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).