Äîêóìåíò âçÿò èç êýøà ïîèñêîâîé ìàøèíû. Àäðåñ îðèãèíàëüíîãî äîêóìåíòà : http://star.arm.ac.uk/~ambn/271dcf.ps Äàòà èçìåíåíèÿ: Fri May 9 19:31:53 1997 Äàòà èíäåêñèðîâàíèÿ: Mon Oct 1 21:10:54 2012 Êîäèðîâêà: Ïîèñêîâûå ñëîâà: m 43

A&A manuscript no.
(will be inserted by hand later)
Your thesaurus codes are:
08.05.3; 08.12.1; 10.15.2 IC 2602
ASTRONOMY
AND
ASTROPHYSICS
9.5.1997
CCD photometry of late­type stars in the young open
cluster IC2602
D.C. Foster 1 , P.B. Byrne 1 , S.L. Hawley 2? , and W.R.J. Rolleston 3
1 Armagh Observatory, College Hill, Armagh BT61 9DG, N. Ireland
dcf@star.arm.ac.uk, pbb@star.arm.ac.uk
2 Dept. of Physics and Astronomy, Michigan State University, East Lansing, MI 48824­1116, USA
slh@pillan.pa.msu.edu
3 Dept. of Pure and Applied Physics, The Queen's University of Belfast, Belfast BT7 1NN, N. Ireland
R.Rolleston@Queens­Belfast.ac.uk
Received date; accepted date
Abstract. We present the results of VRI photometry of
the young open cluster IC 2602. Two 15arcmin\Theta15arcmin
fields were observed in February and May 1991 using
the 1­m Swope telescope at Las Campanas. Using theo­
retical isochrones obtained from D'Antona & Mazzitelli
(1994), and allowing for observational and other uncer­
tainties, we identify 78 primary candidate members with
12 ! V ! 18:5 from their positions on colour­magnitude
diagrams. We compare the cluster field with an offset field
of similar galactic latitude and estimate the contamination
due to background stars to be large, – 50%, as might be
expected given its low galactic latitude. We also compare
our photometry with that given for the X­ray detected
stars of Randich et al (1995) present complimentary nar­
row band Hff photometry for a subset of the stars. 1
Key words: Stars: evolution -- Stars: late­type -- Open
clusters: individual; IC 2602
1. Introduction
Ever since the pioneering work of Wilson (1963) and Kraft
(1967), it has generally been accepted that stars of solar
mass or less lose their angular momentum with time. Fur­
ther research by van den Heuvel & Conti (1971) and Sku­
manich (1972) suggested that low­mass stars arrive on the
main­sequence rotating rapidly, as a consequence of the
Send offprint requests to: D.C. Foster
? NSF Young Investigator
1 Table 3 is only available in electronic form at the CDS
via anonymous ftp to cdsarc.u­strasbg.fr (130.79.128.5) or via
http://cdsweb.u­strasbg.fr/Abstract.html
conservation of angular momentum during the pre­main­
sequence contraction phase. Drawing together both Ca ii
emission data (indicative of chromospheric activity in late­
type stars) and projected rotational velocities (used to de­
termine angular momenta) for a small number of stars in
three open clusters, viz. Pleiades, Ursa Major and Hyades,
and using the solar values, Skumanich derived empirical
relationships between the afore­mentioned properties and
age, which simply stated that both Ca ii H & K emission
reversals and stellar rotation declined with time according
to an inverse power law. Such relationships were consistent
with the theoretical predictions of Durney (1972) based on
models in which rotational braking was caused by a stellar
wind. This view remained unchallenged for more than a
decade.
More recently, new observations have given rise to a
new paradigm. Stauffer et al. (1984, 1985, 1987), follow­
ing up the discovery of rapidly rotating K stars in the
Pleiades (Van Leeuwen et al. 1987), measured rotational
v sin i for GKM­type dwarfs in the ff Persei (age¸50 Myr),
Pleiades (age¸70 Myr) and Hyades (age¸800 Myr) open
clusters. These results showed that all late­type members
of ff Per exhibit a very large spread of rotational v sin i val­
ues, ranging from approximately 15--200 kms \Gamma1 . By con­
trast, G­type Pleiads had v sin i close to or less than the
observational limit of 10 kms \Gamma1 with rapid rotation only
observed amongst the K­ and M­types.
Inter­comparison of these two clusters suggests that
a more rapid braking mechanism than that provided by
the classical stellar wind scenario must be at work. This
mechanism must be capable of braking the rotation of G­
dwarfs on a time­scale of the order of the age difference
between ff Per and the Pleiades, ie. ¸20 Myr. Further­
more, only moderately rapid rotation was detected in the
oldest of the clusters, the Hyades, and then only in the
M­type dwarfs. Comparison of these results supports the

2 D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602
idea that, once this rapid phase of braking is complete, a
power law relation may then apply.
These important conclusions have been based on the
comparison of results found for three open clusters. In or­
der to place constraints on possible braking mechanisms,
further observations are required of young clusters with
ages distributed over the critical range 10--200 Myr. Clus­
ters younger than, and of similar age to, ff Per are needed
to confirm that the rapid braking of G­dwarfs is univer­
sal, while those intermediate in age between the Pleiades
and Hyades will yield information on the time­scales for
braking of progressively lower mass stars.
The authors have undertaken such a programme to
investigate the distribution of stellar rotation in a number
of open clusters. However, as a consequence of the intrinsic
faintness of late­type dwarfs (MV ú 6 at K0, MV ú 9 at
M0), it is necessary to restrict the study to clusters that
are within approximately 400 pc of the Sun; otherwise the
measurement of v sin i from high­resolution spectroscopy
will not be observationally feasible for a sample of their
late­type members. Due to their relative proximity, these
clusters have a large extent on the sky and unambiguous
identification of bona fide members is difficult. With this
in mind, we have selected several target clusters for which
we have obtained BV RI CCD photometry. Further details
of the observational programme and background material
can be found in Rolleston (1995).
2. IC2602
The cluster IC2602 (ff = 10 h 43 m :0, \Gamma64 ffi 24', J2000.0) is a
group of stars distributed about the B0Vp star ` Carinae.
Early studies (Whiteoak 1961, Braes 1962) determined the
cluster to be at a distance of 150­155pc. They estimated
the cluster to be approximately 8­12Myr old, derived both
from the nuclear age of ` Car, and the contraction of clus­
ter members on to the main sequence. More recent studies
(Mermilliod 1981) estimate the age to be nearer 36Myr.
The reddening has been determined to be EB\GammaV =
0:04 (Whiteoak 1961, Braes 1962) corresponding to
EV \GammaI = 0:044 (Randich et al 1995). Photometric measure­
ments for the brighter stars were obtained (Braes 1962,
Whiteoak 1961 and Hill & Perry 1969), and some of these
stars have been studied spectroscopically (Whiteoak 1961,
Abt & Morgan 1972). From measurements of 22 confirmed
members, Braes (1962) has estimated the mean cluster
proper motion to be 15¯ ff cos ffi = \Gamma8:8 \Sigma 1:0mas/yr,
¯ ffi = 3:5 \Sigma 1:5mas/yr.
More recently the cluster has been studied in the X­ray
wavelength region (Randich et al. 1995). The X­ray data
from ROSAT PSPC pointings detected 110 objects, 68 of
which were identified with optical counterparts. The study
also included CCD photometry in V and I of the central
3.3 square degrees of cluster, with a magnitude limit of
V ! 18, and the photometry of the optical counterparts
was used to determine the likelihood of the counterparts
being cluster members. Additional photometry was pub­
lished by Prosser et al. (1996). 44 of the X­ray sources
were deemed to be possible cluster members. The X­ray
luminosity distribution function for the F,G and early K
stars showed the cluster to be more X­ray luminous than
the Pleiades. There was little difference for the late K
and M­type stars. Randich et al. argue that, unless their
sample is affected by incompleteness, their result is con­
sistent with the younger age of IC2602 compared with the
Pleiades, the younger earlier type stars having had less
time to spin down and thereby decrease X­ray luminosity.
The later types, having longer braking timescales, would
be rapid rotators in both clusters and hence have similar
X­ray luminosities.
3. Observations
Our photometry was obtained using the 1­m Swope tele­
scope at Las Campanas Observatory on the nights of 9--11
February and 17--18 May 1991. The detector was a Ford2
2048\Theta2048 pixel CCD, rebinned 2\Theta2. At the Cassegrain
focus each rebinned pixel corresponds to 0:87arcsec on the
sky. Observations were made through a Johnson V filter,
Gunn r and i filters and a narrow band (70 š A) Hff inter­
ference filter.
Two cluster fields were observed. Another field, offset
approximately 1 degree from the cluster, but at a similar
galactic latitude, was also observed in order that some
estimate of the contamination by galactic field stars might
be made. The unsuitability of this field for this purpose
is discussed in Section 6.2. The coordinates of each of the
field centres are given in Table 1 and their positions on
the sky are shown in Figure 1.
For each field, four exposures were made, one short and
three long to record respectively the bright and faint stars
at optimum exposure levels. Exposure times were 15:10:10
sec for the short and 600:600:600 sec for the long through
V:R:I filters respectively. Three exposures of 1800 seconds
each were made of the first cluster field using the Hff filter.
Photometric standard stars chosen from Landolt
(1992) to cover a range of spectral types were observed
during each night. Dome flat fields for the V,R and I fil­
ters, sky flat­fields for the Hff filter, and bias frames were
also taken.
Table 1. Coordinates of the field centres
Field ff ffi l b
(J2000.0)
IC 2602a 10:44:44.4 ­64:29:04.0 289.792 ­4.924
IC 2602b 10:41:01.5 ­64:11:27.2 289.328 ­4.802
``Offset'' 10:35:15.7 ­63:54:46.6 288.635 ­4.853

D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602 3
Fig. 1. The region of IC2602 showing the positions of the two
observed cluster fields
4. Data reduction
The basic CCD reduction was carried out using IRAF's
CCDRED package (Tody 1986). The raw images were
trimmed and overscan subtracted. No additional bias sub­
traction was found necessary. Multiple dome flat­fields
(sky flat­fields in the case of the Hff fields) were aver­
aged, and the data flat­fielded using the resultant images.
Further details of this process can be found in Massey
(1992).
Since the CCD pixel size (0.87arcsec) was large com­
pared to typical seeing conditions encountered at Las
Campanas we measured the full­width at half­maximum
(FWHM) of a number of moderately exposed stellar im­
ages on each frame and found the average value to be
close to 2 pixels. Thus we judged the stellar images to
be undersampled for determining and applying the point­
spread function (PSF) method of photometry needed for
the crowded cluster fields.
To overcome this we first used the aperture photome­
try routines within the DAOPHOT (Stetson 1987, Stetson
et al. 1990) package in IRAF for measuring the Landolt
standards. An aperture of radius 16 pixels or 14 arcsec
was chosen to match Landolt's original aperture as closely
as possible. These measurements were used to transform
measurements from the instrumental system to standard
magnitudes (Johnson­Cousins). The equations used were
of the form:
v = V + k v1 + k v2 X + k v3 (V \Gamma R)
r = R+ k r1 + k r2 X + k r3 (V \Gamma R)
v \Gamma i = k v\Gammai 1
+ k v\Gammai 2
X + k v\Gammai 3
(V \Gamma I)
where vri are the instrumental magnitudes, V RI are the
standard magnitudes, X is the airmass and the k are the
transformation coefficients. It should be noted that since
the reddest observed standard had V \Gamma I = 2:268, there is
some uncertainty associated with the colours determined
for stars redder than that. Isolated stars were then located
on the cluster frames and these were measured with this
``standard'' aperture and a variety of smaller apertures
to determine the dependence of the derived magnitude on
aperture size. This we refer to as an ``aperture correction''.
Because of crowding, particularly a problem for the
fainter stars, it was necessary to use PSF photometry on
the cluster frames. A best fit PSF was determined and sub­
tracted for each object frame independently. The residuals
from the PSF subtraction were found to be mainly in the
core of the profile and these were measured using aper­
ture photometry with a radius equal to the FWHM of the
PSF. This latter quantity was typically 10% of the PSF
flux itself, independent of magnitude. A PSF correction
was accordingly determined for each cluster star individ­
ually and applied to the PSF­derived magnitude followed
by an ``aperture correction'' as described above. The accu­
racy of this method was tested using isolated stars on the
object frames and was found to produce errors which were
less than those from other sources (see Table 2 below).
These instrumental magnitudes were then transformed to
the standard system using the transformation equations
described above.
The three sets of V,R,I photometry, resulting from the
three individual deep CCD exposures in each colour, were
then combined in order to eliminate spurious detections
resulting from the undersampling (eg. cosmic rays). Only
stars which were common to at least two of the three sets
were used. The photometry was then averaged. 77% of
the stars thus selected had a range in measured V magni­
tude less than the averaged internal photometric error for
that star. The resulting deep­frame photometry was then
merged with that determined using the short­exposure
frames. The positions of each star was determined using
the starlink package ASTROM and star positions from
the UKST/COSMOS database, and as such are accurate
to ¸ 0:2arcsec.
5. Results
The selection of candidate cluster members was made by
placing theoretical isochrones on colour­magnitude dia­
grams and selecting stars from their positions with re­
spect to the isochrones. The theoretical isochrones of
D'Antona & Mazzitelli (1994), and in particular those
using Alexander, Rodgers & Iglesias opacities with the

4 D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602
Canuto & Mazzetelli convection model, remain the most
comprehensive for the low­mass stars with which we are
concerned here. However, the calibrations used to trans­
form the isochrones from the theoretical quantities of T eff
and log L
L0 to the observed colours and magnitudes are not
well defined.
For 4000K ! T eff ? 3500K we used temperature
scales and bolometric corrections from Bessell (1995) and
for T ? 4000K, temperature scales and bolometric cor­
rections of Kurucz computed by Wood & Bessell (pri­
vate communication) which are available via anonymous
ftp from mso.anu.edu.au. For cooler stars Stauffer et al.
(1995) have made comparisons of a 70Myr isochrone, using
several different temperature scales and bolometric correc­
tions, with known Pleiades members. The best agreement
was achieved using temperature scales from Kirkpatrick et
al. (1993) with bolometric corrections from Bessell (1991).
Thus, for T eff Ÿ 3500K, we have used temperature scales
from Kirkpatrick et al. (1993) with bolometric corrections
from the more recent paper by Bessell (1995).
Given the uncertainty in the literature regarding the
age of the cluster, ranging from 8Myr to 36Myr, it was
decided to use 10 and 40 Myr isochrones as the limits for
selection of cluster members. The isochrones were trans­
formed to allow for a distance modulus of 6:0 and red­
dening of EV \GammaI = 0:044 (Randich et al. 1995) and broad­
ened to allow for the 0:2mag uncertainty in the distance
modulus. These limits were further broadened to allow for
the photometric errors. These errors for field IC2602a are
shown in Table 2.
We have also taken into account the effect of binarity
on the location of stars with respect to the theoretical
isochrones. The size of the effect depends of the composi­
tion of the binary. We assume that if the unresolved com­
panion has a lower mass and hence redder colour, it will
have the effect of shifting the position to a brighter and
redder position in the colour­magnitude diagram. For a
companion of equal mass, the increase in brightness cor­
responds to 0:75mag. However, Dabrowski and Beardsley
(1977) have shown that the increase in magnitude in the
case of some binaries is as large as 0:8mag, so we have
decreased the bright selection limit by 0:8mag to allow for
the presence of binaries. We note that the sequence of ex­
isting members from Prosser et al. (1996) show a width
of ¸ 1:5mag. Given the increase in photometric error in­
troduced though undersampling we feel that our broader
selection criteria are justified.
Primary candidate members of the cluster were those
stars which were located between the selection limits in
both (V,V--I) and (R,R--I) colour­magnitude diagrams.
Stars were selected as secondary candidate members if they
fell between the limits in one or other of the diagrams but
not both. Closer inspection of the secondary candidate
members reveals that many of them are very unlikely to
be true cluster members, having colours that place them
far from the selection limits in the non­selecting diagram.
Table 2. Photometric errors for field IC2602a.
Magnitude range Err(V) Err(V--R) Err(V--I)
12 ! V ! 13 0.02 0.04 0.03
13 ! V ! 14 0.03 0.06 0.04
14 ! V ! 15 0.04 0.08 0.05
15 ! V ! 16 0.05 0.10 0.06
16 ! V ! 17 0.07 0.15 0.10
17 ! V ! 18 0.08 0.17 0.12
18 ! V ! 19 0.13 0.14 0.13
19 ! V ! 20 0.13 0.14 0.13
A V vs V­I colour­magnitude diagram for field IC2602a
can be seen in Figure 2, for the second cluster field
(IC2602b) in Figure 3 and for the ``offset'' field in Fig­
ure 4.
Its is clear from the colour­magnitude diagrams that
the number of stars increases dramatically for V ? 18:5 in
field IC2602b, due to an increased scatter from the main
field population. In view of this increased level of con­
tamination we will consider only the primary candidate
members with V Ÿ 18:5. In field IC2602a we detect 45
primary candidate members. In field IC2602b we detect
33 primary candidate members. The primary candidate
cluster members are listed in Table 3.
The results of the Hff data are shown in Figure 5, as a
plot of R minus an instrumental Hff magnitude (with an
arbitrary zero­point) versus R for field IC2602a. The fig­
ure also shows the position of our photometrically selected
primary candidate members. Figure 6 shows a similar di­
agram for field IC2602b.
Photometry of all the stars observed can be obtained
electronically from the Armagh Observatory WWW server
(http://star.arm.ac.uk/¸dcf/ic2602.html) or by anony­
mous ftp upon request.
6. Discussion
6.1. Background Contamination
Given the low galactic latitude of IC2602 (b = \Gamma4:9) it
is to be expected that the stars selected using theoretical
isochrones will be contaminated to a large extent by back­
ground stars. In order to gauge the scale of this contami­
nation a similar selection process was applied to the pho­
tometry of the ``offset'' field. The selection process yielded
43 ``primary candidate members'' with V ! 18:5.
It should be noted that in using the offset field to gauge
the background contamination we are assuming that the
distribution of stars in it is representative of the back­
ground stars to be found in the cluster fields. This compar­
ison suggests that nearly all of the stars selected are back­
ground stars. However, see the discussion in Section 6.2
for further comments on the location of the offset field.

D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602 5
1 2 3 4
20
18
16
14
12
Fig. 2. A V vs V­I colour magnitude diagram for the first cluster field showing the stars selected as primary (filled circles) and
secondary (open circles) candidate cluster members. The solid lines are 10Myr and 40Myr isochrones, the dashed lines show the
selection limits including all sources of uncertainty discussed in the text, and the dotted line shows the bright selection limit
before any allowance was made for binarity.
6.2. Comparison with the results of Randich et al.
In their study of IC2602 using the ROSAT X­ray satel­
lite, Randich et al (1995) detected 110 X­ray sources in
an 11deg 2 area, 68 of which they identify with at least
one optical counterpart. 4 of these X­ray sources lie in our
field IC2602a: R64, R69, R76, R80. No detected X­ray
sources lie in our second cluster field. One X­ray source,
R25, lies in our offset field. Randich et al identify 7 stars
as possible optical counterparts to these X­ray sources (6
in field IC2602a, 1 in our offset field) based on photometri­
cally selected cluster members and additional bright stars
located close to the X­ray position.
Table 4 shows the results of our photometry for the
stars which Randich et al indicate as optical counterparts.
Photometry for the X­ray selected stars is taken from
Prosser et al. (1996) where available. The X­ray source
R80 is clearly identified with the bright star HD308016
whose magnitude of V = 10:66 places it beyond the bright
limit of our photometry.
Comparing the remaining 6 stars, there are differences
in both the V magnitude and the V \GammaI colours between the
two sets of photometry somewhat larger than the errors
in Table 2. Given that Randich et al's objects are active
stars, however, they are quite likely to be variable. Van
Leeuwen et al. (1987) have found variable stars in the
Pleiades with amplitudes of up to \DeltaV ¸ 0:2. Walter et al.
(1992) have recorded variability on a naked T­Tauri star
of amplitude \DeltaV ¸ 0:5, \Delta (R \Gamma I) ¸ 0:2. Thus a part of
the discrepancy at least might be reasonably attributed
to such variability. Since the majority of Randich et al's
photometry is unpublished we are unable to make a more
detailed comparison between the photometric datasets at
this point.
The positions of the detected X­ray sources in the re­
gion of IC2602 are shown in Figure 7. The figure also
shows the positions of our observed fields. It should be
noted that there are no X­ray sources in the the field
IC2602b and that there are relatively few in the region

6 D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602
1 2 3 4
20
18
16
14
12
Fig. 3. A V vs V­I colour magnitude diagram for the second cluster field with symbols as per Figure 2
surrounding it. This lack of X­ray sources is consistent
with the lack of primary candidate members found from
the photometry of that field. It should also be noted that
the offset field contains one X­ray source, R25. Both our
photometry and that of Randich et al indicated that this
star is a probable cluster member. If this star is a true
cluster member then the offset field may be positioned too
close to the cluster to give an accurate indication of the
background contamination. If this is the case then we are
over­estimating the level of background contamination as
the ``primary candidate members'' will contain real clus­
ter members as well as background stars which lie in the
appropriate areas of the colour­magnitude diagrams.
6.3. Comparison with the Pleiades
The Pleiades, by virtue of its proximity, is one of the
most extensively studied young (age¸70Myr) open clus­
ters. Hambly and Jameson (1991) have studied its mass­
distribution and luminosity function. We may use their re­
sults to independently examine the effects of background
contamination on our data.
By taking the star numbers for the inner 0:6 ffi radius
of the Pleiades as an approximate model for our region
of IC2602, and scaling to allow for the different distance
moduli and angular extents of the two clusters, we should
see roughly 22 cluster members in our cluster fields. Inso­
far as this comparison is valid, this suggests that the level
of contamination is less than that suggested by the num­
ber of ``primary candidate members'' located in the offset
field, assuming that the two clusters have similar star den­
sities and mass functions. We have made no allowances for
the differences in ``richness'' between the two clusters.
In summary, the evidence suggests that the level of
contamination due to background stars lies somewhere be­
tween 50%, as suggested by Pleiades mass functions, and
73% as suggested by comparing the ``offset'' field with
field IC2602a, although it is likely to be well below the
latter figure, given that the ``offset'' field appears to be
located within the cluster. Such levels of contamination
might well be expected given the location of the cluster,
and the broad selection limits used.

D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602 7
1 2 3 4
20
18
16
14
12
Fig. 4. A V vs V­I colour magnitude diagram for the offset field with symbols as per Figure 2
6.4. Luminosity Function
In Figure 8, which shows the distribution of primary can­
didate members with V magnitude, several features are
apparent. Firstly, for 12 ! V ! 14 there is a clear excess
of cluster members in field IC2602a in comparison with
cluster field IC2602b and the offset field. This excess is
not apparent in the range 14 ! V ! 16 where an effec­
tively similar number of primary candidate members was
selected in each cluster field. A larger number of stars were
selected in the offset field in this magnitude range. In each
colour­magnitude diagram there is an obvious field­giant
branch with 1:4 ! V \GammaI ! 1:9. This may be a cause of con­
tamination in the candidate list for V ! 14:5, but a reason
why this should be worse in the offset field in comparison
with the other two fields is unclear. There is a sharp falloff
in stars in all fields in the range 16 ! V ! 19. Because of
the likely high level of field contamination, further discus­
sion of the cluster luminosity function is not appropriate
at this time.
6.5. Hff Brightness
Given that the Hff filter had a passband of 70 š A, that the
equivalent width of Hff in an active late­M dwarf is ¸9š A
and that the equivalent width in absorption for an inac­
tive M dwarf is ¸1š A, we would expect a Hff magnitude
difference between active and inactive stars of ¸ 0:14mag,
assuming that cluster late­type stars are similar to so­
lar neighbourhood M dwarfs. This magnitude difference is
comparable to both the scatter in the plot and the errors
in R \Gamma Hff. Thus, we were unable to make any further
selection on the basis of the Hff magnitudes.
We were able to determine Hff magnitudes for 20 of the
45 photometrically selected stars in cluster field IC2602a
(the rest being too bright). 13 of the 20 stars lie to the Hff
bright side of the mean R \Gamma Hff level, 7 of these more than
3oe from the mean level (oe =¸ 4:2 š A). Of the stars that lie
below the mean, none are more than 1oe from the mean.
Similarly, for cluster field IC2602b we have determined Hff
magnitudes for 18 of the 33 primary candidate members.
The scatter in the data is much larger (oe =¸ 10 š A), and
so only two stars lie more than 1oe from the mean, both

8 D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602
13 14 15 16 17 18
­7
­6.5
­6
­5.5
­5
Fig. 5. Instrumental r vs instrumental r \Gamma Hff for cluster field IC2602a. The solid dots are the photometrically selected primary
candidate members for that field.
Table 4. A comparison of our photometry with that listed in
Randich et al. and Prosser et al. For details of stars with IDs
``f1'' or ``of '' refer to tables on the WWW site.
ID V V \Gamma R V \Gamma I Memb.
R25 16.38 --- 2.75 Y
of­6078 16.15 1.00 2.67 Y
R64 16.53 --- 2.51 Y?
F41 16.66 0.93 2.62 Y
R69A 18.0 --- 5.40 ?
f1­514 16.71 2.08 5.12 N
R69B 14.12 --- 2.36 N
f1­10 14.20 1.11 2.42 N
R69C 17.07 --- 2.87 Y
f1­738 17.16 1.04 2.94 Y?
R76 16.81 --- 2.97 Y?
F71 16.81 1.07 2.88 Y
in emission. The data for the primary candidate members
are shown in Table 5.
For field IC2602a the Hff data reinforces the likelihood
of a significant fraction of the photometrically selected
star being true cluster members; 35% of those stars with
Hff magnitudes being well in emission. For field IC2602b
the results are also consistent with the reduced number of
photometrically selected stars, and the lack of any X­ray
detections in that field.
7. Conclusions
We have presented VRI and narrow band Hff photom­
etry of two 15 \Theta 15arcmin fields near the centre of the
young open cluster IC 2602. The VRI photometry was
used to prepare colour­magnitude diagrams. Through the
use of theoretical isochrones, we have identified 78 primary
candidate cluster members. Using similar techniques on a
nearby offset field, and through comparison with data on
the Pleiades cluster from the literature, we have estimated
the level of background contamination to be – 50%. The
Hff data was also used to prepare colour­magnitude dia­

D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602 9
13 14 15 16 17 18
­8
­7.5
­7
­6.5
­6
Fig. 6. Instrumental r vs instrumental r \Gamma Hff for cluster field IC2602b. The solid dots are the photometrically selected primary
candidate members for that field.
grams. The location of the photometrically selected stars
on these diagrams reinforces the likelihood of their being
cluster members. Independent methods, such as a proper
motion study, or a study of radial velocities, are needed
to confirm the selected stars as cluster members.
Acknowledgements. We would like to thank Dr. M. Read
for providing the UKST/COSMOS positional data used to
determine star coordinates. Research at Armagh Observa­
tory is funded by a grant­in­aid from DENI. Data reduc­
tion was performed on the PPARC funded Northern Ireland
starlink node. DCF acknowledges support of a studentship
from Armagh Observatory. SLH was partially supported by
NSF Young Investigator award AST94­57455. WRJR acknowl­
edges financial assistance from the PPARC, grant number
GR/J25352.
References
Abt, H.A., Morgan, W.W. 1972, ApJ 174, L131
Bessell, M. 1991, AJ, 101, 662
Bessell, M.S. 1995, in ``The Bottom of the Main Sequence''
ESO Workshop, ed. C. Tinney, ESO, Garching, p. 123
Braes, L.L.E. 1962, Bull. Astron. Inst. Neth. 16, 297
D'Antona, F., Mazzitelli, I. 1994, ApJS 90, 467
Dabrowski J.P., Beardsley, W.R. 1977, PASP 89, 225
Durney, B., 1972, NASA SP­308, 282
Gould, A., Bahcall, J.N., Flynn, C. 1996, ApJ, 465, 759
Hambly, N.C., Jameson, R.F. 1991, MNRAS 249, 137
Hill, G., Perry, C.L. 1969, AJ 74, 1011
Kirkpatrick, D., Kelly, D., Rieke, G., Liebert, J., Allard, F.,
Wehrse, R. 1993, ApJ, 402, 643
Kraft, R.P., 1967, ApJ 150, 551
Lada, C.J., Margulis, M., Dearborn, D. 1984, ApJ, 285, 141
Landolt, A.U. 1992, AJ, 104, 340
Massey, P. 1992, ``A User's Guide to CCD Reductions with
IRAF'', NOAO Laboratory
Mermilliod, J.C. 1981, A&A 97, 235
Prosser, C.P., Randich, S., Stauffer, J.R. 1996, AJ, 112, 649
Randich, S., Schmitt, J.H.M.M., Prosser, C.F., Stauffer, J.R.
1995, A&A 300, 134
Reid, N. 1992, MNRAS, 257, 257
Rolleston, W.R.J. 1995, IAJ, 22, 181
Sagar, R., Myakutin, V.I., Piskunov, A.E., Dluzhnevskaya,
O.B. 1988, MNRAS, 234, 831
Skumanich, A., 1972, ApJ 171, 565

10 D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602
a
b
o
Fig. 7. The region of IC2602 showing X­ray sources corresponding to photometric cluster members (filled circles) and
non­members (open circles) from Randich et al (1995).
Stauffer, J.R., Hartmann, L., Soderblom, D.R., Burnham, N.,
1984, ApJ 280, 202
Stauffer, J.R., Hartmann, L., Burnham, N., Jones, B.F., 1985,
ApJ 289, 247
Stauffer, J.R., Hartman, L., Latham, D.W., 1987, ApJ 320,
L51
Stauffer, J., Hartmann, L.W., Barrado, D. 1995, ApJ, 454, 910
Stetson, P.B. 1987, PASP, 99, 191
Stetson, P.B., Davis, L.E., Crabtree, D.B. 1990, in ``CCDs in
Astronomy'', ASP Conference Series, 8, 289
Terlevich, E. 1987, MNRAS, 224, 193
Tody, D., 1986, iraf User Manual, NOAO Laboratory
van den Heuvel, E.J., Conti, P., 1971, Science 171, 895
Van Leeuwen, F., Alphenaar, P., Meys, J.J.M. 1987, A&AS,
67, 483
Walter, F.M., Vrba, F.J., Brown, A. 1991, in ``Cool Stars, Stel­
lar Systems, and the Sun'', ASP Conference Series, 26, 340
Whiteoak, J.B. 1961, MNRAS 123, 245
Wilson, O.C., 1963, ApJ 138, 832
This article was processed by the author using Springer­Verlag
L A T E X A&A style file L­AA version 3.

D.C. Foster et al.: CCD photometry of late­type stars in the young open cluster IC2602 11
Table 5. Hff colour and equivalent widths for the primary candidate members.
ID R R­Hff W Hff ID R R­Hff W Hff
F1 14.55 ­6.79 ­0.36 F36 14.78 ­6.13 2.92
F2 14.22 ­6.74 2.92 F38 14.47 ­6.24 ­4.11
F4 14.34 ­6.82 ­2.26 F41 15.73 ­5.83 25.25
F6 13.70 ­6.87 ­3.50 F42 17.01 ­5.92 19.30
F7 14.25 ­6.81 ­1.63 F46 16.12 ­5.88 20.96
F8 17.02 ­6.85 ­7.65 F49 14.46 ­6.23 ­3.50
F9 14.88 ­6.72 3.60 F53 14.79 ­6.16 0.93
F10 14.55 ­6.80 ­1.00 F54 16.92 ­5.91 20.13
F13 14.49 ­6.80 ­1.00 F57 16.89 ­5.98 14.50
F14 13.80 ­6.87 ­3.50 F59 14.17 ­6.18 ­0.36
F16 14.95 ­6.74 2.92 F60 15.09 ­6.10 4.28
F18 14.27 ­6.65 9.22 F61 16.88 ­5.79 28.82
F19 16.48 ­6.40 24.37 F62 15.00 ­6.18 ­0.36
F20 14.41 ­6.59 13.72 F64 14.53 ­6.19 ­1.00
F22 13.97 ­6.68 9.22 F66 14.45 ­6.22 ­2.88
F24 14.14 ­6.81 ­1.63 F67 14.09 ­6.15 2.25
F30 14.04 ­6.78 2.25 F69 14.67 ­6.16 0.93
F32 14.23 ­6.66 8.50 F70 14.06 ­6.12 4.28
F34 13.58 ­6.24 ­3.50 F71 15.74 ­5.84 24.37
12 14 16 18
0
5
10
15
Fig. 8. The distribution of primary candidate members in
IC2602a (solid line), IC2602b (dashed line) and the offset field
(dotted line) with magnitude.