With knowledge of the spectral types and IR colors, 
could be calculated for a subset of the stars.
Since B-V colors are not useful, one instead works from V-K as well as
R-I
(Carney 1983). For those stars with spectra, the intrinsic colors for a
given spectral type are taken from
Bessell (1979) and Bessell and Brett (1988). The Seaton (1979)
extinction curve is used for reddening. Extinction is solved for by
dereddening the stars, so that the dereddened colors match the observed
spectral type. By varying the
spectral class and spectral type (by one class and one subtype
respectively), I calculate the extinction
(Av) for each color. The best fit Av was then used to deredden the
colors. Extinction was calculated for all of the colors. However, Av
determined from the R-I color was used preferentially in the
dereddening procedure because it would be least affected by
chromospheric activity.
Dereddened colors could only be obtained for those stars
with either IR data or spectra.
The dereddened colors for those stars are
given in Table 13.
To determine luminosity, one can simply integrate the dereddened SED
and assume a blackbody law beyond the range of the photometric coverage.
Alternatively, if one assumes the stars have normal photospheres, then
the bolometric flux, in units of ergs/cm
/s, is given by:
In equation 2.3, 
is the dereddened V magnitude and BC is the bolometric correction.
The bolometric correction
is taken from Lang (1994). The bolometric correction is dependent on
spectral classification, the value of the correction tabulated in
Lang varies by up to 10% between luminosity classes III and V, with
class V sources having lower values than class III sources.
Walter et al.\
(1994) explicitly solved for the bolometric correction of
luminosity class IV objects. These value are systematically lower
than the Lang values for class V stars by 5 -- 10%. The estimation of
the bolometric correction based on tables from Lang for class V sources
introduces an error of about 5%.
The bolometric flux thus obtained was divided by the solar value (Allen 1973) at a distance of 380 pc to produce the luminosity given in the table.
Figure 5 shows the photometric
data from all Orion OB1 stars for which extinctions were calculated in this
manner overplotted with evolutionary tracks from D'Antona and
Mazzitelli (1994).
Many authors (e.g., D'Antona and Mazzitelli 1994, Swenson 1996,
Vandenberg 1985)
have recently published evolutionary
models for low--mass PMS stars. These models generally attempt to
reproduce the temperature and luminosity of a PMS star as a function
of time. The D'Antona and Mazzitelli tracks employ mixing length theory and
convection (following the prescription of
Canuto and Mazzitelli 1990). High temperature opacities are taken from Rogers
and Iglesias (1992) and low temperature
opacities are from Alexander et al.\
(1989).
In general, it is difficult to use such evolutionary models to obtain
exact ages for PMS sources because of difficulty in setting the zero
point. This difficulty mainly arises because it is unclear how stars form
in the natal stages. Most authors assume that these stages move very
quickly, so the exact starting point is unimportant. D'Antona and
Mazzitelli (1994) begin their calculations when the core temperature in
hydrostatic equilibrium reaches 
K.
Some young stars sit at the birthline (Stahler 1983) for an extended
period of time while still accreting material (e.g., the stars in Upper
Scorpius; Walter et al.\
1994).
It is not
appropriate to start the model evolving until after the star has started
down the Hayashi track, at this point the evolution of the star is
treated as a function of the mass of the star alone, and the disk is
discounted.
Other differences among the models, such as the treatment of convection and
opacities can generate age differences of up to 40%.
The models can be used to measure relative parameters
of the samples in the two sub--associations examined. The data given in
Table 13 indicate that, in general, the redding of
the stars near 
Orionis is small (A
). Based on this
observation, and the data in Figures 1 and
2,
the X--ray sources in Orion OB1a are
about 0.5 dex
older than the stars near 
Orionis. The stars near 
Orionis also
have a much narrower age dispersion. In both cases, the X--ray sources
are less than 1.5 M
. Both samples show some sources
which appear to be far younger than the rest of the sample. One
explanation for these stars may be that they are part of binary
systems. Binary systems of stars occupy
higher positions on the H--R diagram than single stars since the
magnitudes of the two components are combined (Simon et al.\
1993).
