3 Data analysis

3.1 Infrared spectrophotometry

Simultaneous infrared and optical observations of V 380 Ori were secured in March 1985, so as to have its IR energy distribution at the same time as the optical one and to check for possible variability in the infrared magnitudes. The observed magnitudes and the 1 $\sigma$ errors are reported in Table 2. Our photometric observations clearly show the presence of the large IR dust excess, which has been reported since the sixties by many authors. To check for a possible IR long term variability of V 380 Ori, we analysed the photometric broad band data available in literature. The data are summarized in Table 2 where, for very close observations we give the average values after checking that all the data fall well inside the observational errors. Though the comparison of observations taken with different instrumentation should be taken with care, it seems that long term oscillation (on a time scale of many months to years) might be present, while no systematic secular variations seem to be taking place.

  

Table 2: Infrared photometry of V 380 Ori^*

Notes to the table: ^* In parentheses the quoted errors.   Sources: ¹ Lorenzetti et al. (1983); ² Present work; ³ Hutchinson et al. (1994); three observations for JHKL; ⁴ Davies et al. (1990); ⁵ Mendoza (1968); ⁶ Kilkenny et al. (1985).

  

Figure 1: The CVF IR observations of V 380 Ori in March 1985. Ordinates are fluxes in 10-10 erg$\;$cm-2 s-1 $\mu^{-1}$

During the March 1985 campaign, we have also secured the infrared spectrum of V 380 Ori which is shown in Fig. 1. Our spectrophotometric infrared data were obtained with a resolution higher than that of the other previous observations, mainly to search for emission and/or absorption features. But, like in the previous lower resolution observations, the spectrum appears featureless as if it were produced by thermal emission of dust grains (see Cohen 1980, and references therein).

3.2 The optical spectrum

The overall optical spectrum of V 380 Ori appears substantially constant during the period covered by our observations, except for some cases discussed below. In particular the equivalent widths of the emission lines of both ESO and OHP spectra are comparable with each others and with those of November 1991 reported by Shevchenko (1997). We have also found that the continuum flux and slope did not change during 1983-1985. In Fig. 2 we present some examples of the spectra in different spectral regions. In each panel only spectrograms with the same resolution are compared. In the upper panel the April 1984 appears weaker than expected being of lower quality, as explained in Sect. 2.

  

Figure 2: Comparison of selected spectral regions of V 380 Ori observed at ESO during 1983-1985. Fluxes are in 10-13 erg$\;$cm-2 s-1 Å-1, not corrected for the interstellar extinction. The upper tracings are vertically shifted by steps of 2 10-13 erg$\;$cm-2 s-1 Å-1. The April 1984 spectrum is of poor photometric quality (see text)

Concerning the January 1995 OHP spectra, as discussed above, we could not perform the absolute calibration. Nevertheless, after the standard reduction procedure, the continuum slope of each OHP spectrogram well agrees with that of the ESO spectrograms. Taking into account the stability of our spectroscopic data and after verification that no major photometric variation occurred in the last decade, we corrected the continuum level of the OHP spectra by constant factors in order to have them overlapped onto the ESO spectra, with a mean continuum level of around 2 10^-13 erg cm^-2 s^-1Å^-1, being confident that this procedure does not produce significant systematic errors on the true values of the absolute fluxes of the lines. In fact, we intend to use in a forecoming paper the higher resolution OHP spectra for the analysis of the flux of metallic emission lines using the statistical method of the Self Absorption Curve (Friedjung & Muratorio 1987), in order to determine the physical conditions of the line emitting region.

Figure 3 shows the 1995 optical spectrum of V 380 Ori which appears very rich in emission lines, mainly due to metals in low ionization stages, and a few absorption lines of photospheric (such as the Stark-broadened Balmer lines and weak Fe I and Si II lines), or interstellar origin (e.g. the Na I yellow doublet, the DIB at 6275 Å). Most emission lines appear broad with a mean FWHM of about 150 km s^-1, which in some cases causes serious problems of line blending.

  

Figure 3: The optical spectrum of V 380 Ori observed in January 1995. The scales are the same as in Fig. 2. The fluxes are normalized as explained in Sect. 3.2. In the tracings dots indicate spikes. In Figs. 3a and 3b the vertical bars correspond to Mg II 4481 Å line. In Fig. 3c the PCygni profiles of Fe II (42) and the two photospheric absorptions of Si II (5) are indicated

 

Figure 3: continued. The optical spectrum of V 380 Ori observed in January 1995. In Fig. 3e the bars put in evidence the [O I] lines

Previous spectroscopic studies of V 380 Ori were mostly limited to the description of some particular lines, or have presented mean values of radial velocities and of line widths. Line lists have also been published (see Shevchenko 1994 and references therein), but in these identifications normally only the line considered to be the main contributor is reported. Since in V 380 Ori the lines are intrinsically broad, close blends cannot be easily separated, but secondary contributors may appear as bumps on the wings of the observed emissions in sufficiently high resolution spectra. In fact, most of the emission lines in our higher resolution OHP spectra present asymmetries that can be explained either by the contribution of the same line from different emitting regions with different radial velocities, or by another element.

For the purposes of our work on metallic emission lines we felt the need of making a line list as accurate as possible, and have therefore made a new line identification using a semiautomatic iterative procedure, based on Kurucz (1994) line list, which was run in the MIDAS environment. For each emission feature we measured the central wavelength and the flux above the nearby continuum by means of a Gaussian fit, and made a first identification of the possible contributors by an automatic selection from the list. Of the lines in the vicinity of the observed spectral feature we accepted only those belonging to transitions of ionized metals, and of a few neutral species, whose excitation potential and line strength are such that the lines are expected to be detectable. The few well known lines of H I, He I, Na I, Ca I, and Mg II  were excluded from this automatic search. We have measured the flux, FWHM and position of the expected contributors to emission blends with multiple Gaussian fits. It turned out that many asymmetric lines, which were attributed to a single line in previous works on V 380 Ori based on smaller resolution or lower S/N spectra, are in most cases well fitted by blends of two or more lines. No correction was made for the possible presence of photospheric absorption lines below the emission observed features, except for the broad H absorptions. The final selection was made in the usual way by checking the multiplet tables and by comparing with previously published identifications.

Table 3 gives the emission line spectrum of V 380 Ori in January 1995, which includes emission lines of H I, He I, [O I], Na I, Mg II, Si II, Ca I, V II, Ti II, Cr II, [Fe II], Fe II, and Ni II. In Table 3 we give for the line blends the wavelength and flux (not reddening corrected) of each contributor to the blend. We estimate an error of less than 10% for the strongest lines, 10% to 40% for the weakest lines. We remember that Table 3 is available at CDS (see footnote to the abstract).

  

Figure 4: Some spectral features in the January 1995 spectrum of V 380 Ori: a-c) Fe II emission lines fitted with two Gaussians. For a better evaluation of the fit the observed spectrum is also reproduced vertically shifted. d) Spectral region of the He I 5876 Å  line and the Na I resonance doublet. The scales are the same as Fig. 2

Most of the emission lines belong to Fe II permitted transitions with heliocentric radial velocity of about +37 km s^-1; the profile of the strongest lines is characterized by a hump on the blue side of the central symmetric emission peak (see Figs. 4a-c), which we were unable to otherwise identify. The line shape suggests the presence of two separate components, rather than a line broadening due for instance to a continuous velocity gradient. In the weaker lines the blue component is confused with the principal emission and produces a blue asymmetry. In order to take this into account, we have fitted each line to two components with the nearly same width ($FWHM~\simeq$ 150 km s^-1), a principal one centred at the stellar velocity, and a blue component. From the fit we found for the latter one a nearly constant blueshift of -140 km s^-1 with respect to the principal emission (Figs. 4a-c). The strong Fe II multiplet 42 lines have P Cygni absorption components centred at -250 km s^-1 with respect to the emission peak, and extending to about -300 km s^-1 (Fig. 3). This signature of outflow is an indication of line formation in a cool wind. Since no P Cygni absorption is seen in the Fe II lines other than multiplet 42, the wind should be rather teneous. On the other hand, the relative strength of the Fe II principal emission components within multiplets, and the preliminary results obtained from the analysis with the SAC method on the whole set of multiplets (Muratorio et al. 1998) clearly indicate that they originate in an optically thick medium, which, if were placed in front of the stellar photosphere, would produce a strong line absorption. Therefore, the absence of P Cygni absorption in the other Fe II lines, combined with the emission line broadness, is strongly suggestive of line formation in a geometrically bounded rotating region such as an equatorial disk, seen at a rather large inclination angle, rather than in a wind.

In addition to the permitted transitions, we have identified several weak narrow Fe II forbidden lines ($FWHM~\sim~70-100$ km s^-1, $v_{\rm hel}~\sim~35$ km s^-1), that might be expected to be formed in the outer less dense parts of the Fe⁺ emitting region. Their narrowness indicates line formation in a region different from that where the Fe II multiplet 42 P Cygni absorptions are formed. They are most probably formed in the less dense outer parts of the disk where the Keplerian velocity is much smaller.

[O I] is present on the OHP spectrum with a narrow 6300 Å line ($FWHM~\sim~84$ km s^-1, not corrected for the instrumental profile); the equivalent width is $W^{\rm e}_{\rm eq}$$~\sim~0.4$ Å; no asymmetry is detectable. Particular care was taken in verifying the goodness of the wavelength scale. We obtained for this line a heliocentric radial velocity of $+22~\pm~4$ km s^-1, the same as the interstellar Na I D lines (see below), which is indicative of a forming region at systemic velocity as also suggested by Böhm & Catala (1994) for other Ae/Be systems.

The 6363 Å component is clearly detectable on the blue wing of the Fe II 6369 Å emission line (Fig. 3). The 6300 line is also visible on the 1983 spectrum. Similar width and strength of the 6300 Å emission were found by Hamann (1994). The line was not detected in January 1992 by Böhm & Catala (1994), while Corcoran & Ray (1998) found the line to be much stronger ($W^{\rm e}_{\rm eq}$= 0.97 Å) in a spectrum taken only one month before our OHP observations. The intrinsic weakness of the [O I] lines, the variety of the instruments used by different groups and the lacking of systematic observations makes it difficult either to confirm a possible short time variability, or to give a realistic interpretation about the origin of this line.

The hydrogen Balmer series is present on our spectra from H$\alpha$  to H$\delta$.For H$\alpha$  a FWHM=260 km s^-1  was measured on the OHP and ESO spectra of 17 and 20 December 1983. Starting from H$\beta$  the emission emerges from the photospheric absorption. In spectra with sufficiently high resolution we measured the equivalent widths with respect to the base of the lines. In the low resolution spectra the equivalent width includes the absorption; in order to make a better comparison with the others we degraded the spectra of higher resolution and measured again the equivalent widths. For each run the differences between the values from original low resolution spectra and those from "degraded" spectra resulted lower than the uncertainties. The measured equivalent widths of the Balmer lines are presented in Table 4. In the OHP spectra all the emissions but H$\alpha$ appear asymmetric with a red peak and a bump on the blue wing. The heliocentrinc velocities from H$\alpha$ to H$\delta$ at the FWHM are +39, +42, +32 and +37 km s^-1 $\pm~5$ km s^-1, respectively, close to that of permitted Fe II. The peaks are redshifted by about 25 km s^-1. The observed differences can only partly be attributed to the different observing conditions, and are clearly indicative of a significant variablity, as also confirmed by the analysis of published data, which for H$\alpha$ range from an equivalent width of 58 Å in November 1981 (Finkenzeller & Mundt 1984), to about 80 Å in January 1992 (Böhm & Catala 1995). A value of 107 Å  was measured by Garrison & Anderson (1977) on plate spectra. Again nothing can be so far said about the time scale of the variations.

Helium is present in all our ESO and OHP spectra with the emission lines at 5876 Å and 6678 Å. No other He I emission lines such as 4471 and 4713 Å  have been found in the OHP spectra. The 4713 Å  line seems to be weakly present in the 1983 ESO spectrum. The 4921 Å and 5015 Å lines, if present, are masked by the very strong Fe II lines of multiplet 42.

The higher resolution OHP spectrum show the 5876 Å line as a wide asymmetric emission ($FWHM\,\sim$ 380 kms, Fig. 4d), the mean of the wavelengths at half maximum intensity gives a heliocentric radial velocity of about -22 km s^-1. A sharp absorption is clearly visible at +48 km s^-1 on the red wing of the line. The flux of the 5876 Å emission line corresponds to an equivalent width of 1.08 Å, a value which is comparable with that measured in our ESO observations of December 1983 ($W^{\rm e}_{\rm eq}$= 1.00 Å), and of April 1985 ($W^{\rm e}_{\rm eq}$= 1.26 Å), while the line appeared weaker ($W^{\rm e}_{\rm eq}$ $\sim$ 0.75 Å) in the spectrum of 15 March 1985. This He I line was also in emission in 1991 (Shevchenko 1994), while it was clearly in absorption in 1981 (Finkenzeller & Mundt 1984). Neither absorption nor emission was found by Böhm & Catala (1995) at 5876 Å on 16 January 1992. Our ESO observations show that also the He I 6678 Å line is in emission with $W^{\rm e}_{\rm eq}$ $\sim~0.5-0.7$ Å. Corcoran & Ray (1995) report that in February 1990 this line was present as a broad emission (FWHM = 340 km s^-1) at heliocentric velocity of +40 km s^-1, splitted into two nearly symmetric components by a central absorption, which again confirms the large variability in both strength and shape of the He I lines in V 380 Ori.

The Mg II $\lambda$4481 emission which is present in all our observations, in the higher resolution OHP spectra has nearly the same profile of He I 5876 Å with a broad emission split asymmetrically by a central absorption (Fig. 3) at heliocentric radial velocity close to that of the He I  central absorption.

The Na I resonance doublet is also present with two broad emissions which are both split by a narrow absorption at +24 km s^-1 which is interstellar in origin (Fig. 4d).

In the ESO spectrum the Ca II h and k violet doublet lines have very strong emission components (Fig. 2), while the near-infrared is characterized by the emission lines of the Paschen series, Mg II $\lambda 7896$, the O I  7771 Å and 8446 Å lines, and the very strong Ca II m.2 triplet. A comparison with previous observations (Herbig & Soderblom 1980, data obtained in 1970 and 1978; Shevchenko 1994, data obtained in 1991; Böhm & Catala (1995), data obtained in 1992) indicates quite constant equivalent widths of O I and Mg II lines and a slight variability of the Ca II triplet. The intensity ratio of the calcium triplet lines, corrected for the contribution of the Paschen lines, is close to unity in agreement with the other observations, and suggests line formation in an optically thick medium.

  

Table 4: Equivalent widths (Å) of the Balmer lines

Notes: The second value of H$\beta$ was obtained from low resolution and degraded spectra (see text).

3.3 The ultraviolet spectrum

The ultraviolet spectrum of V 380 Ori was observed with IUE at many epochs within different programmes: 1978 July 31 and October 28, 1979 April 17 and September 26, and 1989 October 19. The ultraviolet spectrum of V 380 Ori is characterized by the presence of many absorption features, and by a sharp flux decrease shortwards of about 1300 Å (see Fig. 5). In spite of the rich emission line spectrum observed in the optical, the Mg II 2800 Å doublet seems to be the only emission feature visible at low resolution, while Fe II is present with strong absorption features. The line identification of the UV spectrum of V 380 Ori includes absorption lines of O I, C I, C II, C IV, Si II, Si III, Si IV, Al II, Al III, and especially Fe II. C IV and Si IV are probably only of interstellar origin. The most prominent features are those of Fe II with a deep and broad absorption at 2750 Å and the flux cutoffs at 2410 and 2630 Å due to the blends of many resonance and low excitation lines. The low resolution does not allow us to determine whether the absorption features are also produced by the same expanding material which produces the P Cygni absorptions of the optical Fe II multiplet 42 lines, as discussed above. Hence the spectral features may not be good indicators for spectrum and luminosity classification. To overcome this problem, we have compared the UV spectrum of V 380 Ori with that of standard stars from the IUE Low-Dispersion Spectra Reference Atlas (Heck et al. 1984). We have found that the short wavelength cut off near 1300 Å is typical of B9-A0 stars (where it is also due to the strong red wing of Ly$\alpha$). The absorption features, and most notably those of Fe II, are more typical of a later spectral type, and are at any rate much deeper than in main sequence stars.

  

Figure 5: Comparison of the UV low resolution spectrum of V 380 Ori with that of spectroscopic standard stars. The offset spectra are, from top to bottom: the dereddened spectrum of V 380 Ori observed with IUE in September 1979, and the IUE spectra of A0V, A0III and A0I standard stars. The spectrum of V 380 Ori is dereddened assuming EB-V=0.5 and a 2175 Å  interstellar band 2.5 times weaker than for the mean Galactic extinction law. Ordinates are fluxes in 10-14 erg$\;$cm-2 s-1 Å-1

We are therefore brought to the conclusion that the ultraviolet absorption spectrum is formed in an envelope, cooler than the underlying stellar photosphere, whose partly overlapping lines hide the emission components.

3.4 Interstellar extinction and energy distribution

A depression near 2175 Å due to the interstellar band is clearly recognisable in all the UV spectra of V 380 Ori, partly affected on both wavelength sides by the shell absorptions (Fig. 5). Using the mean galactic extinction law, we have found that the feature disappears for E_B-V= $0.20~\pm~0.05$. This value is however in disagreement with the colour excess which can be derived from the optical energy distribution. In fact the literature generally gives a colour index of $B-V~\approx~+0.5$ for the star (e.g. Finkenzeller & Mundt 1984; Hillenbrand et al. 1992), which, if associated with an effective temperature of $\sim$10000 K, would imply a colour excess of E_B-V $\sim~0.5$. This is also in agreement with the strength of the DIB at 6280 Å  observed in the OHP spectrum. No multicolour photometry has been made simultaneously to the IUE and optical spectroscopic observations.

It should be recalled that the broad band photometry does not directly reproduce the continuum because of the contribution of the many emission lines. Using our spectrograms we have found that the emission lines contribute by nearly the same amount, $\sim$10%, to both B and V, so that the colour excess is not significantly altered by the emission lines. We think that the difference between the ultraviolet and visual colour excesses must be attributed to an anomalous interstellar extinction law in the V 380 Ori region. This is not an unexpected result since anomalous extinction curves have been in many cases found in regions of star formation.

We have analysed the depth of the 2175 Å band with a code kindly provided us by A. Cassatella, which adopts the parametrization proposed by Fitzpatrick & Massa (1986). We found satisfactory results assuming a colour excess of E_B-V = 0.5 with a 2175 Å  band 2.5 times weaker than the galactic one.