem V B-V V-IK Spt vsini As A$ ML Offax Cnts Err logLx logLbol log B-Vo Notes

(km s -I) (")(") (arcmin) (s -l) (s -1) Lx/Lbol

K62 Y 6.87 0.00 ... B3IV

K58 Y 7.59 0.01 ... B6V

K49 Y 7.69 0.02 ... B6V

K64 Y 7.35 0.02 ... B5IV

K73 Y 7.13 0.02 ... B5IV

K43 Y 9.08 0.17 ... A1V

K67NS Y 9.6 ...... A2V

K53 Y 11.43 0.51 ... F2V

P38 Y 12.42 0.74 0.77 G0V

P60 Y 13.43 0.88 0.80 G8V

P71 Y 13.68 0.92 0.91 K0V

P100 Y 14.37 1.06 1.10 K0V

P207 ? 15.50 1.24 1.33 ...

P290 ? 14.88 1.41 1.48 ...

P232 ? 15.92 1.61 1.85 ...

P248 Y? 17.92 ... 2.49 ...

P230 Y? 17.57 ... 2.60 M3

P257 Y? 17.14 ... 2.60 ...

50 +2 -7 13.8 9.0 0.0004 0.0001 29.61 36.66 -7.05 -.18

200 +2 -2 592.1 4.8 0.0041 0.0003 30.64 36.28 -5.64 -.17

25 -6 +2 6.5 5.2 0.00018 0.00007 29.28 36.24 -6.96 -.16 SB1

220 +4 -3 12.1 7.0 0.0004 0.0001 29.60 36.38 -6.78 -.16

210 +4 -1 88.9 16.5 0.0032 0.0003 30.53 36.47 -5.94 -.16

185 +0 +7 5.1 6.7 0.0004 0.0001 29.67 35.29 -5.62 -.01

250: +3 -2 50.8 11.4 0.0013 0.0002 30.12 35.02 -4.90 0.06 VB 3"

38: +5 -4 61.1 4.7

67 +3-4 169.4 12.9

13 -11-6 7.4 17.4

17 +3 +0 83.8 11.0

21 +4 +0 82.1 7.8

... +3 -4 29.5 10.0

... +2 -2 45.5 7.1

•.. +3 +0 38.6 9.9

<I0 +2 -3 8.0 8.8

12

52

0.0008 0.0001 29.92 34.22 -4.30 0.33

0.0032 0.0003 30.52 33.81 -3.29 0.56

0.0010 0.0003 30.00 33.44 -3.44 0.70

0.0017 0.0002 30.24 33.35 -3.11 0.74

0.0012 0.0002 30.11 33.11 -3.00 0.88

0.0009 0.0002 29.95 32.76 -2.81 1.06

0.0007 0.0001 29.84 33.10 -3.26 1.23

0.0009 0.0002 29.99 32.82 -2.83 1.43

0.0003 0.0001 29.53 32.28 -2.75 1.54

+2 -3 8.2 6.0 0.00028 0.00009 29.47 32.48 -3.01 1.55

+3 -4 44.0 7.6 0.0008 0.0001 29.95 32.65 -2.70 1.55

P270 Y? 17.27 ... 2.57 M2.5 <10 +2 -4 23.6 8.4 0.0006 0.0001 29.76 32.58 -2.82 1.55

P283 Y? 18.25 ... 2.55 M3 ... -1 -1 5.9 7.2 0.000200.00008 29.32 32.18 -2.86 1.55

P315 Y?. 17.85 ... 2.78 M3 ... +9 -2 11.5 10.2 0.0005 0.0001 29.74 32.45 -2.71 1.58

P237 ? 19.36 ... 3.03 ...... +2 -4 10.5 5.3 0.00029 0.00009 29.49 31.97 -2.48 1.63

K56 N 7.51 0.12 ... A0 ...

K60 N 11.0 ...... G1 ...

P12 N 12.72 0.77 0.80 G8V 70

P20 N 12.24 0.46 0.60 F6V ...

P30 N 12.68 0.61 0.65 G3V ...

P108 N 13.75 0.78 0.77 G8V ...

÷0 -1 465.1 16.3 0.0083 0.0004 30.94 36.03 -5.09 -.06

+2 -2 40.5 4.7 0.0007 0.0001 29.86 34.39 -4.53 0.60

+1 -3 627.5 6.1 0.0042 0.0003 30.64 33.70 -3.06 0.59

+3 -4 29.1 2.3 0.00043 0.00009 29.66 33.91 -4.25 0.28

-6 -7 24.3 18.3 0.0016 0.0003 30.23 33.71 -3.48 0.43

+5 -3 210.5 10.8 0.0027 0.0002 30.45 33.29 -2.84 0.60
TABLE 7. Upper Limits

Star Mem Cnts |ogLx log B-Vo Offax Star Mere Cnts logLx log B-Vo Offax Star Mem Cnts logLx log B-Vo Offax

(sec -i) (Lx/Lbol) (arcmin) (sec -l} (Lx/Lbo|) (arcmin) (sec -1) (Lx/Lbol) (arcmin)

K 37 Y 0.00034 29.55 -4.70 0.35 18.8 P197 ? 0.00052 29.74 -2.09 1.67 14.6 P282 ? 0.00022 29.37 -3.30 1.10 12.4

K 39 Y 0.00039 29.61 -5.47 0.12 12.7 P198 ? 0.00033 29.54 -2.28 1.74 12.8 P284 ? 0.00024 29.40 -2.40 1.57 13.5

K 43 Y 0.00052 29.74 -5.55 -0.01 6.7 P199 ? 0.00029 29.49 -3.52 0.83 11.5 P286 ? 0.00016 29.23 -3.04 1.52 6.6

K 49 Y 0.00026 29.44 -6.80 -0.16 5.2 P290 ? 0.00040 29.63 -2.18 1.64 13.2 P287 ? 0,00031 29.51 -3.54 0.98 11,8

K 50 Y 0.00047 29.70 -5.51 0.08 13.0 P201 ? 0.00017 29.25 -3.37 1.79 10.6 P288 ? 0.00041 29.64 -2.62 1.36 13.6

K 57 Y 0.00033 29.54 -4.80 0.30 10.9 P203 7 0.00032 29.53 -3.40 0.97 10.6 P289 ? 0.00030 29.50 -3.44 1.04 14.4

K 65 Y 0.00032 29.53 -5.03 0.26 19.8 P204 ? 0.00045 29,68 -2.17 1.54 17.0 P291 ? 0.00064 29.83 -1.84 1.63 16.7

K 66 Y? 0.00044 29.67 -5.00 0.12 13.6 P205 ? 0.00031 29.51 -2.17 1.60 12.7 P292 ? 0.00017 29.25 -2.82 1.42 7.9

K 76A Y 0.00068 29.86 -5.90 -0.07 16.5 P206 ? 0.00050 29.72 -3.20 1.43 19.0 P293 ? 0.00077 29.91 -3.12 1.50 18.5

K 76B Y 0.00068 29.86 -5.14 9.00 16.5 P208 ? 0.00035 29.57 -3.58 1.09 11.8 P295 ? 0.00031 29.51 -3.52 1.01 16.4

K 82 Y 0.00019 29.30 -6.68 -0.12 19.8 P209 ? 0.00043 29.66 -2.11 1.66 11.9 P296 ? 0.00015 29.20 -3.53 1.05 7.3

P 19 Y 0.00022 29.37 -4.63 0.46 4.6 P210 ? 0.00015 29.20 -2.58 1.58 9.0 P297 ? 0.00027 29.45 -3.99 1.39 10.5

P 27 Y 0.00095 30.00 -3.73 0.59 16.2 P211 ? 0.00028 29.47 -3.77 1.37 9.9 P298 ? 0.00024 29.40 -2.55 1.56 11.0

P 39 Y 0.00094 30.00 -3.61 0.57 15.5 P212 ? 0.00060 29.80 -2.07 1.64 13.9 P299 ? 0.00062 29.82 -2.44 1.51 17.8

P 75 Y 0.00069 29.86 -3.48 0.71 14,1 P213 ? 0.00034 29.55 -3.42 0.96 14.2 P300 ? 0.00012 29.10 -2.52 1.64 7.9

P 77 Y? 0.00060 29.80 -3.16 1.07 14.7 P214 ? 0.00045 29.68 -3.11 1.10 13.8 P301 ? 0.00036 29.58 -3.06 1.43 14.5

P 83 Y? 0.00032 29.53 -3.31 1.11 13.8 P215 ? 0.00034 29.55 -3.53 1.03 11.4 P303 ? 0.00015 29.20 -3.00 1.50 9.1

P 94 Y 0.00010 29.02 -4.12 0.83 3.1 P216 ? 0.00038 29.60 -2.37 1.54 12.5 P304 ? 0.00051 29.73 -2.07 1.54 19.0

P 99 Y? 0.00025 29.42 -3.49 1.17 12.1 P217 ? 0.00077 29.91 -2.24 1.54 19.4 P306 ? 0.00029 29.49 -3.53 0.96 10.9

PI20 Y? 0.00031 29.51 -3.68 1.19 18.1 P218 ? 0.00028 29.47 -3.02 1.46 10.2 P308 ? 0.00050 29.72 -3.50 1.15 18.7

P121 Y? 0.00033 29.54 -3.56 1.19 18.8 P220 ? 0.00021 29.35 -2.26 1.62 10.5 P310 ? 0.00054 29.76 -3.11 1.15 14.6

P126 Y? 0.00072 29.88 -3.16 1.26 17.8 P222 ? 0.00066 29.84 -3.15 1.04 18.7 P311 ? 0.00025 29.42 -2.52 1.54 11.3

P243 Y? 0.00017 29.25 -3.20 1.57 9.3 P223 ? 0.00020 29.32 -2.82 1.38 7.6 P312 ? 0.00036 29.58 -2.36 1.75 15.0

P265 Y? 0.00019 29.30 -3.24 1.50 5.6 P224 ? 0.00037 29.59 -2.48 1.63 15.2 P313 ? 0.00035 29.57 -2.24 1.64 13.1

P283 Y? 0.00031 29.51 -2.72 1.55 7.2 P225 ? 0.00044 29.67 -2.69 1,28 15.6 P314 ? 0.00039 29.61 -2.74 1.35 I6.4

P305 Y? 0.00039 29.61 -2.74 1.56 17.3 P227 ? 0.00040 29.63 -2.63 1.52 16.9 P316 ? 0.00041 29.64 -3.13 1.38 13.6

P309 Y? 0.00070 29.87 -2.69 1.41 11.9 P228 ? 0.00017 29.25 -2.94 1.49 6.6 P317 ? 0.00041 29.64 -3.17 1.38 13.6

P335 Y? 0.00031 29.51 -2.61 1.63 13.1 P231 ? 0.00018 29.28 -2.46 1.71 6.9 P320 ? 0.00035 29.57 -2.45 1.52 14.1

P336 Y? 0.00033 29.54 -2.93 1.50 14.9 P233 ? 0.00033 29.54 -2.54 1.71 16.0 P321 ? 0.00033 29.54 -1.75 1.73 11.5

P343 Y? 0.00040 29.63 -2.76 1.56 15.5 P234 ? 0.00012 29.10 -3.18 1,33 5.8 P322 ? 0.00021 29.35 -3.48 1,09 11.4

P348 'Y? 0.00053 29.75 -2.43 1.60 17.2 P235 ? 0,00027 29.45 -2.09 1.73 6.6 P323 ? 0.00023 29.39 -2.77 1.43 12.2

P354 Y? 0.00057 29.78 -2.56 1.56 17.4 P238 ? 0.00030 29.50 -2.36 1.85 19.4 P324 ? 0.00055 29.76 -3.42 1.28 19.9

P373 Y? 0.00070 29.87 -2.16 1.70 18.4 P240 ? 0.00050 29.72 -2.10 1.93 18.7 P326 ? 0.00046 29.69 -3.24 0.97 15.0

K 38 ? 0.00063 29.82 -4.68 0.35 17.9 P241 ? 0.00046 29.69 -2.02 1.64 11.9 P328 ? 0.00027 29.45 -3.10 1.01 14.0

P 2 ? 0.00070 29.87 -4.27 0.25 15.9 P244 ? 0.00021 29.35 -2.57 1.61 3.1 P329 ? 0.00028 29.47 -2.77 1.42 14.9

P 15 ? 0.00014 29.17 -4.32 0.48 5.0 P245 ? 0.00007 28.87 -3.16 1.44 3.2 P331 ? 0.00017 29.25 -3.58 0.96 12.9

P 59 ? 0.00024 29.40 -3.91 1.09 16.2 P246 ? 0.00012 29.10 -2.63 1.69 3.0 P332 ? 0.00033 29.54 -3.51 0.91 16.5

P 63 ? 0.00027 29.45 -4.24 0.89 10.7 P250 ? 0.00018 29.28 -3.71 1.08 10.8 P333 ? 0.00026 29.44 -2.40 1.73 13.7

P 64 ? 0.00034 29.55 -3.55 0.77 15.5 P252 ? 0.00012 29.10 -3.98 1.01 5.6 P337 ? 0.00029 29.49 -2.40 1.60 13.6

P 65 ? 0.00036 29.58 -3.86 1.04 13.8 P253 ? 0.00017 29.25 -2.68 1.49 6.5 P338 ? 0.00045 29.68 -2.09 1.61 14.0

P 67 ? 0.00033 29.54 -3.30 0.96 13.3 P254 ? 0.00057 29.78 -3.13 1.20 14.6 P339 ? 0.00063 29.82 -2.99 0.97 17.1

P 78 ? 0.00033 29.54 -3.64 0.63 13.5 P255 ? O.00010 29.02 -3.09 1.44 1.6 P340 ? 0.00043 29.66 -2.23 1.70 19.2

P 93 ? 0.00009 28.98 -4.49 1.07 3.6 P256 ? 0.00024 29.40 -2.83 1.39 I0.I P341 ? 0.00019 29.30 -2.47 1.61 14.4

P101 ? 0.00029 29.49 -3.76 1.33 12.0 P258 ? 0.00012 29.10 -2.83 1.51 6.9 P342 ? 0.00057 29.78 -2.02 1.65 17.5

P104 ? 0.00041 29.64 -3.67 1.20 14.9 P259 ? 0.00058 29.79 -2.06 1.63 14.5 P344 ? 0.00023 29.39 -2.45 1.75 17.9

PII4 ? 0.00030 29.50 -3.96 1.04 12.2 P260 ? 0.00060 29.80 -2.48 1.48 16,0 P346 ? 0.00052 29.74 -2.98 1.04 15.4

P119 ? 0.00037 29.59 -3.77 1.14 19.2 P262 ? 0.00066 29.84 -1.88 1.63 17.6 P347 ? 0.00047 29.70 -2.90 1.39 15.2

P122 ? 0.00040 29,63 -3.18 0.95 17.0 P263 ? 0.00021 29,35 -3.66 0.96 9.0 P349 ? 0.00031 29.51 -3.49 0.93 19.2

P184 ? 0.00036 29.58 -3.06 1.58 17.2 P264 ? 0.00011 29.06 -2.56 1.64 2.2 P350 ? 0.00035 29.57 -2.35 1.73 16.8

P185 ? 0.00041 29.64 -3.55 1.12 17.1 P267 ? 0.00035 29.57 -3.35 0.86 12.9 P352 ? 0.00041 29.64 -3.35 0.90 16.1

P186 ? 0.00023 29.39 -2.04 1.66 13.4 P268 ? 0.00014 29.17 -3.31 1.28 8.4 P359 ? 0.00034 29.55 -2.36 1.60 16.7

P187 ? 0.00090 29.98 -2.06 1.61 18.3 P269 ? 0.00028 29.47 -2.72 1.62 12.3 P360 ? 0.00049 29.71 -2.36 1.54 19.4

P188 ? 0.00045 29.68 -2.35 1.48 17.3 P271 ? 0.00008 28.93 -3.85 1.28 5.8 P362 ? 0.00063 29.82 -2.40 1.36 16.7

P189 7 0.00031 29.51 -2.34 1.74 12.9 P272 ? 0.00024 29.40 -2.38 1.60 14.7 P364 ? 0.00070 29.87 -2.17 1.47 18.4

P190 ? 0.00055 29.76 -3.06 1.06 12.3 P273 ? 0.00029 29.49 -3.17 1.42 14.4 P366 ? 0.00082 29.94 -2.71 1.05 18.2

P191 ? 0.00039 29.61 -2,50 1.58 18.8 P274 ? 0.00027 29.45 -2.17 1.73 8.2 P367 ? 0.00065 29.84 -1.44 1.70 19.0

P192 ? 0.00056 29.77 -2.17 1.58 12.2 P275 ? 0.00009 28.98 -3.15 1.47 4.8 P369 ? 0.00080 29.93 -1.89 1.68 18.1

P193 ? 0.00054 29.76 -2.97 1.35 16.7 P276 ? 0.00027 29.45 -3.46 0.97 13.1 P370 ? 0.00040 29.63 -2.60 1.50 19.7

P194 ? 0.00023 29.39 -3.47 1.30 12.0 P277 ? 0.00009 28.98 -2.95 1.51 4.7 P372 ? 0.00075 29.90 -1.82 1.67 18.1

P195 ? 0.00035 29.57 -2.54 1.39 16.4 P278 ? 0.00008 28.93 -3.34 1.39 4.7

P196 ? 0.00045 29.68 -3.13 0.97 12.1 P281 ? 0.00044 29.67 -2.86 1.36 12.9
TABLE 8. Bright Cluster Members

Kopff _ Spt b X-ray Activity logLx

62 B3IV detected 29.61

73 B5IV detected 30.53

64 B5IV detected 29.60

105 B6V SB2 beyond field ...

58 B6V detected 30.64

49 B6Vp CII SB1 detected 29.28

72 B7V SB1 beyond field ...

82 B6Vp CII upper limit <29.30

23A BSV beyond field ...

76A B8.5Vp Hg-Mn upper limit <29.86

32 B8V shell beyond field ...

22 B8.5V beyond field ...

23B B9.5V beyond field ...

81 BSV beyond field ...

43 A1V detected 29.67

50 AI+A2V SB2 upper limit <29.70

102 A2V beyond field ...

39 Am (classical) upper limit <29.61

67A A2V VB (detected) 30.12

a Kopff (1943)

b Abt & Levato (1975)
TABLE9.SaturationColorsforOther Clusters

Cluster Age B-Vo Reference

(Myr)

IC2602 ........... 30 --,0.65

a Per ............ 50 --'0.6

Pleiades ......... 70 ,-_0.75

NGC6475 ....... ,-,220 ,_1.0

Praesepe/Hyades ,_600-800 ,,_1.4

Randich et al. 1995

Prosser et al. 1996

Fig. 4

Prosser et al. 1995

Randich & Schmitt 1995

Stern et al. 1995
TABLE 10. Candidate Variable X-ray Sources

IC4665-A IC4665-B

Name Mem Rate Error Rate Error ML(A) ML(B)

(s-1) (s-,)

K56 N 0.0070 0.0005 0.0097 0.0007 200.0 248.2

K43 Y 0.0007 0.0002 <0.00003 ..... 5.9 0.1

P230 Y? 0.0003 0.0001 <0.0003 ..... 5.2 3.0

P232 7 0.0013 0.0002 <0.0003 ..... 41.3 1.7

P237 ? 0.0005 0.0001 <0.0001 ..... 10.7 1.1

P248 Y? <0.0002 ..... 0.0006 0.0002 1.2 8.4

R6 N 0.0005 0.0002 .......... 10.5 0.0

R12 N <0.0002 ..... 0.0007 0.0002 2.6 10.9
TABLE 11. Additional X-ray Sources

ID RA (2000) DEC ML Offax Cnts Err logLx

(arcmin) (sec -1) (sec- 1)

R1 17 45 01.7 5 39 27 11.9 10.4 0.0005 0.0001 29.75

R2 17 45 05.7 5 47 39 9.4 13.1 0.0007 0.0002 29.84

R3 17 45 13.9 5 37 53 15.3 7.3 0.0005 0.0001 29.72

R4 17 45 14.4 5 30 43 11.1 10.5 0.0005 0.0001 29.72

R5 17 45 18.5 5 45 00 7.9 9.0 0.0003 0.0001 29.55

R6 17 45 24.4 5 33 06 7.1 7.1 0.0004 0.0001 29.57

R7 17 45 43.2 5 55 16 12.2 16.9 0.0011 0.0002 30.05

R8 17 45 46.7 5 37 26 18.8 1.3 0.0005 0.0001 29.71

R9 17 45 48.5 5 44 48 27.7 6.5 0.0005 0.0001 29.70

R10 17 45 55.2 5 37 55 18.2 3.0 0.0003 0.0001 29.50

l:tll 17 45 57.3 5 46 56 5.7 9.2 0.0003 0.0001 29.47

l:t12 17 46 03.8 5 39 38 13.8 5.3 0.0004 0.0001 29.65

R13 17 46 06.3 5 32 32 8.7 8.2 0.0003 0.0001 29.45

R14 17 46 07.6 5 53 54 123.5 16.6 0.0041 0.0003 30.63

R15 17 46 37.2 5 43 34 12.7 14.4 0.0009 0.0002 29.96
TABLE 12. Stars Associated with X-ray Sources

Star V B-V V--IK Mem Star V B-V V--IK Mem

R1A 14.37 0.86 0.79 N R10B 18.80 ... 0.80 N

R3A 17.23 1.08 1.00 N R10C 16.89 ... 0.63 N

R3B 18.27 0.92 1.13 N R10D 17.58 ... 0.95 N

R3C 17.98 0.87 0.62 N R10E 14.03 ... 0.36 N

R3D 17.31 0.76 0.47 N RllA 16.81 0.80 0.69 N

R3E 13.60 0.84 0.74 _ RllB 17.94 0.81 0.68 N

R3F 15.21 0.97 0.84 N RllC 18.78 1.05 1.07 N

R3G 17.75 0.74 1.01 N RllD 18.40 1.28 1.17/' N

R3H 19.10 1.57 1.83 N RllE 17.77 1.01 0.87 N

R4A 15.05 1.17 1.06 _ RllF 19.55 ... 1.29 N

R4B 15.83 0.90 0.69 N RllG 19.28 ... 1.05 N

R4C 15.45 0.89 0.82 N R12A: 14.91 ... 0.71 N

R4D 16.66 0.94 0.85 N R12B: 16.53 ... 0.61 N

R4E: 19.17 ... 1.01 N R12C 16.78 ... 1.42 N

R6A 15.80 1.19 1.02 N R12D 18.33 ... 0.68 N

R6B 16.97 0.70 0.76 N R12E 18.40 ... 0.87 N

R6C 15.79 0.93 0.85 N R12F 18.87 ... 2.37 N?

R7A: 14.26 1.16 0.97 Y? R12G 17.47 ... 0.51 N

RTB 18.98 ... 1.27 N R13A 17.21 ... 1.13 N

R7C 15.33 1.01 0.91 N R13B 19.81 ... 1.53 N

R8A 19.18 ... 0.36 N R13C 19.49 ... 0.81 N

R8B 20.12 ... 0.85 N R13D 18.05 ... 0.98 N

R8C 19.55 ... 0.93 N R13E 18.76 ... 1.12 N

RSD 20.50 ... 1.97 N R13F 17.44 ... 0.87 N

R8E 15.57 ... 0.96 N R13G 18.18 ... 1.40 N

R8F 16.78 ... 0.65 N R13H 16.21 ... 0.78 N

R9A 14.80 0.92 0.81 N R13I 18.30 ... 1.77 N

R9B 15.61 0.59 0.54 N R13J 16.80 ... 0.68 N

R9C 18.78 0.55 0.54 N R14 a 13.60 1.36 1.62 N

R9D 13.32 ... 0.67 _ R15A 19.36 ... 2.46 N?

R10A 16.92 ... 0.13 N R15B 18.72 ... 1.07 N

a companion to R14: IK -----16,V>20-21, sep _ 3" south
s _ _ s

"'l'il .... I

,,It,J:l,,,, i,i

t_

.=_

"" E

b_O
4_

0

%,,

0

0

t.L.

m

0

x-I 6Ol

0

!

o

I

m

p_

ox

Ox

E

t,_

r_

0

_() .,,._

. ,,.._
i I I I

B

=-.=IP R

I

m

, 1 1 1 1

I

O-_Ib

Q

---.-D

w

v

i

0 i

i

0

C_

I

_I- U'} ¢&) I_

I i I I

0

u')

!

o

>

I

O3

r._

ct)

0

('°"-I/X7)5oi
o

0

o

0

0

0

o

o o _

0 _

0 0 0 0

o _o

o o °0 0

o _o o

o o o

o

o

o o

0%0 %

o

0

0 o°o _

o

oo_

Ooo

o

ooo oo Do

000----3,

O0 0 0

o

o

o

o

%

x-1 6o I

I

m

p.....

ox

ox

o

0

....., I::z,

_ r,..0
a)

u) to

L_ (D

G) CO

0

0 O0

0

%o

0

_ I_ 0

i>_1_ _ 0

_...,, _-_ _.o 0 %.

i _' _OI_GD b..j _,

• ___7 °o_°o_

,.ooO;-.,.,- oe:O o

• _w._ o o 0 0

0

0

• o

0

• og-" _" 0 _

0

0

00o 0

Ib OO_o_

0 0 D. _

6) o o

Q 0 00_

o e ---e.p D,

0

0 o 0 _0 _

o _ o

0 0 0 _

0

0000 0 _0 _ 0

q) 0 _.

o _o 0 •

0 • 0 _ 0

• 0

_0

O_

0

o e_. o _

0

0 _' _'

ee ,8 "e

o x__

tO

tO

0

t_

i e

00

o

>

I

m

r_

o_

O_

0
m

I

c_

!

I I I i

0

I I I I i

_00_ _i

Oob-_o _ %

0 ioo_1O- 01> 0

o--_oi_._ =, _ =,.

t>lp

o ,_ o c_"

o_c_

00( _ 0 00---_ _0

I _ 0

@ ,o--_ °

oOi 0 0

0 0

0 _ O_ o

'cPO o

6 o

I O0 Oco

o_ o

r-"

I

0

! ! I I

I

0

Q__--b o

0

I I I

I I I

GI)

c_o D.

0

c_

o° °_

Oo

@

0

0

0

If,,

0 Oo

0

"_0

@

0

0

I ! I I I i I

•_ u')

I !

''1''''1

_. _o

o @

D,

b b

----_0I_ 0

o

---b

0

• @ o

p

0 -

@ 00_

,,,Ill

I !

o

>

I

o

Lo

I"

5oi
I I I I

0

0

D

! I

O4

|

I I

Oe-_ _

°

I0 0

I

n 0

aO

I .._

d

I 0 0 -'-'_

i

i

Oo_o---;

I I I I I I

0

0

O0

0

0 0

0

0

_ o o

i O0

I0 0

' _8

C_ 0 0 0

CO 0 0

I

_ o Oo uO o

o _o°r_ 0

_o o

_ o "-*

0 0 0

(ID 01_ 0

_l,J,,l_,,,l'!

0

e

0

I 1

!

,,,,I,,,,I,,

I |

em

ID

I... ¢d_

0

0

0

0

0

• • O° O0_

m

,111,,,1

II al

0

LID

I

o

>

I

m

C

E

r,..T.,

0

('°q-I/x-I) 5Ol
I

%

%

I

I i I

%

%

%

! I

%

i I I I I

%

,%

%

%

%

%

%

%

%

f i I I I I

L_

C)

I I I i I i

::> %

%.

%

%.

%

%,

.%

%

%

>

>

I

I I I !

i i i I

n._

6

A

o

I

m

%

%

I I I I I

C_

i I I

- > %

%

%.

%

@4

C',

l:m

°w,,_

E

0

_ °_

_ C

e_

ffl

>

v

"-" 0
IIIIIIIII

C) A

O

I

rn

oE

0

z...

¢)0

L_

,._ c_

0

(X) 0

I

Ev

v

ro_

f$)

0>

00
m,

|

(D

s

s

J

d

J

s

J

s

f

s

J

f

s

#

J

A_

S

S

S

S

S

S

J

S

A

C_L

|

I

J

$

I I I I

#

J

.e

s'

s

f

#

f

s

s

s s i

s

I

!

0

t-.,

Q_

v

I

I

!

G',

E

O

0
I I I I

o

m

D

m

-T

[]

[]

D

[]

[]

O

[] []

n

n

[3_-_ D

[]

m

m

m

u

m

rl

DODO

[]

[]

D

D

[]

[]

i..U_

tO

O_ CO

[] •

I

m

m

I

oO

_v

rl

,_ _o

cO

.=_

E

0

C_

!

!

14")

!

C° /x7) 5oi

!

!
_l

-- O

I

B

B

m

m

B

m

m

m

B

B

B

m

m

B

B

m

B

B

m

m

B

m

U)

O) u_

"IO _O

o •

i i I l

O

O

O

l l I l

0

0

0

i l I i

0

Q 0

0

0

0

o @

0

0 0

0 0

0 0

i I '

0

0

0

0

O _ O O

O __

O

O

m

O _

o-"

o_

OD

c_

!

cO

Od

!

O_

!

04

cd

|

_Y

v

00

CO

!

_t

|

04

_f

!

_t

_f

!

L_

o_

u_

o

O_

c_

_o_

cO £O xt 04 O
Fig. 9

Giampapa, Prosser & Fleming (1997