EDP Sciences
Free Access
Issue
A&A
Volume 550, February 2013
Article Number A93
Number of page(s) 13
Section Stellar structure and evolution
DOI https://doi.org/10.1051/0004-6361/201117912
Published online 31 January 2013

© ESO, 2013

1. Introduction

The star system V383 Sco (HV 7021) was discovered to be an eclipsing binary at the beginning of the twentieth century during a photographic study of variable stars in a field of the Milky Way near the Galactic centre. As a partial result of those studies, Henrietta Swope (1936) presented photometric data containing observations of three eclipses whose minima occurred in 1901, 1914, and 1928. The very long orbital period, approximately 4900d (13.5 yr), is one of the longest known among eclipsing binaries. Its light curve has a wing-like shape at the beginning and at the end of the eclipses, which is characterized by slower photometric changes. Henrietta Swope considered it to be an effect of atmospheric absorption and suggested that V383 Sco is similar to ζ Aurigae-type stars. Later, the spectral type of the primary was estimated as F0Ia (Popper 1948). V383 Sco was then neglected for many years and O’Connell (1951) barely mentioned it during his analysis of the light curve of the V644 Cen system. Currently V383 Sco is a very poorly-studied system.

During the past decade the ASAS-3 survey has monitored V383 Sco in the standard V and I photometric bands. These observations cover the last eclipse, from 2007 to 2008 (see Figs. 1 and 2). While studying the V light curve, one of us pointed out the asymmetry of this eclipse (Kato 2008) and its similarity to the asymmetry of eclipses observed in the EE Cep system (see Gałan et al. 2012). Inspired by this discovery, we began the study of V383 Sco to verify the possibility that the eclipses could be caused by a dusty disk, as in the unique EE Cep and ε Aur systems. Here, we present an analysis based on the ASAS-3 survey V- and I-band photometry, low and high-resolution spectra, and all available visual, near- and far-infrared photometric data.

2. Observations and data reduction

thumbnail Fig. 1

Top: V- and I-band light curves from ASAS-3. Bottom: the V − I colour index. Crosses denote colour indices V − I obtained by interpolation in the cases when V and I measurements were not simultaneous, i.e. not taken during the same nights. Arrows and numbers show the moments of the pulsation maxima visible in V- and I-band light curves, used for timing analysis of pulsations (see Sect. 3.2). The dotted arrow shows a postdiction for the pulsation maximum at epoch Ep = 10 (corresponding to a maximum that occurred during the most recent eclipse – epoch E = 8) according to the ephemeris from Eq. (4).

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thumbnail Fig. 2

Expanded view of the 2007/8 eclipse of V383 Sco. The symbols are the same as in Fig. 1. The vertical dashed line marks the mid-eclipse moment2 (Table 2).

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The ASAS-3 survey (see Pojmański 2004) has monitored V383 Sco from the Las Campanas Observatory (Chile) with two standard filters since 6 Feb. 2001 (JD 2 541 947) in V and since 5 Sept. 2002 (JD 2 452 404) in I. The data were extracted from the database in April 2010 and processed using the pipeline described by Pojmański (1997). We selected data obtained with diaphragm numbers 2 and 0 in the cases of bands V and I, respectively, which had the smallest statistical errors. The resulting V and I-band light curves are presented in Fig. 1 with the colour index V − I which has been calculated for measurements made on the same nights (marked with circles) and by interpolation of close but non-simultaneous measurements (marked with crosses). These are available as Tables B.1–B.4 at the CDS.

High-resolution spectra (R = 84   000) of V383 Sco and the nearby long-period (P = 6545d) eclipsing binary V381 Sco in the spectral range 3800–6900 ÅÅ were obtained at La Silla with the High Accuracy Radial velocity Planet Searcher (HARPS) spectrograph using the EGGS mode on two consecutive nights in mid-October 2009. The spectra were extracted and the wavelength calibrated using the HARPS pipeline. Additionally, a low-resolution spectrum was acquired at the end of October 2009 with the Grating Spectrograph with a SITe (Scientific Imaging Technologies, INC.) CCD mounted at the 1.9 m Radcliffe telescope at the South African Astronomical Observatory (SAAO). Grating number 7 with 300 lines mm-1 and a slit width of was used. To calibrate the spectrum, the spectrophotometric standard stars LTT 2415, LTT 7987, and LTT 9239 were used. All the data reduction and calibrations were carried out with standard IRAF1 procedures. The extracted and flux-calibrated spectrum of V383 Sco covers the range  ~ 3800–7700 Å with a resolving power of R ≃ 1000. The journal of our spectroscopic observations is given in Table 1.

Table 1

Journal of spectroscopic observations.

3. Results and discussion

3.1. The eclipses and orbital period

The V- and I-band light curves and V − I colour index of the last eclipse of V383 Sco (in 2007/8) are shown in Figs. 1 and 2. There is a strong dependence of the eclipse depth on the photometric band. The mid-eclipse on 14 Jan. 2008 (JD 2 454 480) occurred about 170d earlier than predicted from the ephemeris (1)which was constructed with the period and the moment of minimum from Swope (1936). We have used the archival photometry of three minima observed before 1930 by Swope (1936, where the upper estimates and uncertain points were excluded from the analysis) and the ASAS-3 data obtained in the V band for timing analyses. Table 2 lists approximate times of mid-eclipses, estimated using a slightly simplified variant of the Kwee & van Woerden (1956) method2. The residuals (O–C, observations minus calculations) for the times of minima were calculated using the linear ephemeris in Eq. (1), are listed in Table 2, and are marked in Fig. 3. The best linear fit to the residuals at epochs 0, 2, and 8, gives the new ephemeris (2)It is not possible, however, to phase the light curve correctly (Fig. 4) with this new ephemeris. Regardless of which linear ephemeris is used, the eclipses are always shifted with respect to each other by up to several weeks. This makes the V383 Sco case reminiscent of the VV Cep system, where the 1997 eclipse occurred by about 1% of an orbital period later than predicted (Graczyk et al. 1999). The reason for such large differences in eclipse contact moments is not clear, but could be due to changes in the orbital parameters. However, in the case of V383 Sco it might be appropriate to consider another possibility. The changes in the eclipse contact moments could be caused by variations in the radius of the eclipsing component caused by pulsations which could explain the observed changes in the durations of the eclipses – the eclipse at epoch E = 2 clearly lasted significantly longer than the most recent one.

Table 2

Moments of minima and the O–C residuals estimated according to the linear ephemeris JDmidecl = 2   415   450 + 4900 × E.

thumbnail Fig. 3

O–C diagram with a linear (solid line) fit to the residuals. The moment of minimum at epoch 1 (vertical bar) is an estimated figure because adequate data were not available.

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thumbnail Fig. 4

Light curves of four eclipses phased with a constant period according to the new linear ephemeris (Eq. (2)). The open symbols denote archival photographic photometric data (squares: E = 0, circles: E = 1, diamonds: E = 2). The V-band ASAS data (filled circles: E = 8) have been shifted arbitrarily by  + 0.6 for facilitate comparison with the light curves from the different photometric systems.

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In the observational data collected so far, no traces of an eclipse of the secondary, cool component, have been detected. There is no way to predict the phase of the secondary eclipse, because through the entire orbital period the spectroscopic observations needed to infer the radial velocity curves are almost completely lacking.

3.2. Pulsations

thumbnail Fig. 5

Power spectra obtained using Fourier analysis on out-of-eclipse V-band (top) and I-band (bottom) data (see Sect. 3.2).

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The ASAS-3 photometric data outside of eclipses show apparent variations which seem to be connected with stellar pulsations. To investigate the amplitude and periodicity of these variations, we studied the out-of-eclipse data in the time intervals JD 2 451 947–JD 2 454 216 and JD 2 454 683–JD 2 455 145. After eliminating trends from the ASAS V, I and V − I data (by fitting and subtraction of a linear function and/or a second order polynomial), a fast Fourier transform was used to search for possible periods of variation.

Table 3

Frequencies and periods obtained with Fourier analysis of the out-of-eclipse V-band light curve.

thumbnail Fig. 6

ASAS-3 V-band (top), V − I (middle), and I-band (bottom) light curves outside of the eclipse, phased with period . The small amplitude of the observed V variation and its strongly non-sinusoidal shape with a wide minimum are related to the presence in the system of the second component which strongly dominates in this wavelength range (see Sect. 3.7).

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The power spectra obtained using the V-band data are presented in Fig. 5 (top) and the frequencies and the corresponding periods are given in Table 3. The dominant frequency is the lowest strong peak at  ~ 0.00505/d, i.e. a pulsation period of . A similar peak dominates in the I (Fig. 5, bottom) and V − I data, but less accurate by a factor of about two, while the peaks at greater frequencies (corresponding to shorter periods) are not visible. We therefore suggest that the period is a reliable value of a pulsation period for one of the components in this system. The V and I-band light curves and variations of the V − I colour index after phasing with this pulsation period are shown in Fig. 6. The amplitude of variations observed in the redder I band is much greater than in the bluer V band.

Pulsational maxima Ep from the V- and I-band light curves, as shown in Fig. 1, were used for a second timing analysis. The moments of maxima are shown in Table 4 with their corresponding O–C residuals, which were calculated using the following ephemeris (3)The pulsation zero epoch (Ep = 0) was adopted arbitrarily at JD 2 452 547.5. The best linear fit (Fig. 7) gives the ephemeris for pulsation maxima (4)The resulting pulsation period agrees very closely with the period calculated using Fourier analysis. The final ephemeris for pulsation maxima (Eq. (4)) can be used to postdict the moment of the local maximum that indeed occurred during the last eclipse, with an accuracy of a few days, corresponding to the pulsation epoch Ep = 10 (see the dotted arrow in Fig. 1). It is worth noting that the photographic observations of the eclipse at epoch E    =    1 show an analogous, local pulsational maximum (see Fig. 4).

thumbnail Fig. 7

O–C diagram for the pulsation maxima with a linear (solid line) fit to residuals.

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Table 4

Pulsational maxima and the O–C residuals estimated according to the initial ephemeris in Eq. (3).

3.3. Spectral Energy Distribution (SED)

Table 5

Photometry used to calculate the spectral energy distribution (SED).

There are not many existing photometric multicolour measurements of V383 Sco that can be used to construct its SED; however, these rare data exist over quite a wide spectral range. The TYCHO-2 catalogue has BT and VT magnitudes, although these are not very good quality. Inaccurate out-of-eclipse mpg data is available from Swope (1936). The ASAS-3 survey gives values of V and I magnitudes in the Johnson photometric system and 2-MASS JHK observations complement the energy distribution in the near-infrared. The revised versions of the Infrared Astronomical Satellite Faint Source Catalogue (IRAS FSC) and the Point Source Catalogue (IRAS PSC) include a potential counterpart for V383 Sco. These data together with new AKARI satellite detections significantly extend the observations to the mid- and far-infrared. All the available photometric out-of-eclipse flux estimations are given in Table 5. The photometry was dereddened using the colour excess value EB − V = 0.46 (see Sect. 3.6) and by adopting the mean interstellar extinction curve for R = 3.1 developed by Fitzpatrick (2004). The magnitudes from BTmpgVTVIJHK bands were transformed to the λFλ fluxes using the Bessell et al. (1998) calibration. The resulting fluxes are shown in Table 5 and the SED is shown in Fig. 8. The spectrum is dominated by an F-type supergiant in the visual and there is a significant excess from the near- to far-infrared.

thumbnail Fig. 8

Dereddened spectral energy distribution of V383 Sco obtained from the data given in Table 5 (points). The dashed line (blue in the electronic version of this paper) shows the F0I-type supergiant from the 1993 Kurucz Stellar Atmospheres Atlas (Kurucz 1993). The dot-dashed line (red in the electronic version) marks the SED of the SRb-type pulsating star RX Boo, built using the VOSpec virtual observatory tool. The dotted line (magenta in the electronic version) shows a black body with T = 500 K. The continuous line (orange in the electronic version) shows the sum of all 3 components appropriately rescaled (by a factor of 1.9 × 10-20 for the Kurucz model of an F0I star and by a factor of 1.3 × 10-4 for the observational SED of RX Boo) to fit the SED of V383 Sco. Triangles denote IRAS data with very low flux density quality (marked with the number 1 in the revised version of the IRAS Point Source Catalog), i.e. upper estimates: from the IRAS FSC catalogue (filled triangles) and from the IRAS PSC catalogue (open triangles).

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3.4. Low-resolution spectrum

A first inspection of the low-resolution spectrum (LRS) of V383 Sco revealed that it is dominated by a late A- or early F-type absorption spectrum. Only Hα appears as a weak emission. In the red part of the spectrum, above λ ~ 5500 Å, traces of molecular bands seem to be present. It was obvious that the V383 Sco spectrum is a superposition of a relatively hot A–F spectrum and a cool M spectrum. In an attempt to better define both components’ spectra, we used the FITSPEC task, from the STSDAS3.SYNPHOT4 package in IRAF, with three free variables – the renormalization of the two spectra and EB − V. The task searches for values of those variables that minimize the residuals between the templates and the observed spectrum. As templates we used the spectra from the Jacoby-Hunter-Christian spectrophotometric atlas (Jacoby et al. 1984) which have a resolution similar to that of our spectrum. The best model was obtained for a superposition of F0 I and M1 I spectra with a flux ratio FF/FM = 3.9 reddened by EB − V = 0.5. The observed spectrum of V383 Sco and the fitted model are compared in Fig. 9. The differences between the models for spectra in the ranges A9 I–F3 I and M1 I–M2 I (not every spectral subclass has a template in the atlas) are very small and the estimate of EB − V varies from 0.4 to 0.6. Based on the LRS, we can infer that the spectra of the components visible in the V383 Sco spectrum are F0 I and M1 I with an accuracy of the order of one or two subclasses and that the reddening in the direction of the star is EB − V = 0.5 ± 0.1.

thumbnail Fig. 9

Observed low-resolution spectrum of V383 Sco (thick line) compared with the model (thin line) that is a superposition of F0 and M1 supergiant spectra (grey lines). The observed spectrum and the model are dereddened by EB − V = 0.5.

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thumbnail Fig. 10

HARPS spectrum of V383 Sco in the region 4450–4550 Å compared with the spectra of HD 90772 and HD 75276 from the UVES library (Bagnulo et al. 2003) and ε Aur from the ELODIE archive (Moultaka et al. 2004).

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3.5. High-resolution spectra

A careful examination of the HARPS spectra of V383 Sco obtained on 2009 October 14 and 16 did not show any noticeable night-to-night changes. Thus, we average the high-resolution spectra, obtaining a high-resolution spectrum (hereafter, HRS) with a higher signal-to-noise ratio. Comparing the V383 Sco spectrum with those from the UVES library of high-resolution spectra (Bagnulo et al. 2003) we found that it falls somewhere between the A9 I and F2 I spectral classes (Fig. 10). The spectrum of V383 Sco is very close to one of the spectra of ε Aur from the ELODIE database (Moultaka et al. 2004), obtained outside the eclipse on 2003 November 1 (Fig. 10). Recently, Hoard et al. (2010) classified the spectrum of the F star in ε Aur as F0 II–III? (post-AGB), and pointed out that it has the appearance of an F0 supergiant. Based on the HRS, we suggest an F0 I spectral class for V383 Sco as well.

The Balmer series lines appear similar to those in the spectrum of a Be star. Hα is present as two emission components divided by an absorption (Fig. 11). The emission components are separated by about 70 km s-1 and the blue component is slightly stronger than the red one. The absorption seems to be a blend of two components separated by about 10 km s-1. In Hβ, in addition to the same emission and absorption components, the wings of a wide absorption profile are clearly visible (Fig. 11). The radial velocities measured on the Hα and Hβ emission wings are 89 km s-1 and 92 km s-1, respectively. They are practically identical with the velocity of the metallic absorptions (see Sect. 3.6). In the profiles of the higher series members, the wide absorption dominates. Traces of emission components in the core of the wide absorption are visible only in Hγ and Hδ.

The most common features in the spectrum of V383 Sco are sharp absorption lines of neutral and singly-ionised metals. Among the most conspicuous are the Ba ii absorptions. All the lines of multiplets 1 and 2 in our spectral region (4554 Å, 4934 Å, 5854 Å, 6142 Å, 6497 Å) are present and are among the strongest features. The Sr ii lines 4078 Å and 4216 Å are strong as well. In contrast, we were not able to identify the Li i line 6708 Å in our spectrum with S/N ~ 35 in this region. The strongest metallic absorptions show asymmetric or split cores which indicate that they could be blends of two components (see for example the Ba ii 6142 Å line in Fig. 12 and the stellar Na i absorptions in Fig. 13). Some lines show weak, inverse P Cyg profiles with a faint blue emission component and wide absorption wings (Fig. 12). Such variety of structures shows that the spectrum of V383 Sco is very peculiar and more observational material is necessary for its better understanding.

To estimate the projected rotational velocity of the F star we used the method described by Carlberg et al. (2011). We divided the spectrum of V383 Sco in the region 4500–5500 Å into eight intervals, avoiding Hβ and the interchip gap around 5320 Å. In each interval, the spectrum of V383 Sco was cross-correlated with a synthetic spectrum (Teff = 7500 K, log g = 1.5, microturbulent velocity 1 km s-1) from the POLLUX database (Palacios et al. 2010). To correct for the instrumental width, telluric lines were used. Because the macroturbulent velocity for luminosity class I, estimated from Fig. 17.10 in Gray (2005), appeared to be higher than the total broadening, we used the formula for luminosity class II developed by Hekker & Meléndez (2007). We found a projected rotational velocity v   sin   i = 7.6 ± 0.5 km s-1 for the F star in V383 Sco. The reliability of this estimation is qualitatively supported by Fig. 10. It is obvious in the figure that the lines in the spectrum of V383 Sco are sharper than those in the spectrum of HD 75276. De Medeiros et al. (2002) estimate v   sin   i = 9.8 km s-1 for HD 75276.

thumbnail Fig. 11

Hα and Hβ profiles in the HARPS spectrum of V383 Sco.

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In addition to the emissions in the strongest Balmer lines and the very weak emission components in the inverse P Cyg profiles in the HARPS spectrum of V383 Sco, we also identified the [O i] 6300 Å and 6364 Å forbidden emissions (Fig. 14). The 6300 Å line is blended with two telluric lines, but is clearly visible while the 6364 Å line is very weak. Both [O i] forbidden emissions are shifted longward with a velocity of  ~90 km s-1 which equals the radial velocity of the absorption lines in the spectrum of V383 Sco (see Sect. 3.6). A Gaussian fit to the [O i] 6300 Å emission gives a FWHM of the order of 31 km s-1.

thumbnail Fig. 12

Profiles of Ti ii 5186 Å, Y ii 5200 Å, and Ba ii 6142 Å in the spectrum of V383 Sco in which weak emissions on the shortward edges and wide absorption wings are clearly visible.

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thumbnail Fig. 13

Comparison of the region around Na i 5890 Å in the HARPS spectra of V383 Sco (stellar absorption at around  + 90 km s-1) and V381 Sco (stellar absorption at about  − 180 km s-1).

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thumbnail Fig. 14

[O i] 6300 Å and 6364 Å forbidden emissions in the spectrum of V383 Sco.

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3.6. Reddening and distance to V383 Sco and V381 Sco – two objects located near the Galactic centre

Like V383 Sco (), V381 Sco () is a poorly-studied system. There are wide discrepancies in the spectral classification of this star in the literature. Its discoverer, Henrietta Swope, assigned it a spectral type F0 and an unknown luminosity class (Swope 1936), while Popper (1948) classified it as an A5 Ia star. In Bowers & Kerr (1978), V381 Sco is listed as an M5 Ia red supergiant. Comparing the spectrum of V381 Sco with the spectra from the UVES library of high-resolution spectra (Bagnulo et al. 2003) we found the best match with an A8 II spectrum.

In Fig. 13 we compare the region around the Na i D2 5890 Å line in the spectra of the two stars. It is evident that despite the proximity of the stars in the sky, the Na i interstellar line absorptions are notably different. In both spectra there are some components with similar radial velocities in the range from  − 30 to 20 km s-1, which most probably originate in the same interstellar clouds. Several absorptions with different radial velocities visible in the spectra indicate that additionally, there are different interstellar clouds in the direction to each star. By fitting Gaussians to the particular absorption components, we measured the total equivalent width of the interstellar Na i D2 absorptions. Using the calibration of Munari & Zwitter (1997), we estimated a reddening EB − V = 0.43 for V383 Sco which agrees with the value obtained from the fitting of the LRS. For V381 Sco we obtained a reddening value EB − V = 0.39. According to Munari & Zwitter (1997), the accuracy of the reddening estimation based on the Na i equivalent width is higher for EB − V ≤ 0.4. For multi-component profiles of Na i, as in our case, the accuracy is generally . The resulting estimate of EB − V is only a lower limit, since the Na i lines can be saturated. Hereafter, we use an average EB − V = 0.46 ± 0.15 for V383 Sco.

Assuming an average out-of-eclipse brightness of V383 Sco , a reddening EB − V = 0.46 ± 0.15, and using for an F0 Ib supergiant from Schmidt-Kaler (1982), we derive a distance to the star of  kpc. If we use the absolute magnitudes for F0 Ia or Iab supergiants (Schmidt-Kaler 1982), then the estimated distance is unrealistic and exceeds 15 kpc. Using EB − V = 0.39 ± 0.15 (estimated above), an out-of-eclipse magnitude , and for an A8 II bright giant (Schmidt-Kaler 1982), we estimate the distance to V381 Sco to be  kpc.

When comparing the spectra of V383 Sco and V381 Sco we see an impressive difference in the radial velocities of the stellar absorption lines. This difference can be easily seen in the stellar Na i absorptions in Fig. 13. Using about 350 metallic absorption lines in the spectrum of V383 Sco, we measured an average heliocentric radial velocity Vh = 89.8 ± 0.1 km s-1. Measuring about 190 absorption lines in the spectrum of V381 Sco, we obtained an average heliocentric velocity Vh =  −178.8 ± 0.2 km s-1. The corresponding velocities relative to the local standard of rest (LSR) are Vlsr = 98 km s-1 and Vlsr =  −171 km s-1 for V383 Sco and V381 Sco, respectively.

The absolute velocities of both stars are high. They cannot be caused by orbital motion in such long-period systems. There is no evidence for their having a high proper motion in the literature. Most probably, these high velocities are connected with the kinematics of the regions close to the Galactic centre. V383 Sco and V381 Sco lie close to the Galactic centre and to the Galactic plane (z ~ 0.8 kpc and z ~ 0.4 kpc respectively). An examination of the Galactic-longitude–velocity diagrams of CO emission (Dame et al. 2001), H i (see Fig. 1 in Weiner & Sellwood 1999), Galactic planetary nebulae (Durand et al. 1998) and OH/IR and SiO-maser stars (Habing et al. 2006) shows that the observed velocities of V383 Sco and V381 Sco fall into regions close to the Galactic centre with very high velocities. The kinematics of the gas in these regions is inconsistent with purely circular motion (Dame et al. 2001; Weiner & Sellwood 1999, and references therein). Thus, we cannot use the standard Galactic rotation curve (e.g. Sofue et al. 2009) to estimate the kinematic distances of V383 Sco and V381 Sco.

Taking into consideration the distances to V383 Sco and V381 Sco estimated above, as well as their Galactic coordinates and the measured velocities relative to the LSR, both stars seem to belong to the bulge/bar structure in the inner part of the Milky Way (Weiner & Sellwood 1999; Habing et al. 2006; Vanhollebeke et al. 2009, and references therein). Moreover, the velocimetric model proposed by Vallée (2008) suggests that the spatial kinematic characteristics of these stars are inconsistent with their location in any of the inner spiral arms of the Galaxy.

One way to roughly estimate the distance to V383 Sco and V381 Sco is to use the kinematic model developed by Weiner & Sellwood (1999) for the inner part of our Galaxy. From the velocity contour plots in Fig. 8 in their paper we infer that V383 Sco is placed about 0.5 kpc behind the Galactic centre and that V381 Sco is placed about 1 kpc in front of it. For a distance to the Galactic centre R0 = 8.0 kpc (Sofue et al. 2009) the estimated distances to V383 Sco and V381 Sco are 8.5 kpc and 7 kpc, respectively. The main sources of errors in these estimates are the determination of the stars’ positions in Fig. 8 of Weiner & Sellwood (1999) by sight and the uncertainty in the R0 estimate (see Table 1 in Vanhollebeke et al. 2009). Thus, we estimate the combined random and systematic error of the kinematic distances as  ± 1 kpc. Thus, these distances are statistically consistent with the luminosity-based estimates indicated above.

To obtain the distance to V383 Sco by a third, independent method, we used the period–luminosity relation for pulsating Mira-type stars according to a formula given by Whitelock & Feast (2000): MK =  −3.47log P + β (β = 0.84 ± 0.14). Using and dereddened infrared brightness we obtained a distance modulus (K0 − MK) = 14.63 ± 0.20 and a distance of 8.4 ± 0.8 kpc. This value agrees closely with the previous two estimates ( kpc from the MV calibration and 8.5 ± 1 kpc from the kinematic model of the Galactic centre). Considering these three estimates to be independent Gaussian distributions and taking a weighted mean (conservatively using 1.9 kpc as the standard deviation for the MV estimate), we obtain 8.4 ± 0.6 kpc as the distance to V383 Sco. Similarly, the weighted mean distance to V381 Sco from the reddening (conservatively using 5.5 ± 1.3 kpc) and kinematic (7 ± 1 kpc) distances is 6.4 ± 0.8 kpc.

3.7. Discussion of the model of V383 Sco

Our interest in V383 Sco arose from the suspicion that it could be a system similar to EE Cep and ε Aur – unique long-period eclipsing binaries with a dusty debris disk as a component causing eclipses (Mikołajewski & Graczyk 1999). We considered the possible similarity between V383 Sco and two competing high- and low-mass models of ε Aur (see e.g. Guinan & De Warf 2002). The analysis of the SED by Hoard et al. (2010) and conclusions drawn from interferometric observations of the disk movement relative to the F-type star by Kloppenborg et al. (2010) seem to validate the low-mass model of ε Aur with a single B5V-type primary embedded in a dusty disk and a much less massive F-type post-AGB secondary. However, there still exist strong arguments in favour of the model with a high mass F-type supergiant and perhaps a binary system at the disk centre (see Chadima et al. 2011). We found that only the low-mass variant can be applicable to V383 Sco which explains the long-lasting eclipses (~11% of the orbital period) that are observed. Nevertheless, one significant observational constraint strongly disagrees with the low-mass ε Aur model. The depth of the V383 Sco eclipses shows a strong dependence on the mean wavelength of the photometric band, which is clearly visible in the V − I colour index (Fig. 2). The colour changes during the eclipses exceed 20, where the eclipse in the I band is more than twice as shallow as it is in the V band, which suggests that the eclipsing body has to contribute significantly to the total flux in the near-infrared. This case is not similar to the systems with a dark, eclipsing dusty debris disk like ε Aur and EE Cep. Instead, the eclipsing object has to be a cool supergiant. A good alternative that satisfies this condition is the resemblance to the BL Tel system. In BL Tel the eclipses caused by an M-type supergiant with orbital period show a similar asymmetric shape to those of V383 Sco and have a similar duration of about 10% of the orbital phase (see van Genderen 1986). Pulsations with a period of 651 observed in BL Tel correspond to a hot component, a UU Her-type variable (van Genderen 1977) – a subtype of SRd-type variables. In the case of V383 Sco, however, the amplitude of the 200-day variations is higher for longer wavelengths, which has not been observed in any type of pulsating stars and is unclear given the brightness changes caused by pulsations. Thus we conclude that the pulsating component of the V383 Sco system cannot be a hot star, but rather a cool component dominating in the IR.

Another, more likely possibility is that V383 Sco is similar to HD 172481 which, according to Reyniers & Van Winckel (2001), consists of an F-type post-AGB star and a cool M-type companion, probably AGB. Whitelock & Marang (2001) showed that the cool component in this system is a Mira-type variable with a pulsation period of 312d.

Several arguments can be considered in favour of the post-AGB nature of the F0 I component in V383 Sco: (i) the system’s location in the bulge/bar structure of our Galaxy; (ii) the high radial velocity; and (iii) the presence of dust that would explain the IR excess in the SED (Fig. 8). Additionally, the [O i] 6300 Å emission in the spectrum of V383 Sco may indicate the presence of a very low excitation nebula. If we assume that the wide and asymmetric profile of this forbidden line originates in an expanding nebula, then the observed outflow velocity will be of the order of 15 km s-1, which is typical for post-AGB expanding envelopes. Nevertheless, the infrared excess observed in the SED is a little too low to be a post-AGB star.

However, much stronger arguments favour the hypothesis that there is a cool, pulsating M-type component in the V383 Sco system. Two arguments follow from the spectral energy distributions, on the basis of spectra and/or photometric data: (i) in Fig. 9 in the red part of the LRS there are traces of molecular absorption band features characteristic of an M-type supergiant. The spectrum in the visual domain can be fitted with a combination of an F0I and an M1I spectrum. (ii) There is an infrared excess visible in the SED (see Fig. 8). At first glance, one might suspect that the system includes a very cool non-stellar component as the companion of the F-type star that cause the eclipses. If we apply the black-body approximation to the SED in the near IR part, it may seem that the observed excess could not be produced by a cool star. However, as a consequence of efficient mass loss accompanied by dust creation, pulsating Miras or semiregular SR-type stars have SEDs with maxima strongly shifted to the near infrared with respect to a black body of the same temperature (see Lobel et al. 1999). Thus the observed SED of V383 Sco can be reproduced by a superposition of two stellar spectra – a hot supergiant of F0I-type and a cool M-type giant/supergiant. In Fig. 8 we compared an SED of V383 Sco with the sum of the two spectra, individually rescaled to match the observed absolute flux levels, and a black body with a temperature of 500 K to explain the excess in the range up to 18 μm. The F0I-type spectrum is from the 1993 Kurucz Stellar Atmospheres Atlas (Kurucz 1993). The spectrum of the SRb-type star RX Boo (Fig. 8) was built using the VOSpec virtual observatory tool5. According to the GCVS database RX Boo is an SRb-type star that varies with an amplitude of  ~25 in the V-band and changes spectral type at least in the range M6.5e–M8IIIe with a typical period of . The presence of a cool star such as RX Boo is sufficient to reproduce the observed SED of the V383 Sco system.

There are two more arguments for the presence of an M-type supergiant in the V383 Sco system, which provide additional evidence for pulsations of this component, come from the photometric behaviour. (iii) The eclipse depth strongly depends on the colour, i.e. on the mean wavelength of the photometric band, which shows that a cool and very bright object dominates in the near IR. (iv) The observed amplitude of pulsations is much smaller at shorter wavelengths. Using the out-of-eclipse light curves, we estimate amplitudes of the pulsation related variations as 041 and 021 for the I and V bands, respectively (see Table 7). Normally the amplitude of pulsations decreases towards longer wavelengths, but in this case it is the opposite. This can be explained by the amplitude of pulsations in short wavelengths being substantially “suppressed” by the presence of a much brighter (in the visible range) F-type supergiant. We are not sure how deep the eclipses really are because the minima were not covered during the eclipses. However, by extrapolating the changes near the mid-eclipse during the most recent epoch we estimate that the apparent magnitudes during the minimum should not be greater than about 111 and 143 for the I and V bands, respectively. If we estimate the mean magnitudes in the minima of pulsations as 983 and 1086 for I and V, and assume that eclipses are total, we can estimate that the flux ratios of both components FM/FF change during the pulsation cycle in the range 0.444–1.107 in I and 0.044–0.267 in V. The upper value 0.267 agrees closely with the upper value from the LRS in the visible domain obtained around the maximal phase of pulsations: FF/FM = 3.9 ⇒ FM/FF = 0.256 (see Sect. 3.4). Hence, the amplitudes of pulsations of the cool M component could reach about and in the I and V bands, respectively, stronger in the shorter wavelengths as is expected for a pulsating star. In this case the cool component could be a semiregular, SR-type pulsating star. The estimated amplitude of pulsations strongly depends on the depths of eclipses. With increasing eclipse depth, the ratio FM/FF will decrease and this is reflected by a higher amplitude of pulsations of the M-type component. Indeed, the observed SED suggests that the ratio of the components fluxes FM/FF in the I band could be a few times (roughly an order of magnitude) lower than estimated from the photometric data. The true amplitude of pulsations could in fact be much higher and the pulsating component could well turn out to be a Mira-type star. Very high amplitude pulsations could provide an explanation for the changes of eclipse contact moments observed during different cycles, which can result from the variations of the pulsating supergiant radius. It could also explain the observed local maxima, visible during eclipses (e.g. at E = 1 and E = 8) as a consequence of an increase in the cold component’s flux (caused by pulsations) together with a decrease of the hot component’s flux (caused by obscuration).

Table 6

Luminosities of components of the SED of V383 Sco.

Table 7

Estimated minimum values of the eclipse depths () and the pulsation amplitudes δ.

thumbnail Fig. 15

Synthetic V- and I-band light curve and V − I colour index curves obtained by modelling, plotted with solid lines (green, red, and orange, respectively in the electronic version of this paper), in comparison with observational photometric data. Between the flux and colour curves, residuals are shown. The model assumes a temperature of the cool star Tcool = 2800 K. The expanded view of the out-of-eclipse V and I light curves (box with a dashed line) is shown in Fig. 16.

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We have prepared a model of the 2007/8 eclipse of V383 Sco taking into account the pulsations of the cool component using our own simple computer code. Our model is extremely simplified – the stars are considered as ideal spheres, limb darkening has been neglected, and stellar fluxes are approximated as black bodies. The changes in the radius of the M-type supergiant and in its brightness have been expressed with cosine functions. A more detailed description of the model assumptions and the parameters that have been derived are presented in Appendix A. As inputs to the model we assumed that the hot component has an effective temperature Tcool = 7500 K and its radius Rhot = 58   R was estimated using the Stefan-Boltzman law by adopting the luminosity Lhot extracted from the SED (see Table 6). The best solution was found for the M supergiant at effective temperature Tcool = 2800 K and with radius , which is changing due to the pulsations by . In our model the supergiant totally obscures the hot component during an eclipse minimum. The resulting amplitudes of the brightness changes in the cold supergiant – if it were observed as a solitary star – would be 366 and 175 in the V and I bands, respectively. The synthetic V- and I-band light curves and V − I colour index curve are shown in Fig. 15 and compared with photometric observational data. The eclipse is very deep, about 56 in V and 2.5 in I and needs to be confirmed by future observations, because the system has never been observed at the exact moment of minimum. It will be necessary to organize an observational campaign for the next eclipse in order to obtain data with a dense sufficiently time coverage. With the parameters of the system components derived from our model, we estimate that if a secondary eclipse occurs, it would be very shallow, i.e. a few 0001 in the V band and not deeper than about 002 in I. Given the strong changes in brightness caused by pulsations that are quite irregular in terms of amplitude it is rather certain that the secondary eclipse will be undetectable in the visible range. The situation is somewhat better in the near-infrared (JHK bands), where we estimate that the eclipse depths should fall in the range 004–006, though this would still be very difficult to observe. Careful inspection of Fig. 16 reveals an additional interesting feature in the synthetic light curves. To our surprise and satisfaction, our model has reproduced flattened pulsational minima in the synthetic V-band light curve, which we had noticed in the phased observational V-band light curve (see Fig. 6, top) without understanding the origin of this phenomenon. Now we know that this is the natural consequence of the fact that the pulsating star is a component of a system with an object which strongly dominates in a given photometric band. The changes which look roughly sinusoidal for a single pulsating star will have flattened minima when we put the pulsating star in a binary system together with a much brighter star6. This result can constitute the additional, fifth (v) argument to support the proposed model of V383 Sco with a pulsating cool supergiant as the eclipsing component. The best results are obtained with a low temperature of the cool component, Tcool = 2800 K, which agrees very closely with the result for the M star based on the SED. The unreddened V − I colour index close to the mid-eclipse is approximately 40, which is close to the index expected for a M7-8 star. We found earlier (Sect. 3.4, Fig. 9) that the M star was better described by an M1-type fit to the LRS. The difference between M1 and M7-8 issignificant and corresponds to a temperature difference of about 700–800 K. On the other hand, the LRS was obtained at a pulsational maximum (on 31 Oct. 2009 – JD 2 455 136), according to the ephemeris from (Eq. (4)) at a pulsation phase almost exactly equal to zero (~0.03). At that time, the star had reached the temperature close to the maximum possible value of the pulsation cycle. Thus, the results are consistent with our model, because during its pulsations, the cool star changes its spectral type, at least in the range from about M1 to about M7–8.

thumbnail Fig. 16

Expanded view of the synthetic fit to the out-of-eclipse V and I light curves and residuals.

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4. Conclusions

Using the ASAS-3 photometric data together with high (HARPS) and low (1.9 m telescope at SAAO) resolution spectra, we have revised our knowledge of the long-period eclipsing binary V383 Sco, for the first time since the 1930s. Contrary to our initial expectations, the system does not resemble the unique systems with dusty debris disks as eclipsing bodies, neither the famous, very long period eclipsing binary ε Aur, nor EE Cep. Instead, we have found a number of arguments in favour of a new model in which V383 Sco could be similar to HD 172481 – a system with a post-AGB F0-type star and a cool, pulsating M-type supergiant. The most important difference between these systems lies in the orientation of the orbits which enables the observation of eclipses only in the case of V383 Sco. While the hot component in V383 Sco shows some features characteristic of post-AGB stars, the detected infrared excess is too small and we can say with certainty that it is a supergiant of spectral type approximately F0. However, in the case of the second component that causes the eclipses, we found strong evidence that it is a cool, pulsating M-type supergiant on the basis of the following arguments:

  • i)

    Traces of molecular absorption bands, characteristic of anM-type supergiant, are present in the red part of the LRS.

  • ii)

    The observed SED can be reproduced by the superposition of two stellar spectra, a hot supergiant of F0I-type, and a cool M-type giant/supergiant.

  • iii)

    The eclipse depth strongly depends on the photometric band – it decreases in the direction of increasing wavelength.

  • iv)

    The pulsation amplitude is lower in the shorter wavelength band.

  • v)

    The flattened “bottoms” of the pulsation minima observed in the V band are consistent with a synthetic model of the eclipses.

    In addition, the presence of the pulsating supergiant in the system could explain the changes in duration and shape of the eclipses.

The forbidden emission [O i] 6300 Å in the HRS indicates the presence of a very low excitation nebula around the V383 Sco system.

V383 Sco and a similar eclipsing binary V381 Sco (by chance located close to each other in the sky), have very high and oppositely directed radial velocities: 89.8 km s-1 vs.  − 178.8 km s-1. These are in agreement with the kinematics of the regions close to the Galactic centre. V383 Sco and V381 Sco lie close to the Galactic centre and to the Galactic plane (z ~ 0.8 kpc and z ~ 0.4 kpc, respectively). The distance to V383 Sco estimated using reddening, radial velocities, and the period-luminosity relation is 8.4  ±  0.6 kpc. The distance to V381 Sco estimated from the reddening and kinematic distances is 6.4  ±  0.8 kpc.

The observational material collected so far is just sufficient to construct a rough model of the V383 Sco system, which reproduces the main features of observed photometric changes in the V and I-band light curves. With our model we obtain very deep eclipses: about 56 in V and 25 in I. This result requires confirmation by future observations, as the exact moment of minimum has never been observed.

The next eclipse should start in the middle of 2020. It is important to carry out an extensive photometric and spectroscopic campaign to observe it. In the meantime, systematic spectroscopic monitoring would be useful to obtain a spectroscopic orbit, with which it would became possible to obtain a more precise estimate of the basic parameters of the system’s components such as their masses and radii. Simultaneous photometric observations would be still valuable for the study of changes in the M-type supergiant radius as a function of the pulsational phase. Observational material obtained for this object over a decade time scale is still quite scant and any additional data should lead to very valuable and interesting results.


1

IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA) under co-operative agreement with the National Science Foundation.

2

The times of mid-eclipse were estimated as follows: for the eclipse at E = 8, T8d is the moment on the descending branch at which the brightness has dropped below the out-of-eclipse mean, and for eclipses at E = 0,2 and 8, TEa is the moment on the ascending branch at which the brightness reaches below the out-of-eclipse mean. The mid-eclipse was defined as TEm = TEa − (T8a − T8d)/2.

3

STSDAS is a product of the Space Telescope Science Institute, which is operated by AURA for NASA.

6

Let us assume that a variable pulsating star shows roughly sinusoidal variations in apparent magnitudes VC(t) ~ sinf(t) =  −2.5log (FC(t)). When we insert this star into the binary system with a non-variable companion, the system as a whole will change its brightness according to the expression VS(t) ~  −2.5log (10 − 0.4sinf(t) + FH), where FH denotes the contribution of the companion to the total flux, and the total flux of the system at any time is FS(t) = FC(t) + FH. When the contribution of the companion to the total flux of the system is small (FH ≪ FC(t) = 10 − 0.4sinf(t)), the system shows almost sinusoidal variations. In contrast, when (FH ≫ FC(t)), the V light curve consists of broad minima alternating with relatively short maxima.

Acknowledgments

This study was supported by MNiSW grants N203 018 32/2338 and N203 395534 and financial assistance was given to DG by the GEMINI-CONICYT Fund, allocated to project 32080008. W.G. and B.P. acknowledge financial support for this work from the BASAL Centro de Astrofisica y Tecnologias Afines (CATA) PFB-06/2007. W.G. also gratefully acknowledges support from the Chilean Center for Astrophysics FONDAP 15010003. We greatly acknowledge the variable star observations from the AAVSO International Database contributed by observers worldwide, and used in this research. This research was conducted in part using the POLLUX database (http://pollux.graal.univ-montp2.fr) operated at LUPM (Université Montpellier II – CNRS, France) with the support of the PNPS and INSU. MG was financed by the GEMINI-CONICYT Fund, allocated to project 32110014. The authors thank the referee, Petr Harmanec, for his very constructive comments.

References

Online material

Appendix A: Simple model of V383 Sco

To make the model of the 2007/8 eclipse of V383 Sco, which we describe briefly below, many simplifying assumptions were made. Stars were considered as spheres with radii Rhot, and effective temperatures of photosferes Thot, Tcool for the hot F-type and cool M-type components, respectively. Limb darkening was neglected and stellar fluxes were approximated as black bodies. We describe the main eclipse (of the hot component by the cool supergiant) by taking into account an impact parameter D which measures the projected distance between the centres of stellar disks at the mid-eclipse point, as well as changes in the radius of the cool star in the range  ± ΔR from the mean value of the radius as a result of the pulsations (Fig. A.1). Pulsational changes in the radius and brightness of the cool component were described simply, similar to the approach used for modelling pulsations in Mira χ Cyg by Reid & Goldston (2002). In our model we assume that the changes in magnitudes and radius of the cool star can be expressed with cosine functions where ΔM is the semi-amplitude of brightness changes caused by pulsations (from the mean value), φpul is a pulsation phase, and Δφpul is a phase shift which expresses by how much the moment of maximum radius of the M star precedes the moment of maximum of its brightness. The period and zero moment of pulsation maxima were adopted from the ephemeris (Eq. (4)). Several parameters were treated as fixed, not subject to change in the process of solution. The effective temperature of the hot component Thot was set to 7500 K and its radius Rhot = 58   R was estimated with the Stefan-Boltzmann law by adopting the luminosity Lhot extracted from the SED (see Table 6). The broad, atmospheric parts of the eclipse were included by introducing a light-absorbing envelope that changes the density distribution as a function of distance from its centre r as r-2. Its radius Renv was roughly estimated from the total duration of the atmospheric eclipse and fixed at 440 R. We assumed that the intensity of radiation I, measured after passing through the envelope, changes from an initial value I0 according to (A.3)where τ(r) is an optical depth calculated (A.4)The absorption coefficients κV and κI were chosen via visual comparison of synthetic curves with photometric observational data and were fixed at the values 0.03 and 0.015 for V and I bands, respectively, assuming that the integration constant C is equal to 1. We carried out the solutions for several temperatures of the cool component Tcool in the range from 2600 K to 3500 K with a step of 100 K. For each Tcool value we calculated the star’s radius from the Stefan-Boltzmann law using the luminosity extracted from the SED from the spectrum of RX Boo fitted in the place of the cool component (see Table 6). The adjustable, free parameters were: the moment of mid-eclipse defined as the time at which the centres of stellar disks are at the minimum

separation JD0, the reciprocal tangential velocity of the stars Vt, the semi-amplitude of brightness changes caused by pulsations ΔMV, ΔMI for V and I bands, respectively, and three parameters described earlier: D, ΔR, and Δφpul. To reduce the number of free parameters we combined the semi-amplitudes of ΔMV, ΔMI using an empirical relationship ΔMI ≈ 0.478   ΔMV that we found when analysing the AAVSO photometric data for Miras- and SR-type stars (see Fig. A.2). The best solution, i.e. corresponding to the minimum value of the sum of square residuals, was obtained for Tcool = 2800 K and . The complete set of resulting parameters of the model is shown in Table A.1.

thumbnail Fig. A.1

Schematic explanation for the geometric parameters of the model.

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thumbnail Fig. A.2

Dependence of the pulsation amplitude of Mira/Semi-Regular pulsating stars in the V photometric band on the amplitude in the I band obtained from AAVSO light curves.

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Table A.1

Input (left) and output (right) parameters of the model of the 2007/8 eclipse of V383 Sco.

All Tables

Table 1

Journal of spectroscopic observations.

Table 2

Moments of minima and the O–C residuals estimated according to the linear ephemeris JDmidecl = 2   415   450 + 4900 × E.

Table 3

Frequencies and periods obtained with Fourier analysis of the out-of-eclipse V-band light curve.

Table 4

Pulsational maxima and the O–C residuals estimated according to the initial ephemeris in Eq. (3).

Table 5

Photometry used to calculate the spectral energy distribution (SED).

Table 6

Luminosities of components of the SED of V383 Sco.

Table 7

Estimated minimum values of the eclipse depths () and the pulsation amplitudes δ.

Table A.1

Input (left) and output (right) parameters of the model of the 2007/8 eclipse of V383 Sco.

All Figures

thumbnail Fig. 1

Top: V- and I-band light curves from ASAS-3. Bottom: the V − I colour index. Crosses denote colour indices V − I obtained by interpolation in the cases when V and I measurements were not simultaneous, i.e. not taken during the same nights. Arrows and numbers show the moments of the pulsation maxima visible in V- and I-band light curves, used for timing analysis of pulsations (see Sect. 3.2). The dotted arrow shows a postdiction for the pulsation maximum at epoch Ep = 10 (corresponding to a maximum that occurred during the most recent eclipse – epoch E = 8) according to the ephemeris from Eq. (4).

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In the text
thumbnail Fig. 2

Expanded view of the 2007/8 eclipse of V383 Sco. The symbols are the same as in Fig. 1. The vertical dashed line marks the mid-eclipse moment2 (Table 2).

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In the text
thumbnail Fig. 3

O–C diagram with a linear (solid line) fit to the residuals. The moment of minimum at epoch 1 (vertical bar) is an estimated figure because adequate data were not available.

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In the text
thumbnail Fig. 4

Light curves of four eclipses phased with a constant period according to the new linear ephemeris (Eq. (2)). The open symbols denote archival photographic photometric data (squares: E = 0, circles: E = 1, diamonds: E = 2). The V-band ASAS data (filled circles: E = 8) have been shifted arbitrarily by  + 0.6 for facilitate comparison with the light curves from the different photometric systems.

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In the text
thumbnail Fig. 5

Power spectra obtained using Fourier analysis on out-of-eclipse V-band (top) and I-band (bottom) data (see Sect. 3.2).

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In the text
thumbnail Fig. 6

ASAS-3 V-band (top), V − I (middle), and I-band (bottom) light curves outside of the eclipse, phased with period . The small amplitude of the observed V variation and its strongly non-sinusoidal shape with a wide minimum are related to the presence in the system of the second component which strongly dominates in this wavelength range (see Sect. 3.7).

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In the text
thumbnail Fig. 7

O–C diagram for the pulsation maxima with a linear (solid line) fit to residuals.

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In the text
thumbnail Fig. 8

Dereddened spectral energy distribution of V383 Sco obtained from the data given in Table 5 (points). The dashed line (blue in the electronic version of this paper) shows the F0I-type supergiant from the 1993 Kurucz Stellar Atmospheres Atlas (Kurucz 1993). The dot-dashed line (red in the electronic version) marks the SED of the SRb-type pulsating star RX Boo, built using the VOSpec virtual observatory tool. The dotted line (magenta in the electronic version) shows a black body with T = 500 K. The continuous line (orange in the electronic version) shows the sum of all 3 components appropriately rescaled (by a factor of 1.9 × 10-20 for the Kurucz model of an F0I star and by a factor of 1.3 × 10-4 for the observational SED of RX Boo) to fit the SED of V383 Sco. Triangles denote IRAS data with very low flux density quality (marked with the number 1 in the revised version of the IRAS Point Source Catalog), i.e. upper estimates: from the IRAS FSC catalogue (filled triangles) and from the IRAS PSC catalogue (open triangles).

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In the text
thumbnail Fig. 9

Observed low-resolution spectrum of V383 Sco (thick line) compared with the model (thin line) that is a superposition of F0 and M1 supergiant spectra (grey lines). The observed spectrum and the model are dereddened by EB − V = 0.5.

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In the text
thumbnail Fig. 10

HARPS spectrum of V383 Sco in the region 4450–4550 Å compared with the spectra of HD 90772 and HD 75276 from the UVES library (Bagnulo et al. 2003) and ε Aur from the ELODIE archive (Moultaka et al. 2004).

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In the text
thumbnail Fig. 11

Hα and Hβ profiles in the HARPS spectrum of V383 Sco.

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In the text
thumbnail Fig. 12

Profiles of Ti ii 5186 Å, Y ii 5200 Å, and Ba ii 6142 Å in the spectrum of V383 Sco in which weak emissions on the shortward edges and wide absorption wings are clearly visible.

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In the text
thumbnail Fig. 13

Comparison of the region around Na i 5890 Å in the HARPS spectra of V383 Sco (stellar absorption at around  + 90 km s-1) and V381 Sco (stellar absorption at about  − 180 km s-1).

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In the text
thumbnail Fig. 14

[O i] 6300 Å and 6364 Å forbidden emissions in the spectrum of V383 Sco.

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In the text
thumbnail Fig. 15

Synthetic V- and I-band light curve and V − I colour index curves obtained by modelling, plotted with solid lines (green, red, and orange, respectively in the electronic version of this paper), in comparison with observational photometric data. Between the flux and colour curves, residuals are shown. The model assumes a temperature of the cool star Tcool = 2800 K. The expanded view of the out-of-eclipse V and I light curves (box with a dashed line) is shown in Fig. 16.

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In the text
thumbnail Fig. 16

Expanded view of the synthetic fit to the out-of-eclipse V and I light curves and residuals.

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In the text
thumbnail Fig. A.1

Schematic explanation for the geometric parameters of the model.

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In the text
thumbnail Fig. A.2

Dependence of the pulsation amplitude of Mira/Semi-Regular pulsating stars in the V photometric band on the amplitude in the I band obtained from AAVSO light curves.

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In the text

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