Rotational Velocities and Chromospheric/Coronal Activity of Low-Mass Stars in the Young Open Clusters Ic 2391 and Ic 2602 John R.Stauffer,1 Lee W

Home , IC 2391, IC 2602

ROTATIONAL VELOCITIES AND CHROMOSPHERIC/CORONAL ACTIVITY OF LOW-MASS STARS IN THE YOUNG OPEN CLUSTERS IC 2391 AND IC 2602 JOHN R.STAUFFER,1 LEE W. HARTMANN,PAND CHARLES F. ROSSER Smithsonian Astrophysical Observatory, 60 Garden Street, Cambridge, MA 02138 SOFIA RANDICH2 European Southern Observatory, Karl-Schwarzschild Strasse 2, D-85748 Garching, Germany SUCHITRA BALACHANDRAN1 Astronomy Department, University of Maryland, College Park, MD 20742 BRIAN M. PATTEN1 Smith College, Five Colleges Astronomy Department, Northampton, MA 01063 THEODORE SIMON Institute for Astronomy, University of Hawaii, Honolulu, HI 96822 AND MARK GIAMPAPA National Optical AstronomyObservatories,3 National Solar Observatory, Tucson, AZ 85726 Received 1996 September 16; accepted 1996 November 12

ABSTRACT We have obtained high-resolution, moderate signal-to-noise ratio spectra for approximately 80 candidate low-mass members of the nearby, very young open clusters IC 2391 and IC 2602. Most of the stars observed are conÐrmed as cluster members based on a combination of photometric and spectroscopic criteria. We provide radial velocities, rotational velocities, and Ha equivalent widths for these stars. From comparison to theoretical preÈmain-sequence (PMS) evolutionary isochrones from DÏAntona and Mazzitelli, we derive an estimated age of the two clusters of D25 Myr. By contrast, the usually quoted upper main-sequence turno† age for the clusters is D35 Myr. We do not believe that this provides evidence for noncoeval star formation within these clusters, but rather that the best age estimate for them given the uncertainties is D30 ^ 5 Myr. In principle, the scatter of stars about the PMS isochrone provides a measure of the age spread among the low-mass stars in these clusters; however, with the data presently available, we are able to derive only a relatively uninteresting upper limit for an age spread of order 20 Myr. We compare the rotational velocity distribution for IC 2391/2602 to that observed for the Pleiades. For the G dwarfs in the IC clusters, we resolve rotation in all but one of the probable cluster members, and thus except for inclination e†ects, our data provide the complete distribution of rotational velocities for solar mass stars on their arrival on the ZAMS. The projected rotational velocities (v sin i) of the G dwarfs in the two IC clusters span the range from D8toD200 km s~1. Comparison of the distribution of rotational velocities for the G dwarfs of the Pleiades and the IC clusters indicates that both the slow and the rapid rotators lose of order half their angular momentum during the Ðrst D35 Myr on the main sequence if they rotate as solid bodies. The low-mass stars in these two clusters exhibit a similar correlation between rotation and coronal activity as is found in several other young open clusters. That is, there is a large spread in coronal activity for stars with v sin i \ 25 km s~1, where we assume there is an intrinsic link between increasing rotation and increasing activity superimposed upon which are a variety of observational and physical mechanisms that act to smear out this relation; above v sin i D 25 km s~1, all of the low-mass stars have log (L /L ) D [3.0, the canonical ““ saturation ÏÏ limit. Our measurements of the Ha equivalent widths are consistentX bol with a similar relationship holding for chromospheric activity. One and possibly two of our spectra for M dwarf members of the IC clusters show broad wings for the Ha proÐle, which we attribute to a Ñare event or to microÑares. Since spectra of a small sample of late-type M dwarfs in the Pleiades also showed similarly broad Ha wings, this suggests that Ñare frequencies for very young M dwarfs may be quite high. Subject headings: open clusters and associations: individual (IC 2391, IC 2602) È stars: coronae È stars: chromospheres È stars: evolution È stars: low-mass, brown dwarfs È stars: rotation

1 Visiting Astronomer at the Cerro Tololo Inter-American Observatory, operated by the Association of Universities for Research in Astronomy, Inc., under contract with the National Science Foundation. 2 Partially based on observations carried out at the European Southern Observatory, La Silla. 3 The National Optical Astronomy Observatories are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation. 776 YOUNG OPEN CLUSTERS IC 2391 AND IC 2602 777

1. INTRODUCTION telescope of the European Southern Observatory in 1994 April. At the 4 m telescope, we used the red, long-camera Open clusters provide an empirical database with which echelle spectrograph with a 31.6 line mm ~1 grating blazed to study the time evolution of a variety of stellar properties. in the red and a 120 km slit width(0A.8 on the sky) providing In particular, open clusters can be used to (1) provide color- a 2 pixel resolution of about 0.15Ó at Ha. The detector was magnitude diagrams that can be compared to theoretical a Tektronix 2048 ] 2048 CCD, read out with four ampli- isochrones and thus enable tests of the theoretical evolu- Ðers using the ARCON controller, allowing a spectral tionary models (see, e.g.,Iben & Talbot 1966; Adams, coverage of roughly 5800È8200Ó. We used the IRAF Strom, & Strom1983; Prosser et al. 1994); (2) determine the QUAD package to bias-subtract and Ñat-Ðeld individually distribution of rotational velocities for low-mass stars as a each of the CCD quadrants and then used standard pack- function of age and thus place constraints on the role of ages within IRAF to extract the echelle orders and preÈmain-sequence (PMS) circumstellar disks and stellar wavelength-calibrate the spectra. Comparison lamp spectra winds in regulating rotational velocities(Cameron & Li were obtained only at the beginning and end of the night, 1994; Keppens,MacGregor, & Charbonneau 1995); (3) and hence the wavelength calibration for each target spec- study internal mixing mechanisms in low-mass stars via the trum is slightly in errorÈmostly as a zero-point shiftÈ mass and age dependence of the surface abundances of owing to Ñexure and thermal variations during the night beryllium, boron, and lithium(Garc•a Lopez et al. 1994; (we correct for this in a subsequent step, as discussed Boesgaard 1991; Chaboyer, Demarque, & Pinsonneault below). At the 3.6 m telescope, we used the CASPEC echelle 1995); and (4) study the age and mass dependence of the spectrograph(Pasquini & Gilliotte 1993) with a Tektronix coronal and chromospheric activity of low-mass stars 512 ] 512 CCD, a 31.6 line mm~1 grating, the short camera (Caillault 1996). Implicit in the use of open clusters for this and the red cross-disperser. A 280 km(2Aon the sky) slit purpose is the availability of accurate membership lists to width was used. The wavelength range for the ESO spectra identify sets of stars that can be conÐdently assumed to was 5500È8300Ó, and the 2 pixel resolution was D0.35 Ó. belong to the cluster(s) under study. Usually, the primary Separate Ñat-Ðeld and comparison lamp spectra were component of such membership studies has been a proper- obtained for each program object. IRAF routines were motion survey. However, in some casesÈparticularly for again used for the standard processing and wavelength cali- clusters in the southern hemisphereÈappropriate Ðrst- bration of the spectra. epoch photographic plates have not been available, and Radial and rotational velocities for the program stars therefore it has been necessary to consider alternate means were determined using cross-correlation routines (see Tonry to identify cluster members. Because young, low-mass stars & Davis1979; Hartmann et al. 1986). The stars observed at are usually strong X-ray sources, one means to accomplish ESO were generally G and K dwarfs, and we were able to this goal is to use ROSAT PSPC images of the cluster choose a single echelle order (jj6410È6495) that has a region to sort the relatively small number of active cluster number of moderately strong iron and calcium absorption members from the vastly larger number of background/ lines for the cross-correlation analysis. The template for the foreground Ðeld stars. This technique has recently been ESO cross-correlations was a high signal-to-noise ratio used for the D35 Myr old clusters IC 2391(Patten & Simon (S/N) spectrum of the daytime sky, and the radial velocity 1996)and IC 2602 (Randich et al. 1995). zero point was set via observation of the radial velocity The identiÐcation of cluster members using ROSAT standards HD 126053, HD 136202, HD 154417, and HD imaging data is not a foolproof techniqueÈthe X-ray error 187691(Latham & Stefanik 1991). By comparison of results circles often contain more than one optical candidate, and obtained using di†erent radial velocity standards as the thus the wrong star can be associated with the X-ray source, template, we estimate that for the ESO spectra the radial or in some cases the optical identiÐcation is made correctly velocity accuracies are 1È2kms~1 for the slow rotators but the star is not in fact a cluster member but is instead a (v sin i ¹ 20 km s~1), increasing to several km s~1 for stars Ðeld star whose photometric characteristics approximately with v sin i [ 50 km s~1. match those expected for members of the cluster. For this The stars observed at CTIO covered a wider spectral type reason, it is necessary to supplement the ROSAT data and range, including G, K, and M dwarfs, and so it was not optical photometry with other data that provide additional possible to use a single order for all of the analysis. For the membership constraints. In this paper, we report the results G and K dwarfs, we used two orders for the cross- of a high-resolution spectroscopic study of a large fraction correlation analysisÈjj7400È7470 and jj6200È6270È of the candidate cluster members identiÐed with ROSAT in both orders dominated by moderate-strength atomic IC 2391 and IC 2602. We Ðrst use the derived radial veloci- absorption lines. For the M dwarfs, we also used two echelle ties, lithium 6708Ó line strengths, and Ha equivalent widths ordersÈjj7770È7870 and jj6640È6750; in this case, the to produce improved membership lists for the two clusters. criterion for selection was for the orders to be dominated by With that list in hand, we then discuss the properties of TiO bands (but to avoid the strongest TiO band heads these stars that can be used to constrain better the ages of where saturation and blending become a problem). The these two clusters, the rotational velocity evolution of stars used as cross-correlation templates were GL 243 (for young low-mass stars, and the dependence of coronal activ- G dwarfs), GL 105A (for K dwarfs), and GL 447 (for M ity on rotation. Lithium abundances for some of the stars in dwarfs). In order to correct for temporal drifts in the wave- our survey will be reported inRandich et al. (1996a). length calibration during the night, we also cross-correlated OBSERVATIONS AND DATA REDUCTION one order dominated by terrestrial molecular oxygen 2. absorption lines (jj7160È7280) and derived a radial veloc- The spectra for this program were obtained during ity zero-point correction for each target object by requiring observing runs at the 4 m telescope of the Cerro Tololo that the radial velocity of the atmosphere above CTIO be Inter-American Observatory in 1995 January and the 3.6 m constant, which is an adequate assumption at our level of 778 STAUFFER ET AL. Vol. 479 interest. For the M dwarfs, the Ðnal calibration of the radial clusters. The mean di†erence between our measured radial velocities was set by observation of several Ðeld dM stars velocities for these stars and the velocities computed from (GL 205, GL 273, and GL 447) with accurate published theStefanik & Latham (1985) Hyades convergent point radial velocities fromMarcy & Benitz (1989). Seven Hyades model is 0.94 km s~1, with a dispersion of 0.75 km s~1. M dwarfs were also observed as secondary standards, where In general, the same cross-correlation analyses used to the ““ true ÏÏ radial velocities were determined from the derive the radial velocities were also used to derive rotation- convergent-point solution provided byStefanik & Latham al velocities following the procedures outlined in Hartmann (1985). For the G and K dwarfs observed at CTIO, the et al.(1986). This was done separately for each of the two radial velocity zero point was set by cross-correlation of our echelle orders, and the results were averaged. Based on pre- template star versus a spectrum of the dusk sky; secondary vious experience and on the comparison of the results for radial velocity standards were HR 1101 and HR 4540 from the two orders, the rotational velocities should typically the radial velocity standard list of Latham & Stefanik. have 1 p accuracies of order 10% of the v sin i. The internal accuracy of the CTIO radial velocities can A few of our stars have such large rotational velocities be estimated by comparison of the di†erence in radial that the cross-correlation routine failed to give reliable velocities between the two echelle orders used for each spec- results. The reason for this is illustrated inFigure 1, where tral type range. For the GK dwarfs, the 1 p dispersion of the spectra of two of these very rapid rotators are compared to di†erence in radial velocities derived from the two orders two more slowly rotating cluster members. Basically, for was 0.98 km s~1 for stars with v sin i ¹ 20 km s~1; the rotational velocities greater than about 150 km s~1, all but same number for the M dwarfs is 0.94 km s~1. The errors the very strongest absorption features are completely for rapid rotators are larger and are a function of both the washed out. For this handful of stars, we have simply con- rotational velocity and S/N of the spectrum. As a rough volved a limb-darkened rotational broadening function estimate, it is likely that the errors increase approximately with the spectrum of an appropriate standard star and linearly with v sin i to about 5 km s~1 for v sin i D 100 km selected the rotational velocity which best matches the s~1. proÐle shape of the Na I D doublet in the program star. The A more stringent test of the internal accuracy of the proÐles change shape signiÐcantly when the broadening CTIO radial velocities is provided by comparing the function is changed by ^25 km s~1, and we therefore derived radial velocities for stars observed more than once assume the 1 p accuracy of these rotational velocities is of during the run. There are eight stars in that category; the that order. We do not list radial velocities for these stars mean absolute di†erence between the two radial velocity since this visual-comparison technique does not provide determinations for these eight stars was 0.65 km s~1. that information. The best test of the external accuracy of our radial veloci- Ha equivalent widths for the program stars were derived ties is provided by the seven single Hyades M dwarfs that using the SPLOT routine in IRAF. With one or two excep- we observed at CTIO during the same run used for the IC tions, this was straightforward.Figure 2 shows the Ha

IC2391-67a 1500 IC2391-50a 1500

1000 1000 Counts Counts

500 500

0 0 5880 5900 5920 5940 5880 5900 5920 5940 Wavelength Wavelength

2500 IC2391-80a IC2602-88a 1000

2000 800

1500 600 Counts Counts 1000 400

500 200

0 0 5880 5900 5920 5940 5880 5900 5920 5940 Wavelength Wavelength

FIG. 1.ÈSpectra of four of the program stars showing the change in appearance of the spectra with increasing rotational velocity. The Na I D doublet lines provide the best discriminant of rotational velocity for the most rapidly rotating stars in our sample. No. 2, 1997 YOUNG OPEN CLUSTERS IC 2391 AND IC 2602 779

1000 region for the star that required special attention, IC 2602- IC2602-R38 R38. A normal, moderately rapidly rotating dMe star has 800 an Ha proÐle that is rounded at the top and is well Ðtted by

600 a Gaussian proÐle whose width is compatible with the starÏs

counts rotational velocity. The top of the Ha proÐle for R38 Ðts 400 this description, but the broad wings on the bottom of the

200 proÐle do not allow a single Gaussian Ðt to be successful, as is illustrated in the second panel ofFigure 2. The third panel 0 6500 6520 6540 6560 6580 6600 shows a two-Gaussian Ðt, which does match the observed wavelength proÐle well. The equivalent widths of the two components 1200 are similar, and the Gaussian width of the broad com- 1000 IC2602-R38 ponent is about 130 km s~1. 800 We carefully examined the Ha proÐles of all of our other emission-line stars for broad components and found only 600

counts one other possible candidate: IC 2391-R71a. The Ha spec- 400 trum for that star is shown inFigure 3. The broad com-

200 ponent is in this case only marginally detectedÈwe suspect such a component is present, but it is not deÐnitely detected. 0 6540 6545 6550 6555 6560 6565 6570 6575 6580 A brief discussion of the signiÐcance of these broad Ha wavelength components will be given in ° 3.2.

800 Tables1, 2, and 3 provide our derived radial and rotational velocities and our Ha equivalent widths for the IC IC2602-R38 600 2391 and IC 2602 target stars. Also provided are V magni- tudes,(V [ I) colors, and coronal activityÈlog 400 counts (L /L )ÈasgivenC by Randich et al. (1995), Patten & X bol 200 Simon(1996), and Prosser, Randich, & Stau†er (1996). Note that the V [I colors used throughout the paper are on the 0 Cousins system. 6540 6545 6550 6555 6560 6565 6570 6575 6580 Table 4 provides Ha equivalent widths for the 10 stars wavelength proposed as members of IC 2391 byStau†er et al. (1989, FIG. 2.ÈT op panel: the observed Ha proÐle for IC 2602-R38; middle hereafterSHJM). In SHJM, we had reported only the pres- panel: the same spectrum but this time including a single Gaussian Ðt to ence or absence of Ha emission; we now provide in Table 4 the Ha proÐle; bottom panel: the Ðt of the proÐle using two Gaussians. The the Ha equivalent widths measured from these 1989 spectra FWHM of the broad component is D300 km s~1. Note that for the bottom spectrum, the continuum of a dM star of similar type has also been (note that in all of the tables, we use the convention that subtracted (from which a weak Ha absorption line was excised prior to positive equivalent widths correspond to Ha emission). performing the subtraction).

1200

1000 IC2391-71a

800

counts 600

400

200

0 6540 6545 6550 6555 6560 6565 6570 6575 6580 wavelength

FIG. 3.ÈObserved Ha proÐle for IC 2391-R71a. This is the only other star in our sample that may have a discernible broad Ha component. 780 STAUFFER ET AL. Vol. 479

TABLE 1 CTIO SPECTROSCOPY OF STARS IN IC 2391

Stara V (V [I) log L /L v sin iv Ha E.W. Li? Remarksb C X bol helio R3a...... 10.95 0.74 [4.21 10. 14.9 [1.1 Y R6a...... 13.77 1.94 [3.38 \6, \7 [20.6, 49.9 ]1.4 . . . SB2 R7 ...... 9.63 . . . [3.98 21.0 15.6 [1.5 Y R14...... 10.45 0.69 [3.59 47. 14. [1.1 Y R16a ...... 11.84 0.94 [3.25 22.0 15.5 D0. Y R18a ...... 13.54 1.53 [3.28 8. 14.6 ]0.4 Y R25a ...... 13.83 1.37 [4.09 \6 48.9 [0.9 N n.m.c R27...... 13.97 2.38 [3.99 \7 [1.5 [0.19 . . . n.m. R30...... 9.88 . . . [4.48 43. 30.1 [1.4 Y SB1? R31...... 11.22 0.73 [4.18 17.0 22.3 [1.1 Y SB1? R44...... 9.69 . . . [3.95 67: 14.2 [1.7 Y R45a ...... 10.70 0.90 [3.62 [200 . . . [1.4 . . . R49b ...... 14.34 1.89 [3.26 12.2 13.7 ]1.0 Y R50a ...... 12.54 0.91 [3.32 64. 17. [0.15 Y R62a ...... 11.73 0.99 [3.30 52. 12: D0. Y R64a ...... 15.32 1.83 [3.18 20.0 6. ]1.3 . . . n.m.? R65...... 14.13 1.79 [3.24 8. 14.5 ]0.8 . . . R67a ...... 11.71 1.03 [3.45 8: 14.9 [0.1 Y R69a ...... 11.67 0.90 [3.48 19.0 15.1 [0.3 Y R70...... 10.85 0.75 [3.77 17.0 13.4 [0.9 Y R71a ...... 15.32 2.41 [3.31 23.5 20.9 ]3.7 . . . SB1? R72...... 11.46 0.84 [3.21 15.0 14.1 [0.6 Y R76a ...... 12.76 1.24 [3.11 8. 14.4 ]0.3 Y R77a ...... 9.91 0.60 [3.77 95: 9: [1.4 Y R78...... 10.44 0.73 [3.14 50: 22.8 [0.9 Y SB1? R80a ...... 11.98 1.04 [3.29 D150 . . . D0? Y?

a Star name fromPatten & Simon 1996. Note that for consistency with the IC 2602 data, we have changed the Patten and Simon naming convention from VXRnnn to simply Rnnn. b REMARKS.Èn.m. \ nonmember; n.m.? \ possible nonmember; SB1? \ possible single-lined spectroscopic binary; SB2 \ double-lined spectroscopic binary. See text for details. Special caseÈR64a (n.m.?): This star has a discrepant radial velocity, but also falls well below the locus of the other IC 2391 stars in a C-M diagram. TiO indices derived from the spectrum conÐrm the photometric V [I color. Thus, despite the Ha emission, this star is probably a nonmember. However, its v sin i is very large for a star this late (early M), and thus it probably is quite young in any case. c A follow-up HRI observation of the core of IC 2391(Simon & Patten 1997) fails to conÐrm the PSPC source R25a. All of the PSPC Ñux attributed to this source should be assigned exclusively to R24 (instead of being split between the two sources).

For the discussion to follow, we assume distances and within the expected accuracy of the cross-correlation tech- reddenings for the two clusters compatible with the papers nique. While this criterion is straightforward for ““ single ÏÏ that discuss their X-ray characteristics(Randich et al. 1995; cluster stars or members of wide binaries, even a cluster Patten& Simon 1996). That is, for IC 2391, we assume member can have a discrepant radial velocity if it is a spec- (m [ M) \ 6.05 andE(V [I)c \ 0.01, and for IC 2602, we troscopic binary and only one spectrum is available. Based assume(m0 [ M) \ 5.95 and E(V [I) \ 0.04. on extensive studies of other open clusters and preÈmain- 0 C sequence stars, all G dwarf and at least early K dwarf 2.1. Assessment of Cluster Membership members of the cluster should have relatively strong lithium The most important initial step in the analysis of the IC 6708Ó absorption lines, whereas most Ðeld stars should 2391/IC 2602 data is the determination of which of our have undetected lithium. However, those same open cluster target stars are cluster members and which are nonmem- studies show that the low-mass members of young clusters bers. Previous membership studies in these clusters based show a large spread in lithium abundance at a given color, on proper motions and photometry extend down only to and there is a rapid decrease in lithium abundance whose about V \ 11 (early G-type stars) and so are not helpful for onset and functional form is mass and age dependentÈand the new candidates since they are almost all fainter than therefore we need to exercise caution when using lithium as this. Perhaps the most inviolate membership criteria for our a membership discriminant for our IC clusters. Finally, also stars are the positions in color-magnitude and two-color based on previous cluster studies (see, e.g.,Stau†er & Hart- diagrams because only unexpectedly large photometric mann1987), we expect the late-K and M dwarf cluster errors or rare or unexpected physical processes (very long- members to all show Ha in emission, while the G dwarfs are lived circumstellar disks, Ñares, or white dwarf companions) likely to have Ha in absorption though ““ Ðlled in ÏÏ relative could cause a star to deviate signiÐcantly from the cluster to older G dwarf Ðeld stars. locus in one of these diagrams. These tests are not generally As is indicated above, we have many useful checks on the likely to exclude many of the current targets, however, membership status of our target stars. And, in most cases, because their selection for spectroscopic observation was all of the criteria agree, and the designation of a star as a based on their having photometry consistent with cluster member or a nonmember is a simple matter. However, for membership. The internal radial velocity dispersion of these some stars, this is not the case, and then the interpretation is clusters is expected to be less than 1 km s~1, and therefore not so clear. For instance, it is plausible that a small number cluster members should share a common radial velocity to of our X-ray sources may be young, active, Ðeld stars rather No. 2, 1997 YOUNG OPEN CLUSTERS IC 2391 AND IC 2602 781

TABLE 2 CTIO SPECTROSCOPY OF STARS IN IC 2602

Stara V (V [I) log L /L v sin iv Ha E.W. Li? Remarksb C X bol helio B56 ...... 9.98 0.67 \[5.32 \6 19.1 [1.4 N n.m.? B57 ...... 10.26 0.76 \[5.17 \6 19.2 [1.3 N n.m.? B93 ...... 10.13 0.80 \[5.72 14.0 36.5 [1.5 N n.m. B94 ...... 10.84 0.79 \[5.34 17: 33.4 [1.5 N n.m. B104 ...... 9.98 0.72 \[5.47 [ 150 . . . [2: . . . n.m.? W79 ...... 11.57 0.85 \4.31 8.0 17.3 [0.9 Y R15...... 11.75 1.06 [2.92 7: 17.4 ]0.2 Y R24a ...... 14.61 1.86 [3.07 34. 18.7 ]1.9 Y R26...... 15.14 2.15 [3.62 \6 17.7 ]1.0 . . . R27...... 14.35 1.80 [3.53 10.0 17.0 ]1.2 . . . R31...... 15.08 2.24 [2.97 35. 16.3 ]3.9 . . . R32...... 15.06 2.16 [3.20 9.0 18.1 ]1.9 . . . R38...... 15.72 2.50 [3.11 48. 17.3 ]11.3 . . . R44...... 14.88 2.03 [3.62 6: 17.2 ]1.4 . . . R50...... 14.75 2.08 [3.18 6: 16.7 ]1.4 . . . R53b ...... 15.39 2.49 [3.11 100. 17.1 ]5.1 . . . R57...... 15.59 2.44 [3.30 \7 19.1 ]3.2 . . . R66...... 11.07 0.83 [3.76 12.0 17.3 [1.0 Y R67...... 14.97 2.16 [3.41 \6 16.4 ]1.5 Y R70...... 10.92 0.71 [4.44 11.0 16.7 [1.2 Y R71a ...... 14.60 2.11 [3.71: \6 [31.2 [0.8 N n.m. R71b ...... 15.34 3.41 [3.90: \6 31.8 [0.7 . . . n.m. R77...... 14.12 1.72 [3.70 \7 17.6 ]0.4 . . . R82...... 14.98 2.51 [3.17 \7 14.6 ]3.4 . . . SB3 R88a ...... 12.71 1.35 [3.56: D200 . . . D0. ... R88b ...... 12.79 1.47 [3.67: \6 16.7 [1.0 N n.m.? R93...... 13.79 1.62 [3.58 8.5 18.5 ]0.6 Y? R96...... 12.94 1.37 [3.19 17.0 15.5 ]0.6 Y R106 ...... 14.22 1.88 [4.02 \6 [29.4 [1.1 Y n.m.?

a Star name fromRandich et al. 1995, Braes 1962, and Whiteoak 1961. Note that B56 \ HD 93796, B57 \ CPD [62 1760, B93 \ HDE 307922, B94 \ HDE 307924, and W79 \ GSC 8965.0599. b REMARKS.Èn.m. \ nonmember; n.m.? \ possible nonmember; SB1? \ possible single-lined spectroscopic binary; SB2 \ double-lined spectroscopic binary. See text for details. Special casesÈR106: this star has relatively strong lithium absorption (unusual since even other members of this color have no detectable lithium) but also Ha is strongly in absorption. Examination of the spectrum shows that it is considerably earlier spectral type than indicated by its V [I color, presumably due to reddening. We estimate E(V [I) D 0.5. This helps to explain the strong lithium absorption, though other IC 2602 members even at the dereddened color still have Ha in emission. We believe this star is probably a background, young Ðeld star and not a member of IC 2602. R82 (SB3): spectroscopic triple system. Listed values are for the primary. At the time of our observation (JD 2,449,733.82), the other two components had radial velocities of [22.5 and 57.0 km s~1. The secondary and tertiary appear to have v sin iÏs slightly above our detection limitÈwe estimate v sin i D 8È10 km s~1. than cluster members. In this case, they could pass the example, the stars B56 and B57 inTable 2 are possible lithium and Ha criteria but fail the C-M diagram and radial cluster members based on photometry compatible with velocity tests. A mid- to late Ðeld K dwarf at approximately membership and radial velocities that are only slightly dis- the cluster distance could be difficult to classify because crepant. However, they also have no detectable lithium neither lithium nor Ha are deÐnitive indicators in this color when we would expect strong lithium absorption features, range, and a single, discrepant radial velocity could still be they have only X-ray upper limits when other members of compatible with cluster membership since the star could be their colors are detected, and their spectroscopic rotational a spectroscopic binary. velocities are below our detection limit in a color range We have made Ðnal membership assessments for each of where all of the unambiguous cluster members are relatively our program stars after careful consideration of all of the rapid rotators. We thus consider these two stars to be prob- relevant data; these assessments are indicated in the column able nonmembers. Notes on a few other ambiguous cases headed Remarks in Tables1È3. If that column is left blank, are provided below the tables. we consider the star to be a member of the cluster, which means that it passes all of the applicable membership tests; 3. DISCUSSION stars that fail all of the spectroscopic criteria are designated as nonmembers (n.m.). Double- (or triple-!) lined spectro- 3.1. Completeness of Our Sample scopic binaries whose average radial velocity matches the In the following several sections, we will use the stars we cluster velocity and that pass the other membership criteria have identiÐed as probable members of the IC 2391 and IC are indicated as SB2 or SB3 members; stars with discrepant 2602 clusters in order to determine physical properties of radial velocities but that pass the other relevant member- these stars and their parent clusters. In order to use cluster ship tests are designated as ““ SB1? ÏÏ and are considered to samples for this purpose, it is desirable to have started with be cluster members during subsequent analysis. Finally, a complete and unbiased membership list. That is certainly objects are marked as ““ n.m.? ÏÏ if we believe it is likely the not the case here. In neither cluster does our ROSAT survey star is a nonmember but the evidence is ambiguous. For cover the entire area of the sky over which cluster members 782 STAUFFER ET AL. Vol. 479

TABLE 3 ESO SPECTROSCOPY OF STARS IN IC 2602

Stara V (V [I) log L /L v sin iv Ha E.W. Li? Remarksb C X bol helio R3...... 11.32 0.89 [3.14 25. 21. [0.2 Y SB1? R5a ...... 9.56 . . . [4.16 . . . 65: [1.2 . . . ? R5b...... 9.68 . . . [3.97 ...... [2.0 . . . n.m. R7...... 9.21 0.46 [4.70 52. 22: [1.5 Y SB1? R12 ...... 9.48 1.20 [5.20 13. [13. [1.2 . . . n.m. R14 ...... 11.57 0.88 [3.40 13. 18. [0.25 Y R15 ...... 11.75 1.06 [2.92 10. 15. ]0.2 Y R18 ...... 12.49 1.20 [3.75 11. 13. [0.9 N n.m.? R21 ...... 9.50 0.62 [3.83 23. 13. [1.35 Y R29 ...... 12.73 1.19 [3.31 22. 17. ]0.5 Y R35 ...... 10.59 0.70 [4.09 20. 16. [1.05 Y R41E ...... 8.86: . . . [4.13: 62. 16. [1.35 . . . R41W ...... 8.86: . . . [4.13: 55. 19. [1.8 . . . R43 ...... 12.14 1.10 [3.05 50. 14. ]0.6 Y R45 ...... 10.73 0.72 [4.29 23. 12. [1.15 Y R49 ...... 11.71 0.94 [4.09 \10 [12 [1.2 . . . n.m. R52 ...... 12.19 1.11 [3.42 95. 15: ...... R53a ...... 14.04 1.85 [3.38 10. 22. [0.7 . . . n.m.? R56 ...... 13.64 1.60 [2.99 17. 14. ]0.9 Y R58 ...... 10.62 0.77 [3.22 93. 14. [0.65 Y R59 ...... 11.86 1.00 [3.13 34. 16. ]0.25 Y R66 ...... 11.07 0.83 [3.76 10. 16. [0.95 Y R68 ...... 11.32 1.09 [3.03 48. 17. ]0.3 Y R70 ...... 10.92 0.71 [4.44 10. 17. [1.05 Y R72 ...... 10.96 0.76 [3.01 49. 16. [0.6 Y R73b ...... 11.06 1.29 [4.28 \10 [7. [1.15 . . . n.m. R79 ...... 9.08 0.51 [4.38 57. 16. [1.8 Y R80 ...... 10.66 1.03 [3.69 10. 20. [0.5 Y n.m.? R83 ...... 10.70 0.78 [3.55 30. 18. [0.8 Y R85 ...... 9.87 0.58 [4.74 45. 14. [1.6 Y R88a ...... 12.71 1.35 [3.54 [ 150 ...... R88b ...... 12.79 1.47 [3.65 ...... n.m.? R89 ...... 12.97 1.35 [3.54 14. 17. ]0.2 Y R92 ...... 10.28 0.78 [3.72 14. 15. [0.8 Y R94 ...... 13.33 1.73 [3.56 23. 16. ]0.9 Y R95a ...... 11.73 0.97 [2.96 12. 16. ]0.45 Y

a Star name fromRandich et al. 1996a. b REMARKS.Èn.m. \ nonmember; n.m.? \ possible nonmember; SB1? \ possible single-lined spectroscopic binary; SB2 \ double-lined spectroscopic binary. See text for details. Special casesÈR53A (n.m.?): this star has photometry consistent with membership but a radial velocity somewhat o†set from the cluster mean and Ha in absorption at a color when we expect members to all have Ha emission. R80 (n.m.?): the photometry and radial velocity indicate this is not a member, but it has a relatively strong lithium feature. If the photometry is accurate, it would have to be at least a triple system to be consistent with cluster membership or a considerably younger than average binary system (i.e., it is located D1.1 mag above the cluster locus in a C-M diagram). If it is a nonmember, it must be a young, foreground star. are likely to be distributed. The X-ray Ñux limits certainly member list will bias our sample toward more coronally also prevent us from identifying the lowest mass members of active stars. Prior to a more detailed discussion of our the cluster, assuming that there is an upper bound to the results, it is therefore useful to attempt to assess the degree ratio of X-ray to bolometric luminosity. More importantly, to which this bias may a†ect our analysis. however, it is necessarily true that our use of an activity Our derived membership list will not be biased if all indicator as the original source of our candidate cluster members of the two clusters have X-ray Ñuxes above the detection limit for our survey. Since our X-ray Ñux limit is TABLE 4 not a single number but varies spatially over the region Ha EQUIVALENT WIDTHS surveyed, there is no single answer to this question. FOR SHJM STARS However, if the cluster member X-ray Ñuxes are above our detection limit over the region from which most of our Star Ha Eq. Width (Ó) members were identiÐed, then it seems likely at least that SHJM 1 ...... [1.1 our membership list will not have been signiÐcantly a†ected SHJM 2 ...... [1.3 by this bias. SHJM 3 ...... 0.0 The simplest and most direct test of the completeness of SHJM 4 ...... 1.9 SHJM 5 ...... 2.2 our X-ray selection method can be made for IC 2391. We SHJM 6 ...... [0.4 have previously used proper motions, photometry, and SHJM 7 ...... [0.9 spectroscopy to identify 10 probable members of IC 2391 SHJM 8 ...... 1.4 (Stau†eret al. 1989). That membership list is, by deÐnition, SHJM 9 ...... 0.8 unbiased relative to activity. The ROSAT survey of IC 2391 SHJM 10 ...... 3.5 detected nine of the 10 proposed members(Patten & Simon No. 2, 1997 YOUNG OPEN CLUSTERS IC 2391 AND IC 2602 783

1993), and the one ““ member ÏÏ not detected in X-rays is star would be predicted to be 13% brighter in X-rays com- probably actually a nonmember (it has no detectable pared to the 30 Myr star. Given the total range of X-ray lithium line at 6708Ó, whereas the other members of the Ñuxes seen in young clusters, di†erences of this order are cluster do have measurable lithium at that colorÈRandich not likely to signiÐcantly bias our survey. et al.1996a). This is admittedly a small sample, but it suggests that our X-rayÈderived membership lists are not 3.2. PMS Isochrone-Ðtting Ages for IC 2391 and IC 2602 extremely biased. One of the original reasons for our interest in these two We can estimate our bias by comparison to the X-ray clusters was that their ages as estimated from their upper survey for another young cluster (the Pleiades) for which main-sequence turno†s (D35 Myr) suggested that they both X-ray data and a good membership list from proper could provide data capable of empirically testing PMS motions are available. The ROSAT PSPC data for both IC theoretical evolutionary tracks. That is, these clusters are clusters are more sensitive toward the center of the cluster, old enough that one expects that both circumstellar gas and with a point-source detection limit at the cluster centers of dust disks and circumcluster gas and dust should have dissi- about log(L ) \ 28.0. For the IC 2602 survey, for a region pated(Skrutskie et al. 1990; Leisawitz, Bash, & Thaddeus correspondingX to half of the area surveyed, the limiting 1989) and thus that it should be possible to place the cluster sensitivity was log(L ) D 28.7. Based on the X-ray survey of members accurately on a theoretical H-R diagram, yet they the core of the Pleiades,X where proper motion membership are also young enough that the PMS turn-on point should lists are expected to be complete, all Pleiades G dwarfs have be at D1.0M_ and the PMS isochrone should be well log(L ) [ 28.9. Assuming only that G dwarfs do not spin displaced above the ZAMS. Therefore, it should be possible up onX the main sequence and that there is a positive corre- to compare accurately the empirical PMS isochrone to the lation between rotation and coronal activity, at 35 Myr the theoretical models and to use that comparison to derive an Pleiades G dwarfs would have been stronger X-ray sources, estimated contraction age for the low-mass members of the and thus a survey with limiting sensitivity of log (L ) \ 28.7 clusters. If enough stars can be identiÐed as cluster members would be complete. Therefore, we expect that for GX dwarfs and if accurate photometry can be obtained, the cluster in the central 50% of the region surveyed in IC 2602Èwith data could also in principle be used to place constraints on a similar argument holding for IC 2391Èwe should have the age spread among the low-mass stars in the clusters. obtained an unbiased cluster membership list. In the outer First, it is useful to examine the color-magnitude dia- portions of the survey area, we will be biased toward more grams for the two clusters qualitatively.Figure 4 (top) active G dwarfs, and thus our sample of G dwarfs will be shows the V versus (V [I) C-M diagram for the stars of slightly weighted toward rapid rotators. However, since for both clusters, with the known and suspected spectroscopic more than half of the region we expect to be complete, our binaries indicated. The Ðgure illustrates several points: (1) full sample should cover the entire rotational velocity range the distance, reddening, and ages of the two clusters must be present for G dwarfs in the two clusters. very similar (or di†erences in one parameter must be for- For K and M dwarfs, we cannot make the same argu- tuitously compensated for by di†erences in another ment, primarily because at the age of IC 2391 and IC 2602 parameter) since there is no discernible di†erence in the stars in this mass range are still evolving moderately rapidly locus of stars for the two clusters; (2) the SB2 and SB3 on their PMS tracks, and thus signiÐcant spin-up of their cluster members are displaced D0.7 mag above the locus of rotational velocities may occur between the age of the IC the other stars, which indicates nearly equal mass com- clusters and the age of the Pleaides. The expected sense of ponents, as must be true for binaries where the lines of each the bias would be that we would not detect the least active star are seen; and (3) the suspected single-lined binaries are cluster members, which presumably would correspond to not displaced signiÐcantly relative to the other starsÈalso the most slowly rotating members. A small piece of evidence as expected, given that the lines of the secondary are not that suggests this is not a huge e†ect is that the ratio of slow seen in the spectrumÈbut supportive of cluster member- to rapid rotators (arbitrarily split at 20 km s~1) is the same ship for these stars despite their discrepant radial velocities. for the Pleiades and IC cluster K and M dwarfs (see Fig. Figure 4 (bottom) shows a similar C-M diagram, but this 10); if we were signiÐcantly biased against slow rotators, time with di†erent symbols for slow and rapid rotators. One presumably the ratio for the IC cluster sample would be option that has been considered for the wide range in rota- smaller than for the Pleiades stars. Nevertheless, some bias tional velocities found at a given mass in young open clus- presumably is present for stars in this mass range, and this ters is that there is a signiÐcant age spread within the cluster should be kept in mind. and that the rotational velocity spread is simply a reÑection There is another possible bias that could enter our of the age spread; for IC 2391/2602, this is obviously not the survey, and that is we could be more (or less) sensitive to the case, since the slow and rapid rotators follow the same older stars in the IC clusters, on the assumption that there is empirical PMS isochrone to within the errors of our mea- a detectable age spread among the members of these two surement. clusters. For plausible age spreads, we believe that this e†ect A more quantitative assessment of our IC cluster photo- must be negligible. Consider two 0.6M_ cluster members, metry can be made by comparing the empirical PMS locus one 30 Myr old and the other 40 Myr old. The older star to theoretical isochrones. To do this, we use the same pro- will have evolved further down its PMS track and will thus cedures outlined previously to estimate the age for HR 4796 have a lower bolometric luminosity (by about 15%). (Stau†er,Hartmann, & Barrado 1995, hereafter SHB) and However, it will have also contracted and will thus have to compare photometry of low-mass stars in the Pleiades to spun up (also by about 15%, using evolutionary tracks from theoretical isochrones(Stau†er 1996). BrieÑy, we use the Vandenberg 1987). IfL /L P v, the two e†ects just cancel, evolutionary models ofDÏAntona & Mazzitelli (1994, here- and the predictedL forX thebol 40 Myr old star would be the afterDÏA&M), which employ Alexander & Rogers opacities same as for the 30 MyrX star. IfL /L P v2, then the older and the Canuto & Mazzitelli turbulent convection model X bol 784 STAUFFER ET AL. Vol. 479

IC2391 + IC2602 4 IC2391/IC2602

10 SB1? SB2 or SB3 IC2391 IC2602 6 12 V

14 8

16 10

.5 1 1.5 2 2.5 .5 1 1.5 2 2.5

IC2391+IC2602 4 Pleiades

10 vsini < 30 km/s vsini > 30 km/s

6 12 V

8 14

16 10

.5 1 1.5 2 2.5 .5 1 1.5 2 2.5

FIG. 4.ÈV vs. (V [I) C-M diagrams for stars in IC 2391 and IC 2602. FIG. 5.ÈH-R diagrams for IC 2391/2602 and the Pleiades compared to The top panel allows comparison of the location of stars in IC 2391 relative theoretical isochrones fromDÏAntona & Mazzitelli (1994)Èsee the text for to stars in IC 2602 in the C-M diagram, as well as indicating the location of a more detailed discussion of the isochrones. The top panel shows the the known and suspected spectroscopic binaries. The bottom panel shows comparison for the IC clusters, where we have assumed (m [ M) \ 6.05 the same stars, but this time allowing a comparison of the location of stars andE(V [I) \ 0.01 for IC 2391, and (m [ M) \ 5.95 and E(V0[I) \ with projected rotational velocities greater than or less than 30 km s~1. 0.04 for IC 2602.C The three isochrones are 3, 10, and0 30 Myr, and the solidC line is the ZAMS from DÏAntona & Mazzitelli. The bottom panel provides the similar plot for the Pleiades, where we assume(m [ M) \ 5.54 and E(V [I) \ 0.06, and isochrones for 3, 10, 30, and 70 Myr are0 shown. The C thick dashed lines are evolutionary tracks for 0.7 and 0.4M_; note also (speciÐcally, Table 1 ofDÏA&M), and bolometric correc- that the highest mass point on the isochrones corresponds to 1.0 M . tions fromSchmidt-Kaler (1982) and Bessell (1991). To _ convert from e†ective temperature to (V [ I) color, we use the ““ tuned ÏÏ temperature scale advocated inSHB, which results in a close Ðt of the empirical Pleiades isochrone to mag would follow the lower envelope of the IC distribution the 70 Myr theoretical isochrone. Over the color range of fairly wellÈand would correspond to a best-Ðt PMS con- our IC cluster data, this temperature scale is quite close to traction age of D25 Myr. This is then about 10 Myr theBessell (1979) temperature scale forT [ 4000 K and younger than the D35 Myr nuclear age for these clusters to theKirkpatrick et al. (1993) temperatureeff scale for T \ advocated by Mermilliod (1981). 4000 K. Because this combination of models and conver-eff Should one conclude from the above that the low-mass sions from the theoretical to the observational plane has stars in these IC clusters are younger than the high-mass been calibrated for photometry on the same system for a stars? Probably not. First, it is likely that systematic errors slightly older cluster, we believe this should be the best (or at least di†erences) between the models used to derive possible way to age-date the low-mass stars in the IC clus- the nuclear ages and those used to derive the contraction ters (and certainly better than use a set of tracks and con- ages plus additional systematic errors in transforming from versions that has not been so ““ validated ÏÏ). the theoretical to the observational plane could account for Figure 5 (top) shows the IC photometry, now trans- age di†erences of this order. For example, we calibrated the formed toMV and(V [I)c , compared to the DÏAntona & color-temperature conversion using an assumed age of 70 Mazzitelli 3, 10, and 30 Myr,0 isochrones and their ZAMS. Myr for the PleiadesÈif, instead, the Pleiades age is of The detailed agreement between the shape of the empirical order 100 Myr(Basri, Marcy, & Graham 1996), this would isochrone and the theoretical isochrones indicates, we modify the color-temperature conversion slightly in the believe, that the calibration of the tracks discussed above sense of making the inferred age for the IC clusters slightly was successful. Of the isochrones published byDÏA&M, the older and hence more in agreement with the nuclear age. closest Ðt is to the 30 Myr isochrone; however, a slightly Errors or di†erences in the assumed cluster distance moduli brighter (younger) isochrone would probably be better. Spe- and reddening could also a†ect the derived ages. Finally, ciÐcally, an upward shift of the 30 Myr isochrone by D0.1 because the IC stars are expected (andknownÈPatten & No. 2, 1997 YOUNG OPEN CLUSTERS IC 2391 AND IC 2602 785

Simon1996) to be very spotted, their mean V [I color at a the spectroscopic rotational velocity data for the cluster given mass or the conversion from V [I color toT could stars with resolution limits of 7È10 km s~1 are not capable be systematically di†erent from that for a quiescenteff star of of independently determining the functional form at low the same mass and luminosity. Our conclusion from Figure rotational velocities. At high rotational velocities 5 is that the PMS and nuclear ages for the two IC clusters (v sin i [ 15È20 km s~1)Èa region not well sampled by agree to within the expected errors of the calibrations. The single, Ðeld starsÈthe cluster data indicate that the activity new data are sufficiently good, however, to suggest that the indices plateau at a relatively Ðxed upper bound generally noncoeval star formation model outlined by Herbig (1962) referred to as the saturation level(Vilhu 1984). This general and developed in greater detail byNorman & Silk (1980) behavior is illustrated for the Pleiades inFigure 6 (top), with whereby the low-mass stars in a cluster form Ðrst and star ROSAT PSPC data as the activity index(Stau†er et al. formation is eventually shut o† by a burst of high-mass star 1994). The Ðgure also illustrates that for stars slightly more formation is not evident in these two IC clusters. massive than the Sun,L /L is no longer obviously corre- At least in principle, it should also be possible to estimate lated with the surface rotationX bol rate, a result most easily the age spread among the low-mass stars in IC 2391/2602 interpreted as being due to the rapid decrease in the thick- from the dispersion of stars about a single isochrone in ness of the outer convective envelope for masses greater Figure 5 (top). The fact that most stars are members of than 1M_ and the corresponding difficulty in maintaining binary systems makes this much more difficult in practice, a solar-like magnetic dynamo (seeRandich et al. 1996b for a as is illustrated by examination of the Pleiades C-M more thorough discussion). diagram inFigure 5 (bottom). Because of direct evidence Figure 6 (bottom) shows the same plot of coronal activity (i.e., the identiÐcation of spectroscopic binaries; see versus spectroscopic rotational velocity for our IC clusters Mermilliodet al. 1992) and indirect evidence (the morpho- stars, also split into the same two color intervals. The IC logical similarity of the Pleiades C-M diagram to C-M dia- cluster stars have the same general characteristics in this grams of much older clusters such as Praesepe or the diagram as illustrated by the Pleiades membersÈsaturation Hyades, where no plausible age spread should have an atL /L D [3 for rapid rotators, lower coronal activity e†ect), we assume that the displacement of stars well above for slowX bol rotators, and little apparent correlation between the single-star main sequence in the Pleiades is due to the rotation and activity for stars more massive than the Sun. presence of unresolved, approximately equal mass compan- However, there are some di†erences also. The saturation ions, and that a signiÐcant portion of the scatter about the level appears to decline for very large rotational velocities single-star MS is due to companions of lower mass. Visual (v sin i [ 100 km s~1), something not seen in the Pleiades comparison ofFigure 5 (top)to Figure 5 (bottom) indicates that there is more dispersion of the low-mass stars about a locus line for IC 2391/2602 than for the Pleiades, possibly pointing to an age spread for stars in the IC clusters. However, other e†ects are also possible contributors to this PLEIADES dispersion: (1) di†erential errors in the distance moduli or -3 reddening for the two clusters; (2) greater variability or spottedness in the younger IC clusters; (3) a real di†erence in the mean ages of the two IC clusters; or (4) a possibly greater binary frequency for stars in the IC clusters com- -4 pared to the Pleiades. Unfortunately, in practice we believe that we are at present unable to derive a useful constraint on the age -5 spread in the two clusters because of our inability to quan- tify the other sources of spread in the C-M diagram indicated above. We have constructed Monte Carlo simulations -6 of the cluster photometry in which we ascribe all additional 0 50 100 150 200 250 spread to an age spread, and we have determined that age vsini spreads up to 20 Myr (this is, a range of ages for the cluster members from 20 to 40 Myr) are allowed. We suspect that IC2391 & IC2602 the real age spread is much smaller than this, but proving -3 that requires both a larger sample of stars and some means to estimate the role of di†erential reddening, spottedness, etc. -4 3.3. Chromospheric and Coronal Activity For clusters that have been well studied previously, it has been shown that for low-mass stars, there is a good corre- -5 lation between chromospheric/coronal activity and rotation (Simon 1990; Soderblomet al. 1993; Stau†er et al. 1994;

Randichet al. 1996b). At low rotational velocities, the corre- -6 lation presumably follows the functional form found for 0 50 100 150 200 250 Ðeld stars, activity Pvn with n between 1 and 4 (Walter vsini 1982; Vilhu 1984;though see also Simon, Herbig, & Boes- FIG. 6.ÈCoronal activity vs. v sin i for low-mass stars in the Pleiades gaard1985, for example, for alternative functional forms); (top) and in IC 2391/2602 (bottom). 786 STAUFFER ET AL. Vol. 479 for lack of any stars that sample this rotation range but that has been seen previously for members of the younger a Persei cluster(Randich et al. 1996b; Prosser & Grankin 1996). We are unaware of any explicit physical prediction of this e†ect, though possible hints for such behavior might be found in the expectation for higher order multipole magnetic Ðelds for very rapid rotation and the distribution of hot gas in extended coronal structures under the competing inÑuence of gravity and centrifugal acceleration for such rapid rotators(Mestel & Spruit 1987; Taam & Spruit 1989). As a second apparent di†erence, the slowly rotating Pleiades low-mass stars seem to reach lower coronal activity levels than do the IC cluster stars, and there is less scatter in the activity levels (that is, the Pleiades stars fall more or less along a single, albeit highly nonlinear, curve, whereas the IC cluster stars are not so well behaved). The relative paucity of very low activity,(V [I) [ 0.70 stars in the IC clusters could either be due to a realC lack of G-star slow rotators (with such stars being present in the Pleiades owing to angular momentum loss on the main sequence) or be due to the selection bias in our survey against low- activity stars. Since most of the low-activity stars in the Pleiades plot are slowly rotating G dwarfs, and for stars of FIG. 7.ÈCoronal activity vs. Rossby number, log (P/q), where P is the that mass we believe our survey is nearly complete for the starÏs rotation period and q is the convective turnover time as estimated by IC clusters, we believe that the lack of low-activity stars in eq. (4) ofNoyes et al. (1984) for stars in the two IC clusters ( Ðlled circles) the IC plot is primarily a physical e†ect, though some con- and stars in a Persei (open circles). Periods have been estimated from the measured v sin iÏs using stellar radii derived from a radius-color relation tribution from our selection bias is probably also present. appropriate to the 30 MyrDÏA&M isochrone, and we have assumed a The larger amount of ““ scatter ÏÏ in coronal activity for mean inclination correction of (4.0/n). 10 \ v sin i \ 30 km s~1 among the(V [I) [ 0.70 IC cluster stars relative to those in the Pleiades is plausiblyC due to the mass dependence of the saturation velocity and the di†erence in the rotational velocity distribution of the IC cluster members and the a Persei members in a Rossby-style clusters relative to the Pleiades, as we will discuss in the plot. The IC cluster members track the locus of the a Persei next paragraph. members quite well, and the F dwarfs in the IC clusters An alternative means to display the correlation between follow the relation deÐned by the lower mass stars as well. rotation and coronal activity is to use the Rossby numberÈ Based on the data for the two clusters, the Rossby number the ratio of the rotation period to the convective turnover at which the relation begins to saturate [log (R) \ log timeÈas the measure of rotation (seeNoyes et al. 1984; (P/q) D [0.8 ^ 0.2] appears to be at least roughly indepen- Simon,Herbig, & Boesgaard 1985; Soderblom et al. 1993; dent of mass. A similar result can be seen in Figure 12 of Randichet al. 1996b). When plotting Ca H and K chromo- Patten& Simon (1996), where Ðeld stars and stars in other spheric activity data, Noyes et al. showed that use of the clusters are included. Using equation (4) ofNoyes et al. Rossby number allowed one to combine stars of a wide (1984) to deÐne the convective turnover time, the saturation range of mass in a single activity versus rotation plot with point occurs at P \ 2d or v(rot) D 25 km s~1 at about 1 much less scatter than when rotation period alone was used. M_ and at P \ 4.5d or v(rot) D 5kms~1 at about 0.4 M_. Ideally, one would like to use stars with measured rotation A quantitatively similar prediction for the mass dependence periods to construct the Rossby index; however, that is of the saturation velocity has been suggested byBarnes & difficult for the open cluster samples generally because only SoÐa(1996) based solely on an attempt to explain the rota- a small fraction of the stars have measured periods. Patten tional velocity data for low-mass stars in a Persei, the & Simon(1996) Ðnessed this problem by combining stars Pleiades, and Hyades. Given this mass dependence of the with periods from a number of clusters (including IC 2391) saturation velocity, one therefore expects ““ scatter ÏÏ for and also including Ðeld stars (which deÐne the long-period 10 \ v sin i \ 25 km s~1 in a plot ofL /L versus v sin i end of the relation). We have chosen another route, which is if one has stars with a range of masses, whichX bol is true for the to estimate the Rossby number from the measured v sin iÏs IC clusters. There are very few Pleiades G and early K and our photometry (to estimate the stellar radii)Èthis dwarfs with rotational velocities in this range for which allows us to include a large sample of stars, even in the IC X-ray data are available (see Figs. 16 and 17 ofStau†er et clusters, at the expense of some imprecision introduced by al.1994), and so it is plausible that it is simply the lack of the lack of knowledge of sin i and the errors in determining stars in the appropriate velocity and mass range that allows the stellar radii. With this technique,Randich et al. (1996b) the Pleiades Ðgure(Fig. 6 [top]) to show less scatter. showed that for the a Persei open cluster, even the F dwarfs Ha has previously been used as a chromospheric activity follow the same coronal activity-rotation relation as the proxy for the low-mass stars in the Pleiades(Stau†er & low-mass stars if one uses Rossby number as the abscissa, Hartmann1987; Soderblom et al. 1993), Hyades (Stau†er et whereas the F dwarfs do not show an obvious correlation of al.1991), and a Persei open clusters (Prosser 1994). These activity and rotation if one uses v sin i as the independent and other studies have shown that even quite young G variable.Figure 7 compares the coronal data for the IC dwarfs have Ha in absorption, though apparently ““ Ðlled No. 2, 1997 YOUNG OPEN CLUSTERS IC 2391 AND IC 2602 787 in ÏÏ relative to older G dwarfs. Young M dwarfs generally the other clusters and extend that trend to a younger age. have Ha in emission, with the average equivalent width We now return brieÑy to the one (or two) stars that increasing to later spectral types.Soderblom et al. (1993) showed broad wings to their Ha emission proÐles. What is showed that by plotting a normalized Ha Ñux versus the explanation for this phenomenon, and why is it Rossby number, one gets a relation very similar to that exhibited by only these two stars in our sample? We believe shown forL /L and with a very similar saturation point. that the most likely explanation for the broad Ha com- Stau†eret al.X (1991)bol suggested that it might be possible to ponent is that it arises from macroscopic turbulence con- use the color at which Ha becomes detectable as an emis- nected with Ñares or microÑares. Similar broad Balmer sion feature as an age indicator for open clusters. Our new emission wings have been seen in dMe stars connected with data for IC 2391/2602 provide further evidence that this Ñare events, with the broad component sometimes blue- may be true. shifted and sometimes redshifted relative to the narrow Figure 8 shows the distribution of Ha equivalent widths component(Eason et al. 1992; Phillips et al. 1988; Doyle et for the low-mass stars of IC 2391/2602, the Pleiades, and the al.1988). Moreover, in a recent survey of chromospheric Hyades. We have used the formula provided byBessell & activity for M dwarfs in the Pleiades, we have found similar Weis(1987) to convert the measured Kron system V [I broad wings in at least two of D15 dMe stars observed. In colors for Pleiades and Hyades stars to the Cousins system. all cases, the FWHM of the broad component is very The Ðgures show that there is indeed a clear age dependence similar and is of order 250 km s~1. In a somewhat di†erent to both the color at which Ha Ðrst appears in emission and context,Linsky & Wood (1994) have reported broad com- the color at which all stars have Ha in emission. There is ponents for the UV lines of C IV and Si iV in the dMe star also good evidence in these plots for a correlation between AU Mic, with FWHM similar to the broad components we rotation and activity, in that at a given color, the slowly see for Ha in our stars. They attribute the broad component rotating stars lie systematically below the rapid rotators. to plasma turbulence generated by microÑares (Dere, Particularly for the M dwarfs in the Hyades, we believe one Bartoe, & Brueckner1989; Cook 1991). If this explanation could use the Ha data as a good predictor of at least relative is correct, then the fraction of dMe stars with broad Balmer rotational velocities. A more detailed explanation of this wings is simply indicative of the frequency of Ñares (or Ðgure would require consideration of how the Ha equiva- microÑares) in the cluster population. Since the Ñare fre- lent width varies as the chromospheric temperature and quency is known to be high in the Pleiades (see Haro, density structure is changed (and as a function of stellar Chavira, & Gonzalez1982) and is expected to be as high or mass), topics that are beyond the scope of this paper. higher in a younger cluster, a detection rate of D15% for Empirically, the IC cluster stars follow the trend shown by the broad Ha components does not seem unreasonable. It would clearly be useful to obtain additional spectra in order to determine better the frequency of the broad Ha com- 6 ponents and to determine if they are transient features as IC2391 & IC2602 would be expected from the Ñare hypothesis. vsini > 8 4 vsini < 8

2 3.4. Rotational Velocities of L ow-Mass Stars at 30 Myr Our original motivation for surveying these two clusters 0 was to determine the distribution of rotational velocities for a population of D30 Myr old stars, corresponding to the

.5 1 1.5 2 2.5 arrival of solar mass stars on the zero-age main sequence and helping to Ðll the gap between older clusters like a Per 6 and the Pleiades and preÈmain-sequence T Tauri stars. PLEIADES vsini < 8 Figure 9 shows the distribution of rotational velocities for 4 vsini > 8 low-mass stars in IC 2391/2602 compared with the Pleiades. Figure 10 shows histograms of the rotational velocity dis- 2 tribution in the IC clusters for three color ranges, corresponding roughly to spectral types G, K, and M. For 0 comparison, the rotational velocity distribution for G dwarfs in the Pleiades is also shown in Figure 10. .5 1 1.5 2 2.5 Note that while we use the shorthand of just saying ““ rotational velocities ÏÏ here and elsewhere, what we have 6 really measured are the projected rotational velocities of HYADES vsini > 8 our program stars. If a star is seen nearly pole-on, the pro- 4 vsini < 8 jected rotational velocity can di†er substantially from the true rotational velocity. If many of the stars had small 2 values of sin i, then inferences drawn from our spectroscopic rotational velocities about the true rotational veloc- 0 ity distribution might be in error. However, the available evidence suggests that the rotational axes of stars in open .5 1 1.5 2 2.5 clusters are randomly oriented (seeKraft 1970), and if that is the case, then the distribution of projected rotational FIG. 8.ÈChromospheric activity as a function of V [I color for IC velocities should di†er in only minor ways from the true 2391/2602, the Pleiades, and the Hyades. rotational velocity distribution. In particular, since the 788 STAUFFER ET AL. Vol. 479

rotators, it is very difficult to draw any conclusions about 200 IC2391(new) the K and M stars other than that there are no gross di†er- IC2391(old) ences between the Pleiades and IC cluster velocity distribu- IC2391(limit) IC2602(ctio) tions. 150 IC2602(eso) Theoretical treatments of the rotational velocity evolu- IC2602(limit) tion of solar-type stars are complicated by the need to include wind torques, possible decoupling between radi- vsini 100 ative envelopes and convective cores, and likely spin-down by the interaction of the stellar magnetosphere with the circumstellar disk during the preÈmain-sequence phase (see, 50 e.g.,Keppens et al. 1995; Cameron, Campbell, & Quaintrell 1994; and references therein). None of these processes are

0 well understood, and often the only real constraints on .5 1 1.5 2 2.5 these processes arise from Ðtting the (incomplete) rotational velocity data, so that it is difficult to separate these individual causes. PLEIADES measured vsini's The principal new constraint that we have found here is 150 vsini limits the evidence for spin-down in G stars of about a factor of 2 in v sin i between 30 and 70 Myr. This is in reasonably good agreement with the prediction of the ““ best choice ÏÏ model 100 for spin-down ofKeppens et al. (1995) (their Fig. 9).

vsini Although this model also predicts slightly too large rotational velocities for Pleiades stars, it manages to reproduce most of the observational constraints on the rotation of 50 solar-type stars over a substantial period of evolution. The other constraint that our data present for the angular momentum evolution of solar-type stars is that the epoch at 0 which slow rotation is observed is pushed to earlier times. .5 1 1.5 2 2.5 To illustrate the signiÐcance of this result as well as some of the issues involved in understanding angular momentum FIG. 9.ÈDistribution of rotational velocities vs. reddening corrected evolution, we show inFigure 11 the rotational velocity color for IC 2391/2602 and for the Pleiades. The IC 2391 (old) data refer to evolution of stars of 1M_ and 0.5M_ for the case of no stars fromStau†er et al. (1989); the IC 2391 (new) points are from Table 1 external torques and solid-body rotation. These results, of the current paper. The Pleiades rotational velocities are from Soderblom et al.(1993) and Stau†er & Hartmann (1987). along with the normalized rate of change of the moment of inertia I, were obtained from calculations provided by D. Vandenberg(1987). Our assumption of solid-body rotation is probably adequate to illustrate the basic problems, since probability for a star having an inclination between i and Keppenset al. (1995) Ðnd a coupling time for enforcing i ] di varies as sin idi, we would expect that half of the solid-body rotation of D10 Myr. stars should have sin i [ 0.86, 35% of them should have Figure 11 (left) shows the large expected spin-up of solar- 0.5 \ sin i \ 0.86, and that only 15% of them should have type stars in the absence of angular momentum loss by sin i \ 0.5. Since the range in projected rotational velocities stellar winds and by coupling with an accretion disk. The is greater than a factor of 20, these relatively small inclina- predicted spin-up can account for the most rapidly rotating tion corrections should not signiÐcantly a†ect the shape of G stars in the IC clusters D 100È200 km s~1, starting from the distributionÈparticularly when we compare one cluster T Tauri stars with rotational velocities D30È50 km s~1, to another. even allowing for some angular momentum loss. However, We have resolved the rotational velocity in essentially all to obtain the slow rotators in the IC clusters at D20 km of the IC 2391/2602 G dwarfs providing an important con- s~1, starting from the slow rotators among T Tauri stars at straint for theoretical models of the rotational velocity evol- D10 km s~1 (Attridge& Herbst 1992; Bouvier 1994), ution of solar mass stars. By contrast, approximately requires very strong torques at the youngest ages (Fig. 11 one-third of the G dwarfs in the Pleiades have rotational [bottom]). velocity limits, and half of the G dwarfs have v sin i ¹ 10 It appears to be difficult to account for the large torques km s~1. For the IC clusters, one has to go to 20 km s~1 to required by the slow rotators in the context of stellar wind include half of the stars, which indicates that if these stars spin-down(Keppens et al. 1995). The slow rotation rotate as solid bodies, they lose of order half their angular exhibited by accreting T Tauri stars led to the idea that momentum during their Ðrst 35 Myr on the main sequence. torques generated by the interaction of the stellar magnetic Similarly, the upper quartile of the G dwarfs in the IC clus- Ðeld with a circumstellar accretion disk cause substantial ters have rotational velocities above 45 km s~1, whereas in spin-down(Ko nigl 1991). This proposal explains why T the Pleiades, the upper quartile includes stars rotating Tauri stars without accretion disks generally rotate much above 15 km s~1. We thus conclude that the rapidly rotat- faster than stars with accretion disks(Bouvier et al. 1993; ing G dwarfs lose between 50% and2 of their ZAMS Edwardset al. 1993). Disk interactions may lead to a nearly angular momentum by Pleiades age if they3 rotate as solid Ðxed rotational period over a wide range of conditions bodies. Because of small number statistics and the possi- (Cameronet al. 1995; Armitage & Clarke 1996). bility that our IC cluster sample is biased toward rapid Typical slow rotators among T Tauri stars have No. 2, 1997 YOUNG OPEN CLUSTERS IC 2391 AND IC 2602 789

8 IC2391+2602 8 IC2391+2602

G dwarfs K dwarfs

6 6

4 4 Number Number

2 2

0 0 0 50 100 150 200 0 50 100 150 200 vsini vsini 12

IC2391+2602 Pleiades 10 20 M dwarfs G dwarfs

8 15

Number Number 10 4

5 2

0 0 0 50 100 150 200 0 50 100 150 200 vsini vsini

FIG. 10.ÈRotational velocity histograms for G, K, and M dwarfs in IC 2391/2602 and for G dwarfs in the Pleaides. The Ðlled regions correspond to stars with v sin i upper limits. v sin i D 10 km s~1 at ages of 106 yr. If evolution proceeds al.(1995) (slightly shorter than the 10È20 Myr proposed by with the disk-magnetosphere interaction producing a nearly Bouvier 1994),the spin-up suggested by Figure 11 (left) constant rotational period, then we would expect to Ðnd would result in a typical rotational velocity of D20 km s~1 rotational velocities of D6kms~1 at an age of 4 ] 106 yr. If at 30 Myr, similar to our results for the IC clusters. A more disk coupling lasts to this age, as suggested byKeppens et precise investigation of the required disk coupling lifetime

FIG. 11.ÈLeft: Predicted rotational velocity evolution of 1.0M_ and 0.5M_ stars during PMS evolution based on the models of Vandenberg (1987), assuming no external torques and solid-body rotation. Right: The corresponding rate of change in the moment of inertia during PMS evolution for stars of these massesÈmost of the evolution in the moment of inertia occurs quite early, and thus to produce slow rotators on the ZAMS, the most e†ective means is to somehow prevent spin-up during the Ðrst several Myr of PMS evolution. 790 STAUFFER ET AL. Vol. 479 will require improved rotational velocity statistics for D1 clusters. At D30 Myr, the IC clusters are at an optimum age M_ T Tauri stars. In any event, by establishing the exis- so that we could hope to place good constraints on the time tence of a substantial population of slow rotators in the IC spread of star formation within open clusters in that clusters, we have emphasized the need for a spin-down members less than about 0.8M_ are displaced well above torque in the earliest phases of stellar evolution. the ZAMS yet the cluster is old enough so that the problems associated with much younger stars (IR excesses, UV 4. SUMMARY veiling, etc.) are not expected to be present. Unfortunately, there are still enough other sources of scatter for the C-M Each new open cluster that is observed o†ers the diagram that we were unable to determine a useful limit on opportunity to either conÐrm or cast doubt upon the the possible age spread. This limitation can probably be current paradigm for the early evolution of low-mass stars. removed if enough additional observations are made of the We cannot usefully observe the evolution with time of a already identiÐed cluster members (radial velocity monitor- single cluster, and so we assume that each cluster we ing and speckle imaging to identify binaries, additional pho- observe is a snapshot in time of some hypothetical Ðducial tometry and/or more detailed analysis of the spectra to cluster. In this paper, we have reported the Ðrst obser- determine better the e†ective temperatures and luminosities vations of low-mass stars in a cluster younger than 50 Myr. of the stars). We had also hoped to discuss in more detail Our observations of IC 2391 and IC 2602 seem to fall the rotation and activity levels of the K and M dwarf into the category of conÐrming the paradigm. The nuclear members of the IC clusters but could do that only in a ages for these two clusters are D35 Myr, whereas our best- limited sense owing to the probable selection biases from Ðt PMS isochrone age is D25 Myr. However, we consider our X-ray survey for stars in this mass range. With the these two age estimates to be consistent with a true age for coming availability of mosaic CCD cameras to allow large- the clusters of D30 Myr given the possible errors in the area photometric surveys to be made and with wide-Ðeld, theoretical models and in the comparison of those models multiobject echelle spectrographs available on 4 m class to the observations. The rotational velocity distribution telescopes, it should be possible to identify an activity- and chromospheric and coronal activities for the low-mass unbiased sample of late-type stars in these two clusters with members of the IC clusters are also similar to those for the a reasonable allocation of telescope time. slightly older Pleiades cluster, with the apparent di†erences in the sense expected given that the IC clusters are younger. In short, we Ðnd no strong surprises in our observations of J. R. S. and C. F. P. acknowledge support from NASA the stars in these clusters. grants NAGW-2698 and NAGW-3690. T. S. acknowledges There are some disappointments in our study of the IC support from NASA grant NAG 5-2022.

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