We constructed the colour-magnitude diagram using the photometry
published by Abraham de Epstein & Epstein (1985)
and Westerlund et al. (1988). First we corrected
for the mean reddening. The
colour excesses for the cluster are small; we used a mean excess
of E(B-V) = 0.02 mag (Westerlund et al. 1988) and assumed a
visual absorption (
) of 0.06 mag. Next we removed
superimposed field stars. Besides photometry the usual criteria
for deciding cluster membership are proper motion and radial
velocity measurements. Unfortunately at present these
measurements are only available for a few of the higher mass
stars (e.g. Perry et al. 1978; Edvardsson et al.
1995), although an extensive proper motion survey of the cluster
is currently being undertaken by Platais et al. (1996). We base
our final selection of cluster members almost entirely on
photometry and consider all the K dwarfs with no measurable Ca
II emission to be non-members.
Figure 2: Logarithm of the lithium equivalent width
(
[Li]) as a function of un-reddened
colour (
for Blanco 1 as compared to the Hyades,
Pleiades and 
Persei clusters
Our revised list of cluster members with their photometry and
positions is given in Table A1. This data was used to
construct the colour-magnitude diagram for Blanco 1 (Fig. 3 (click here)).
We have superimposed the ZAMS of Lang (1991), assuming a
distance modulus (
mag (Paper I). The
diagram exhibits no turn-off and little if any evolution
from the main-sequence. A gap is evident in the
main-sequence between 
0.08 mag and
0.27 mag. This been noted by other authors (e.g. Eggen 1970;
Abraham de Epstein & Epstein 1985) and is a common
feature of young
open clusters.
Figure 3: Blanco 1 colour-magnitude diagram
constructed from
our revised list of cluster members given in Table A1. The ZAMS
line is taken from
Lang (1991)
Our final list of non-members is presented in Table A2. Many
of these stars were rejected by previous authors whilst the others all
lie well above the ZAMS and are either field giants or nearby
dwarfs. Abraham de Epstein & Epstein (1985)
suggested the latter could possibly be
pre-main-sequence (PMS) stars still in a state of contraction to
the ZAMS. This is
unlikely given that even low-mass stars (
) reach
the vicinity of
the ZAMS only after a few million years (e.g. Mihalas &
Binney 1981).
Moreover this feature is not seen for the 
Persei cluster
(age 50 Myr), except perhaps for the very cool M-dwarfs (see
Prosser 1992).
In order to justify the existence of PMS stars, the cluster
would have to be very young indeed. In fact we have removed
virtually all the stars that deviated by more than 2 magnitudes
from the lower portion of the ZAMS (B - V > 0.6 mag) and
included them in our list of non-members as these are clearly
field stars. We recognise that the width of our main-sequence
still remains very large (given that the standard 0.75 mag
spread is due to binarity), and therefore our membership list
needs to be refined further to remove the remaining field star
contaminants.
The lack of an obvious turn-off is indicative of a very young
age and has lead previous authors to estimate ages 
50
Myr: e.g. 30 Myr (Westerlund 1963), 20 - 40 Myr
(Eggen 1970), 45
Myr (Edvardsson et al. 1995), 50 Myr (Epstein
1968; Perry et al. 1978; Abraham de
Epstein & Epstein 1985;
Westerlund et al. 1988). Typically the cluster is believed to
belong to the same age group as the Pleiades (Harris 1976).
Throughout this paper we adopt the absolute ages of Mermilliod
(1981) placing the age of the Pleiades at 
70 Myr and the
Persei cluster at 
Myr. Therefore on the
basis of previous estimates, before considering our new age
measuring technique, we can tentatively assume an age of 50 - 70
Myr for the Blanco 1 cluster.
A close examination of the upper main-sequence, however, reveals that all the latter age estimates hinge on the membership status of the B8 star HD 225187 (W71, Table A1), the bluest star on the colour-magnitude diagram. Perry et al. (1978) note that the proper motion of this star exceeds their criterion for cluster membership yet they include it in their final list of cluster members. The proper motion study of Platais et al. (1996) does not include this star as a member. Finally it is interesting to note that Perry et al. (1978) suspect HD 225187 to be a spectroscopic binary on the basis of its variable radial velocity. Recent studies (e.g. Kazarovets & Samus 1990) show the star to be slightly variable. The existence of a companion star cannot be ruled out.
Figure 4 (click here) compares the upper main-sequence
of the Pleiades and 
Persei clusters with Blanco 1,
where the turn-off from the main-sequence is clearly visible for
both the Pleiades (age 70 Myr) and 
Persei clusters (age
50 Myr). The dichotomy caused by HD 225187 (W71,
Table A1) is clearly apparent. If we include the star on the
colour-magnitude diagram then its position with respect to the
ZAMS imposes a very young age on Blanco 1 (i.e. 
50
Myr). If we remove this star then ages comparable and older
than the Pleiades become feasible. It is conceivable that the
lack of a turn-off is not necessarily a feature caused by a
young age but could also be due to a deficiency of high-mass
stars.
Figure 4: Colour-magnitude for the clusters
Blanco 1, 
Persei and the Pleiades. The dashed lines
show the isochronous
curves for 
Persei and the Pleiades (from Mermilliod
1981). The arrow
indicates the B8 star HD 225187
and cluster age
The lack of apparent evolution on the colour-magnitude diagram
of Blanco 1 clearly supports a young age
(Sect. 3.1 (click here)). If we consider the age of the
cluster to be 50 Myr, comparable with the age of
Persei, then we must expect the low-mass stars of the
cluster to have a similar rotational velocity distribution and
to exhibit a similar proportion of K dwarfs showing H
in emission.
There exists a definite correlation between H
emission
and rotation. Stauffer & Hartmann (1987) have shown that to
(
, ((
) the more rapidly
rotating K-dwarfs of the Pleiades cluster deviate from the
general trend, showing consistently greater
[H
]. This is consistent with the hypothesis
that activity increases with rotation. Greater activity implies
that either the mean chromosphere is hotter or that there is
more surface inhomogeneity, both giving rise to increased
H
emission (Paper I).
Figure 5 (click here) shows the Blanco 1 H
emission
equivalent width as a function of (B-V), compared with other
young clusters and the field. The rotational velocity
distributions for the Pleiades and 
Persei clusters are
well known (e.g. Stauffer 1991). Following Stauffer &
Hartmann (1987) we crudely split these stars into fast and
slow rotators to highlight the correlation between activity
and rotation. A cursory examination of the plot shows that the
Blanco 1 stars are distributed in the same manner as the
Pleiades. Of interest, however, is the distribution of the
young rapidly rotating 
Persei stars on this plot,
none showing the high levels of emission exhibited by the
Pleiades and Blanco 1. Assuming these stars generate their
magnetic flux by a shell dynamo process, then their high
rotation rates should generate more magnetic flux, more
chromospheric heating and hence stronger H
emission.
This appears not to be the case and may reflect a saturation
of magnetic surface activity (e.g. O'Dell et al.
1995) that acts to suppress the magnetic-heating
mechanism in the very rapid rotators. Also, this may be
evidence that magnetic heating processes are not entirely
responsible for the observed chromospheric emission seen in late-type stars.
The strongest H
emitting Pleiades star HII 1883,
(
[H
] = 5.5 Å) has a projected equatorial
velocity of 
= 140 km/s (Stauffer et al. 1984).
Similarly HII 347, with 
[H
] = 2.8 Å, has a
(Stauffer & Hartmann
1987). At equatorial velocities of 75 km/s, the line
broadening would be at least 1.6 Å, well above our
resolution. No Blanco 1 stars were observed to have such broad
H
profiles, so we expect that the rotation
distribution of Blanco 1 not to peak at such high velocities
characterised by young clusters such as 
Persei or the
Pleiades. Nevertheless, it is likely that ZS 142
(
[H
] = 3.14 Å) and ZS 37
(
[H
] = 2.13 Å), the two strongest
H
emitters are relatively rapid rotators amongst the
Blanco 1 population itself.
Figure 5: H
emission equivalent width as a
function of (B-V) for Blanco 1 compared with other young
clusters and the field.
The 
Persei and Pleiades stars are crudely split into
fast and slow rotators
to highlight the correlation between activity and rotation. U
indicates unknown
rotational velocity. Note that the distribution of
Blanco 1 stars is very similar to that of the Pleiades
Figure 5 (click here) reveals that the onset of H
emission in Blanco 1 is at (B-V)=0.93 (ZS 61), whereas the
older solar neighbourhood dwarfs (taken from
Panagi & Mathioudakis (1993)) show an onset of
emission at (B-V)=1.10. Similarly the Pleiades and
Persei clusters also exhibit an onset of H
emission earlier than the older field stars (Prosser et
al. 1991). Studies of these and other young clusters
have shown a correlation between 
[H
] and
spectral type and a trend toward H
emission at early
spectral types. Prosser et al. (1991) compared
the H
emission strengths of selected M-dwarfs for the
Pleiades, Hyades and field. They found that the mean strength
of H
, on average, declines as one passes from the
Pleiades through the Hyades to the field, and therefore may
act as a proxy for age. This can also be seen for the
K-dwarfs.
Figure 5 (click here) shows that the distribution of Blanco 1 K-dwarfs is similar to the Pleiades yet considerably different from the Hyades and field. This is indicative of the youth of Blanco 1 and the Pleiades. In a study of the chromospheric activity of low-mass members in the Hyades, Stauffer et al. (1991) noted that the calibration of this effect could yield a new tool for the determination of the relative ages of open clusters.
In order to quantify the decline of H
emission as an age indicator,
the ratio of H
emitters to the total number of observed
cluster members (
)
within the spectral range 
was
calculated for a selection of ``reference" clusters of differing ages.
The spectral range was chosen to cover only the K dwarfs, the region
where a strong link exists between the activity-age and rotation-age
relations.
Table 2: Cluster, the fraction of H
emitters amongst K
dwarfs (
) and the age, assuming published
ages for all clusters. The age of Blanco 1 has been interpolated
from Fig. 6
The ratio (
) is plotted as a function of age for the
``reference" clusters in Fig. 6 (click here). The points were well
fitted by an exponential
function therefore we show 
) versus
ln(Age). The straight line was fitted to the 
Persei, Pleiades and
Hyades points. The values (
were
attributed poissonian errors and
dashed lines added to approximate its upper and lower bounds.
No uncertainty was assumed for the ages.
The ages of the Pleiades, 
Persei and Hyades clusters
were all taken from Mermilliod (1981). The ratio
) for the
three ``reference" clusters and Blanco 1 are shown against their
respective ages
in Table 2.
Figure 6: Logarithmic plot of the ratio of H
emitters to the total number of
observed cluster members (
) as a function of age for
Persei, the Pleiades and Hyades. The plot allows us to
estimate the age of Blanco 1 (note the 
1 
uncertainty
shown by the horizontal error bar)
The ``reference" clusters have all been subject
to extensive spectroscopic surveys. The Pleiades and Hyades have been
popular targets for many years, in the main,
due to their close proximity to the solar neighbourhood. In particular,
comprehensive searches have been made for H
emission lines amongst
the late-type populations of these clusters (e.g. Prosser
et al. 1991). It is therefore unlikely that many more
H
emitters
remain to be discovered, particularly amongst the relatively bright K-dwarfs.
The (
) values are therefore to be
regarded as extremely
reliable and unlikely to change significantly.
Although the cluster Blanco 1 has not been studied to the extent
of the Pleiades, 
Persei and Hyades it is nevertheless
apparent that the cluster contains relatively few H
emitters. This spectroscopic survey has examined 
60%
of the Abraham de Epstein & Epstein (1985) sample
in the 
spectral range. The
estimated value of (
is therefore considered
realistic for the entire K-dwarf population of the cluster.
Plotting this value in Fig. 6 (click here) reveals that the
cluster is closer in age to the Pleiades than 
Persei
and is considerably older than previous estimates. An estimate
of 
90 Myr was extracted from the plot, with a 68%
confidence interval of 
25 Myr. We recognise that this
estimate is likely to change once cluster membership becomes
better defined. Preliminary membership information from the
proper motion study of Platais et al. (1996) of a number of K
dwarfs has allowed us to obtain a lower estimate of 60 Myr for
the age of the cluster.
and rotation
We consider here how H
emission and absorption act as
proxies for stellar rotation. Looking at
Fig. 5 (click here), the variation of H
emission
against (B-V), amongst the K dwarfs, for the four clusters and
the field. There is a large spread in equivalent width, much
more than for the H
absorption stars
(Fig. 1 (click here)). H
emission is more sensitive
to rotation and stellar effective temperature, for K dwarfs than
G dwarfs. The sensitivity is due to the contrast between the
chromospheric and photospheric temperatures, as at higher
effecive temperatures greater chromospheric electron densities
(and hence temperatures) are needed to drive H
into
emission (Cram & Mullan 1979; Panagi
1994). Once in
emission, H
emission has an approximate cubic dependency
on the chromospheric electron density, over the region of its
formation. Thus, although H
in emission indicates rapid
rotation, its high sensitivity to atmospheric conditions makes
intepretation of the rotation distribution very difficult.
This is not so for H
absorption, whose general trend as
a function of (B-V) is shown for both Blanco 1 and the solar
neighbourhood dwarfs relationship in Fig. 1 (click here). The
narrower relative spread, when compared to the H
emission at a given (B-V), is consistent with H
absorption being much less sensitive to variations in
chromospheric temperature and density.
We now consider the importance of the scatter in absorption
equivalent widths at a given (B-V). Two, perhaps
complementary, explanations can be used to interpret this
spread, assuming that the stars exhibit a spread in rotation
rates. Equating increased rotation rate with increased surface
activity, we would expect that the faster rotators will show
increased levels of surface inhomogeneity, implying that local
variations in temperature and density would dominate the
H
profile. This is the filling-in scenario, and so
increased rotation in this case equates to decreasing (in
magnitude) equivalent width. Alternatively (or in tandem),
increased rotation could yield an overall greater mean
chromospheric temperature (and density), so that the H
absorption equivalent width could actually increase in
magnitude. These conclusions follow naturally from the
theoretical work by Cram & Mullan (1979,
1985) on the formation
of H
in cool stars. They showed that, for a given
effective temperature, increasing the chromospheric temperature
gradient initially gives rise to a greater H
absorption,
before, at some critical chromospheric temperature (and hence
density), H
is driven into emission.
Figure 7: A 4-frame composite digitized sky survey map of the
Blanco 1 region centred near
ZS 134, with RA = 
and
DEC = 
. Approximately
80% of the cluster members of
Table A1 are shown
stars
The transition from H
absorption to emission predicted
by Cram & Mullan (1979) is an intriguing observational problem.
One would expect to see stars with apparently strong
chromospheres, evidenced by strong Ca II emission, but
little or no H
. They have been termed the zero-H
stars. A small number have indeed been observed in the solar
neighbourhood (Robinson et al. 1990;
Mathioudakis & Doyle 1992).
Mathioudakis & Doyle (1992) also looked at Mg
II emission
in a number of late-type solar neighbourhood dwarfs, and found
five stars with generally weak emission. These stars have
(
and have also been termed dM(e), dK(e) or
zero-H
stars. Thus two sub-classes of zero H
stars have been identified, weak and strong chromosphere stars.
In a Mg II emission with colour diagram, Mathioudakis &
Doyle (1992) observed that the hottest dK(e) or dM(e) stars have
fluxes that lie intermediate between those that show either
H
emission or absorption, and so argued that they are in
fact active stars with strong chromospheres. They similarly
argued that the cooler dM(e) stars are in fact quite close to
radiative equilibirum, with little or no signs of a
chromosphere. The results are based on the computatons by Cram
& Mullan (1979) who have shown that radiative equilibrium models
of late K to mid-M dwarfs produce decreasing H
absorption, ranging from 0.24 Å to 0.08 Å. The parallel
between Mg II emission and Ca II emission can be
drawn by noting that there exists a strong linear relationship
between the two, for stars in the solar neighbourhood (Panagi &
Mathioudakis 1993). Based on their relatively strong Ca II emission, we would therefore expect the Blanco 1 stars ZS 46
(
K5) and ZS 61 (
K2) (Table 1) to be active
stars, and their lack of a measurable H
absorption
feature implying that filling-in of the H
core by
active region emission dominates the line profile. Hence they
would prove to be extremely useful in the study of
chromospheric physics, as H
could flip between
absorption and emission under small changes in chromospheric
temperature and density.
An interesting suggestion has been put forward by Bopp (1985)
and Saar (1992) regarding the evolutionary status of the dK(e)
and dM(e) stars. They suggested that these stars are marginal BY
Dra stars, rapidly rotating binaries, showing both a high level
of surface inhomogeneity, manifest in stellar spots, and showing
strong chromospheric emission in H
. They also argue that
the dK(e) and dM(e) stars are in a short evolutionary phase
between the H
emission stars and the H
stars,
brought about as the stars are rotationally braked; the slower
rotating stars have weaker chromospheres and so show H
in absorption. There is no discussion of the time-scale of this
evolutionary phase. Based on this scenario, and given the
relatively young age of Blanco 1, we would expect to see this
transitional phase in K dwarfs, should it exist, whereas in the
much older solar neighbourhood dwarfs the effect would be seen
only for the late K and M dwarfs. The two stars ZS 46 and 61
also seem to satisfy this criteria.