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Several low-resolution UV-spectra of BN Ori have been obtained with the International Ultraviolet Explorer (IUE) satellite in 1984 and 1986. Details of the observations are given in Table 5 (click here). Before each exposure the brightness of BN Ori was measured by the Fine Error Sensor (FES) of IUE in the Field Camera mode (Barylak 1989). In all cases the resulting tex2html_wrap_inline4296 magnitude was around 968, which is close to the results of the ground-based photometry in the periods around the IUE observations (Tables 3 (click here) & 12 (click here)). The low-resolution UV-spectra of BN Ori, HR 5999 and several other Herbig Ae stars show strong emission lines of OI, CII, SiIV and CIV (Tjin A Djie et al. 1982; Brown et al. 1986) shortward 1600Å. The high intrinsic radiation flux in these lines can probably only be supplied by accretion (Tjin A Djie & Blondel 1997) and it has been possible to derive information on the accretion parameters of A and F stars from the UV-continuum over the full wavelength range (1150- 3200Å) of IUE (Blondel & Tjin A Djie 1997). After the choice of a stellar spectral type and luminosity-class and addition of the observed visual magnitude, the distance and the foreground colour excess, we have matched the observed UV-continuum with standard star spectra from the IUE spectral atlasses for the photospheric and boundary-layer components and a optically-thick model for the accretion-disc (UV3C, Blondel & Tjin A Djie 1994, 1997).

Table 7:   Parameters used in the UV3C-model calculations with the low-resolution UV-spectra. For the B9I ab component we used tex2html_wrap_inline4300 = 10300K (Schmidt-Kaler 1982). The value given by {tex2html_wrap_inline4302 is the interstellar (foreground) E(B-V), while the tex2html_wrap_inline4306 value is the total excess

The results of the parameter adjustments required for the matching of the low-resolution UV-spectra of BN Ori, HR 5999, BF Ori and FU Ori are collected in Table 7 (click here). The masses of the stars (tex2html_wrap_inline4362) were estimated from the stellar radii (tex2html_wrap_inline4364 derived from the model calculation) and the photospheric tex2html_wrap_inline4366 (Schmidt-Kaler 1982) with the use of the evolutionary tracks of Palla & Stahler (1993). With this mass and radius, the width of the boundary layer (tex2html_wrap_inline4368) and its effective temperature, the corresponding mass-accretion rate (tex2html_wrap_inline4370) was derived. Since the mass of FU Ori cannot be derived from the evolutionary tracks, we have assumed that its precursor was a 1.0tex2html_wrap_inline4372 T Tauri star. In general the stellar spectral type and luminosity-class appear to be strongly constrained (to 1 or 2 subtypes) by the observed UV-continuum. For BN Ori we estimated a distance of tex2html_wrap_inline4374400pc from the data of Artyukhina (1959) and a foreground E(B-V) = 007 from the NaI D interstellar components.

Figure 8:   Low-resolution UV-spectrum of BN Ori after correction for the foreground extinction (grey line) with the accretion model calculation (black line). The used model parameters are listed in Table 7 (click here), and the "W" indicates the Walraven W-band flux (Table 2 (click here)) after correction for the foreground extinction. The low S/N in the observed spectrum (2000Å- 2400 Å) is due to the low sensitivity of the LWP camera in this range

Figure 9:   SED of BN Ori. The fluxes derived from the photometry of Sect. 2 (click here) are indicated by squares. The (*) points correspond to the fluxes calculated with the UV3C-model (Table 7 (click here)), obtained by matching the low-resolution UV-spectrum (Fig. 8 (click here)). The K- and L-band fluxes (open star) show the effect of removing the disc contribution in the model-calculation. Circles indicate the extiction-corrected model results

The calculated spectrum and the SED corresponding to the parameters of BN Ori in Table 7 (click here) are given in Figs. 8 (click here) and 9 (click here). These results show that the UV3C-model gives a good agreement with the observations from the far-UV up to the L-band in the NIR. In general the model (with an optically thick disc up to 40tex2html_wrap_inline4388) gives higher values than observed for the fluxes in the K- and L-band. The disagreement could not be reduced by taking the disc optically thin but when we limit the outer radius of the disc to 2.0- 2.5tex2html_wrap_inline4396 the NIRE almost disappears (Fig. 9 (click here)). This suggest that the disc has been largely dissipated during the outburst of 1947. For the parameters in Table 7 (click here) the E(B-V) is 012, which with an observed tex2html_wrap_inline4400 045 corresponds with tex2html_wrap_inline4402 = 033. This value is slightly higher then the 030 for spectral type F0III given by Schmidt-Kaler (1982) but not unacceptable as BN Ori may be not exactly F0III. Both the small circumstellar colour excess (005) and the negligeable NIRE indicate that BN Ori has almost no circumstellar dust left.

Table 8: Identification, EW (Å) and FW (tex2html_wrap_inline4404) of lines in the high-resolution UV-spectra of BN Ori, HR 5999 and FU Ori. See Table 10 (click here) for the used symbols

In addition to the continua of the UV-spectra of BN Ori, we have also information on the individual lines from the high-resolution LWP image obtained with IUE. In spite of its long exposure (13tex2html_wrap_inline4406) the S/N ratio is very low in the less sensitive part of the LWP camera (shortward 2400Å) but the region longward 2560Å is better exposed. Since the three low resolution UV-spectra of BN Ori show no differences (although they were taken at different times) we assume that the high-resolution spectrum of Nov.13, 1986 is representative for at least the period Jan'84- Nov'86. The FES-magnitude (tex2html_wrap_inline4410) during this observation agrees with a simultaneous ground-based photometric observation from DASA. From the comparison of the high-resolution UV-spectra of BN Ori, FU Ori and HR 5999 we can draw the following conclusions:

  1. The presence of an extended cool shell around the three stars is revealed by the many low-excitation absorption lines of FeII, CrII and MnII in their UV-spectra longward 2300Å. The lines are strongest in the spectra of HR 5999 and the correlation of their EW-variations with the variations in brightness of HR 5999 (Blondel et al. 1989) shows that these lines are formed in the shell of the star. Because of the lower brightness of BN Ori and FU Ori the S/N ratio in their spectra is lower than in the spectrum of HR 5999. In spite of this we can give upper limits for the EWs of the lines in the BN Ori spectrum and rough estimates (tex2html_wrap_inline441230%) of the EWs of the lines in the FU Ori spectrum. Table 8 (click here), in which the EWs and FWs of the three stars are listed, shows that the EWs and FWs of the BN Ori lines are smaller than those of the corresponding lines in the spectrum of FU Ori, and both are smaller than the EWs and FWs of HR 5999. Since the spectral types of the three stars are not too different we expect that the column densities in the shells will also decrease from HR 5999 through smaller values for FU Ori to almost negligible values for BN Ori.
  2. The MgI resonance line (2852.120Å) of BN Ori has two components: an interstellar one (with EW = 0.21Å) which is displaced by +26.5tex2html_wrap_inline4414 from its heliocentric position (determined from the NaI D\ interstellar lines) and a circumstellar one (with EW = 0.09Å) with a velocity of -53tex2html_wrap_inline4418 with respect to the interstellar one (Fig. 11 (click here)). Similar to the NaI D lines the continuum in the neighbourhood suggests that there exists a third, broad (FW = tex2html_wrap_inline4420400tex2html_wrap_inline4422) and shallow component, centered at the location of the two narrow components. Because of its width this component could be formed either at the stellar photosphere or near the inner accretion-disc. The MgI line of FU Ori also appears to be double. Both components have an EW of about 0.26Å and the circumstellar component has a velocity of tex2html_wrap_inline4424-67tex2html_wrap_inline4428 with respect to the interstellar component. The EW of the circumstellar component is larger than the corresponding one for BN Ori, while the foreground interstellar components of FU Ori and BN Ori are of comparable strength. The MgI components in the UV-spectra of HR 5999 are not resolved because of the large width of the circumstellar component. Its EW is about 10 times larger than the upper limit of the EW of the corresponding component for BN Ori and 3- 4 times the corresponding EW for FU Ori.
  3. The MgIIh (2802.7Å) and MgIIk (2795.5Å) lines in the high-resolution spectrum of BN Ori have broad absorption-troughs with weak emission-components inside (Fig. 10 (click here)). An accurate determination of the emission-components is difficult because of the presence of a reseau mark at 2795Å and an echelle order overlap between 2798Å and 2801Å. Nevertheless we have made an attempt to derive some information on the shell component by dividing the (normalised) observed fluxes in the MgII profiles of BN Ori by the corresponding ones in the profiles of 21 Vul, obtained from the IUE archives (LWR 9037). This star has a spectral type A7- F0IV, a tex2html_wrap_inline4434 of tex2html_wrap_inline4436200tex2html_wrap_inline4438 and a very weak shell component in the MgII profiles.

    Figure 10:   Observed MgII h&k profiles of BN Ori (LWP 9512) and the residual profiles after normalisation to the 21 Vul (LWR 9037) profiles. Top-axis give the velocity (tex2html_wrap_inline4440) relative to the wavelengths of the MgII h&k lines

    The result is given in Fig. 10 (click here), which shows that the shell component of MgII has a P-Cygni profile with outflow-velocities up to -250tex2html_wrap_inline4444. We have made a rough estimate of the MgIIk emission-component (EW tex2html_wrap_inline4446 15Å) by assuming that the intrinsic emission (without the blue absorption trough) is symmetric. With the continuum flux near 2790Å from the low-resolution UV-spectra of BN Ori (2.5 tex2html_wrap_inline4448tex2html_wrap_inline4450), the distance to radius ratio from the UV3C-model (4.35 tex2html_wrap_inline4452) and the total extinction correction (a flux factor of 2.08) we derive an extinction-free surface flux of 8.8 tex2html_wrap_inline4454tex2html_wrap_inline4456 for the emission-component of the MgIIk line of BN Ori. In the same way we derive from the corresponding profile of HR 5999 in the phase of maximum visual brightness (Blondel et al. 1989) an extinction-free surface flux of 6.45 tex2html_wrap_inline4458tex2html_wrap_inline4460 (tex2html_wrap_inline4462 = 68) and 1.2 tex2html_wrap_inline4464tex2html_wrap_inline4466 (tex2html_wrap_inline4468 = 76). The ratio of the MgIIk emission fluxes of HR 5999 and BN Ori is therefore close to 70- 140. This is higher than the ratio of the Htex2html_wrap_inline4472 fluxes of the two stars (Sect. 4 (click here)).

    Figure 11:   MgI 2852Å components of BN Ori

    Table 9:   Comparion of emission fluxes (tex2html_wrap_inline4474)

    The high-resolution MgII profiles of FU Ori have been published by Ewald et al. (1986). The profiles have a P-Cygni shape, but the noise in the neighbouring continuum does not permit to determine an outflow velocity. The extinction-corrected surface flux in the symmetrised MgIIk emission-profile is 8.6 tex2html_wrap_inline4502tex2html_wrap_inline4504 and the corresponding flux of the MgIIh emission is 7.1 tex2html_wrap_inline4506tex2html_wrap_inline4508. This total MgII emission flux of FU Ori is close to the value obtained by Ewald et al. With the use of the observed MgII profile of the F0I b star tex2html_wrap_inline4510Car (Praderie et al. 1980), which has a tex2html_wrap_inline4512 = 0 tex2html_wrap_inline4514, we corrected the profile of FU Ori for the photospheric contribution. The extinction-corrected surface-flux of the symmetrised MgIIk-profile then becomes 1.67 tex2html_wrap_inline4516tex2html_wrap_inline4518. Table 9 (click here) gives a survey of the emission fluxes of the 3 stars. Although the Htex2html_wrap_inline4520 and MgII profiles show variations on a short time-scale (tex2html_wrap_inline45221tex2html_wrap4538) the time variations in the integrated emission fluxes are small. This gives us a justification to use Table 9 (click here) in the following discussion, in spite of the fact that the Htex2html_wrap_inline4526 and MgII fluxes were not observed simultaneously.

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