Spectroscopic and photometric study of the eclipsing interacting binary V495 Centauri
J. A. Rosales R. E. Mennickent G. Djurašević I. Araya M. Curé
astro-ph.SR





2cm

[
\begin@twocolumnfalse
	Spectroscopic and photometric study of the eclipsing interacting binary V495 Centauri
    
Rosales, J.A.^1, Mennickent, R.E.^1, Djurašević, G.^2,3, Araya, I.^4,5, Cur, M.^4
 

^1 Departamento de Astronoma,Universidad de Concepcin,Casilla 160-C, Concepcin, Chile.

^2 Astronomical Observatory, Volgina 7, 11060 Belgrade 38, Serbia.

^3  Isaac Newton Institute of Chile, Yugoslavia Branch.

^4 Instituto de Fsica y Astronoma, Facultad de Ciencias, Universidad de Valparaso, Chile.

^5 Ncleo de Matemticas, Fsica y Estadstica, Facultad de Ciencias, Universidad Mayor, Chile.


    
=======================================================================================================================================================================================================================================================================================================================================================================================================================================================================================================



Double Periodic Variables (DPV) are among the new enigmas of semi-detached eclipsing binaries. These are intermediate-mass binaries characterized by a long photometric period lasting on average 33 times the orbital period. We present a spectroscopic and photometric study of the DPV V495 Cen based on new high-resolution spectra and the ASAS V-band light curve. We have determined an improved orbital period of 33.492 ± 0.002 d and a long period of 1283 d. We find a cool evolved star of M_2=0.91± 0.2 M_, T_2= 6000± 250 K and R_2=19.3 ± 0.5 R_ and a hot companion of M_1= 5.76± 0.3 M_, T_1=16960± 400 K and R=4.5±0.2 R_. The mid-type B dwarf is surrounded by a concave and geometrically thick disc, of radial extension  R_d= 40.2± 1.3 R_ contributing ∼ 11 percent to the total luminosity of the system at the V band. The system is seen under inclination 84.^∘8 ± 0.^∘6 and it is at a distance d= 2092 ± 104.6 pc. The light curve analysis suggests that the mass transfer stream impacts the external edge of the disc forming a hot region  11% hotter than the surrounding disc material.  The persistent V<R asymmetry of the Hα emission suggests the presence of a wind and the detection of a secondary absorption component in He I lines indicates a possible wind origin
in the hotspot region. 





Keywords: Binaries: general- stars: early-type: emission-line: evolution-stars: mass-loss
 

	
\end@twocolumnfalse

]




1 INTRODUCTION


Double Periodic Variables are interacting binary stars of intermediate mass that show two closely linked photometric variations being the long period roughly  33 times longer than  the orbital period (<cit.>,<cit.>). The nature of the second period is unknown but suspected to reflect the strength variations of a wind generated in the stream-disc impact region (<cit.>, <cit.>).   Also, DPVs are considered as one specific evolutionary step for more massive Algols, one possibly involving mild mass transfer and systemic mass loss (<cit.>).  An interesting property of these objects is the surprising constancy of their orbital periods, which is not expected in Algols undergoing RLOF mass transfer (<cit.>). Also the DPVs seem to be hotter and more massive than classical Algols and seem to have always a B-type component;  their orbital periods typically run between 4 and 100 days. DPVs have been found in the Galaxy (MW), Large Magellanic Cloud (LMC) and the Small Magellanic Cloud (SMC). The components of DPVs are denoted by the suffixes g and d appended to the parameter designation, g denoting the primary component (C1) or gainer (hot) and d the secondary component (C2) or donor (cold), respectively. Recently, a mechanism based on cycles of a magnetic dynamo in the donor star, the so called Applegate mechanism, was proposed as an explanation for the DPV long  cycles consistent with the almost constancy of the orbital period (<cit.>).

The importance of studying DPVs is the fact that they are interesting astrophysical objects for the study of mechanisms of mass transfer in semidetached binaries, they share similarities with some Algol systems, and can potentially provide us a wealth of information on the density, tidal friction, wind processes and stellar dynamos in massive close binaries (<cit.>,<cit.>). Hitherto only a couple of DPVs, less than 5%, have been studied spectroscopically (e.g <cit.>; <cit.>a,b;<cit.>;<cit.>;<cit.>).


The eclipsing interacting binary V495 Cen (ASAS ID 130135-5605.5, α_2000= 13:01:35.0, δ_2000= -56:05:30.0, V= 9.95 mag, B-V= 0.52 mag, spectral type Be)[http://simbad.u-strasbg.fr/simbad] was discovered as a DPV by <cit.>, who find that the system is characterized by a long photometric cycle of 1283 d, resulting in the system with the longest period among Galactic DPVs. V 495 Cen has been poorly studied and shows an orbital period of 33.48177 d in the ASAS[http://www.astrouw.edu.pl/asas/] catalogue (<cit.>). Our study will contribute to the overall understanding of DPVs, especially those of long orbital period.
 
 In Section 2 we present a photometric analysis of V495 Cen. In Section 3 we give a brief review of the  methods of data acquisition and data reductions. In Section 4 we present the orbital and physical parameters of the system. In Section 5 we model the light curve with a special code including the light contribution of both stars and the accretion disc, and get  luminosities, radii, temperatures, surface gravities, masses and the system's inclination.  
In Section 6  we present an analysis yielding the reddening and distance of the system. Finally,  the main results of our research are summarized in Section 7. 






0.5cm 
2 PHOTOMETRIC EPHEMERIDES


We have re-analyzed the ASAS light curve considering only those 622 better quality data points labeled as A-type by ASAS (Fig.  1) and rejecting outliers points, contrary to the previous analysis (<cit.>) in which were considered data points labeled as A-type and B-type, the best and normal quality respectively. The data set was analyzed by Phase Dispersion Minimization (PDM-IRAF, Stellingwerf 1978)[IRAF is distributed by National Optical Astronomy Observatories, which are operated by Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation], revealing an orbital period P_o=33.492 ± 0.01613 d (the error corresponds to the difference between the first approximation by PDM and the best value found by visual inspection of the periodogram minimum) and the epoch of minimum HJD = 2454609.71384 ± 0.032 d. We have noticed that the orbital period shown in ASAS (33.4817 d) does not match the value obtained by us. This is probably due to the fact that the algorithm used by PDM to search for periods is different to that used by ASAS at the moment of carrying out the periodogram.

With the long and short periods we disentangled the light curve with a code specially designed for this purpose by Zbigniew Kołaczkowski. The code fits the data with a Fourier series consisting of fundamental frequencies and harmonics and removes them. As a result we obtain the cleaned light curve with no frequencies and two light curves for the isolated orbital and long cycles as shown in (Fig. 2). The process reveals an orbital modulation typical of an eclipsing binary type (EB), with rounded inter-eclipse
regions revealing proximity effects, and a longer cycle characterized by a quasi-sinusoidal variability typical of Double Periodic Variables,  whose full amplitude in the V-band is ∼ 20 % with respect to the total brightness. We determined the following ephemeris for the light curves:







    HJD_min, orbital =2454609.713(32)+33.492(16) ×E,




    HJD_max,long = 2454926.699(42) + 1283 ×E,



These are used for the spectroscopic analysis in the rest of the paper. 









0.5cm 
3 SPECTROSCOPIC OBSERVATIONS


To investigate the fundamental properties of V495 Cen we collected a series of high-resolution optical spectra between February and May 2015  with the CHIRON[http://www.ctio.noao.edu/noao/content/CHIRON] spectrograph. We obtained 10 spectra with R ∼ 25000 (fiber mode) and 30 spectra with  R ∼  80000 (slicer mode). The spectral regions covered were 4500-8500 Å and 4580-8762 Å at the fiber and slicer mode, respectively, with a typical signal to noise ratio (SNR) of ≈ 160.

The corrections with flat and bias, wavelength calibration and order merging were done with IRAF. We also normalized all spectra to the continuum and corrected them to the heliocentric rest frame. 

The spectra obtained with CHIRON are not sky-subtracted. We have not flux calibrated our spectra but this does not affect the strength measurements and radial velocities included in this paper.
 



0.7cm 
4 SPECTROSCOPIC ANALYSIS


0.5cm
4.1 Spectral disentangling 


The spectra were disentangled with a method that is quite good for separating the absorption-lines widths of both stellar components <cit.>. This method turns to be effective when working with absorption-lines widths less than the radial velocity amplitude of the corresponding star, i.e. for most lines except for the strong and wide Hα emission. For this reason all spectral regions of interest  except the Hα region were disentangled. However, as we can see in Fig. 3, the underlying donor absorption is small compared with the strong double peak Hα emission.  The presence of emission in Hα confirms the interacting binary nature for this system and suggests it is in a semi-detached stage and possesses a circumstellar disc. We notice the deep central absorption between the emission peaks revealing large amounts of obscuring material through the line of sight and also the larger intensity of the red peak at this specific epoch. 








0.5cm 
4.2 Determination of donor physical parameters


To determine the physical parameters of the secondary star, we compared our observed average donor spectrum with synthetic spectra constructed with SPECTRUM[http://www.appstate.edu/ grayro/spectrum/spectrum.html] which uses atmospheric models computed with grids of ATLAS9[http://wwwuser.oats.inaf.it/castelli/grids.html] model atmospheres (<cit.>). The models used are in Local Thermodynamic Equilibrium (LTE). 

The theoretical models were calculated for effective temperatures from 4250 to 19000 K with steps of 250 K, surface gravities from 0.5 to 3.5 dex  with steps of 0.5 dex, solar metallicity, micro-turbulent velocity from 0.0 to 2.0 km s^-1 with steps of 1.0 km s^-1 and macro-turbulent velocity from 1.0 to 20.0 km s^-1 with steps of 1.0 km s^-1. 





Our analysis aimed to find the best synthetic spectrum was based on a chi-square optimization algorithm. The method allowed  simultaneous determination of various parameters involved with stellar spectra and consists of the minimization of the deviation between the theoretical normalized flux distribution and the observed normalized spectrum, the latter corrected by a veiling factor  η. The synthetic spectrum depends on the stellar parameters, such as effective temperature T_eff, surface gravity logg, rotational velocity vsini, micro-turbulence v_mic and macro-turbulence velocity v_mac. The optimization of the method converged successfully at T_2= 6000 ± 250 K, logg_2= 2.5 ± 0.5, v_2rsini= 26 ± 1 km s^-1, v_mic=0.0 ± 0.5 km s^-1 , v_mac= 11 ± 0.5 (Fig. 4, 5). We observed a slow divergence of the method for higher temperatures and surface gravity.  





0.5cm 
4.3 Radial velocities for the donor


The radial velocity of the donor was measured by cross correlation of a spectral region plenty  of 
metallic lines with a spectrum chosen as a template. Then the velocities were translated to the absolute 
heliocentric system by adding the template velocity obtained by simple Gaussian fits to selected absorption lines. The donor velocities are given in Table 4.

To find the orbital elements of V495 Cen, we used the genetic algorithm PIKAIA [http://www.hao.ucar.edu/modeling/pikaia/pikaia.php] developed by <cit.>. The method consists in finding the set of orbital parameters that produces a series of theoretical velocities that minimize the function χ^2, defined as:


χ^2(P_o,τ,ω,e,K_2,γ)

    = 1/N-6∑_j=1^N(V_j^obs-V(t_j;P_o,τ,ω,e,K_2,γ)/σ_j)^2,



Where N is the number of observations, V_j^obs is the radial velocity observed in the data set, and V(t_j; P_o, τ, ω, e, K_2,γ) is the radial velocity at time t_j given the parameters. P_o is the orbital period, ω the periastron longitude, τ the time of passage per the periastron, e the orbital eccentricity, K_2 the half-amplitude of the RV for the donor, and γ the velocity of the system centre of mass. The radial velocity is given by:



    V(t)=γ + K_2(cos(ω+θ(t)) +e cos(ω)),





(<cit.>, eq. 2.45) where θ is the true anomaly obtained solving the following two equations involving the eccentric anomaly E:



    tan(θ/2)= √(1+e/1-e)tan(E/2),




    E-e sin(E) = 2π/P_o(t-τ),



(<cit.>, eq. 2.35) To calculate this, we must first solve equation (6) for E, then solve Equation (5) for θ, and finally solve equation (4) for V(t_j; P_o, τ, w, e, K_2,γ). The error was estimated using Monte Carlo simulations, by perturbing the best-fitting solution obtained with PIKAIA and computing a χ^2 of these perturbed solutions. We note that our eccentricity solutions give a small value of e=0.007. We have implemented a test for the significance of the eccentricity (<cit.>). According to this test the condition p_1 < 0.05 means a significant ellipticity, where:
 



    p_1= (∑ (o-c)^2_ecc/∑(o-c)^2_circ)^(n-m)/2




where ecc indicates the residuals of an eccentric fit, and circ the residuals of a circular fit, n the total number of observations, m the number of free parameters in an eccentric fit. Finally p_1 is the probability of falsely rejecting the circular orbit. We obtained p_1=0.0178  which is in principle compatible with a very small ellipticity. However, the 1σ error still allows a circular orbit. The orbital parameters and errors are given in Table 1 and the best fit to the radial velocities is shown in Fig. 6.




The secondary star dominates the metallic and Balmer line spectrum. It was quite difficult to find lines from the gainer. However, we found a little group of characteristic lines that represents the movement of the gainer star and we choose to use the most characteristic lines for a star of B spectral type, these were found as components of the He I 5015 and He I 5875 lines (Fig. 7), as also happened in the DPV HD 170582 (<cit.>). It should be noticed that in HD 170582 the origin of the anomalous He I component is placed near the hotspot region, consistent with a hotspot wind.  The radial velocities for the C1 component (gainer) of the  He I 5015 line and the C2 component (donor) of the He I 5875 line  were measured with a gaussian fit and are shown in Fig. 8, along with fits for both orbits. We observe that the technique of gaussian fit provides velocities with a larger scatter than those obtained with the cross correlation method.







The RV of the component C1 for He I 5015 is fitted with a sine function of amplitude 16.92 ± 0.84 km s^-1 and zero point 0.00± 0.58 km s^-1. It moves in opposition to the donor and it could reflect the motion of the gainer around the binary center of mass.  The RV of the component C2 of He I 5875 was fitted with a sine of amplitude 106.80 ± 0.39 km s^-1 and zero point -0.91± 0.28 km s^-1, both solutions assume a circular orbit   (Fig. 8). The reason for which the radial velocity curve of the component C1 has a small phase lag might be due to the brightness contribution of the disc in the absorption line.




0.5cm
4.4 The Hα emission line profile


An important result of our spectroscopic analysis is that the Hα profile shows a double peak emission with the red peak always of larger intensity than the blue one while  the central absorption does not follow the motion of none of the stellar components (Fig. 9). Representative Hα profiles at some epochs and intensity/width measurements performed at the mean profile are shown in Fig. 10.
The full width of the line at the level of the continuum is 776 km s^-1.

The double emission and the deep absorption core is typical of a disk seen at large inclination, however the persistence of the larger red peak is not usual. In binaries with accretion discs, it is normal to find  variations of the relative intensities of the emission peaks. This is quantified with the parameter V/R  usually defined as V/R = (I_v-1)/(I_r-1), i.e. in terms of the peak intensities relative to the normalized continuum. Due to bright zones in the disc of Algols this ratio tends to vary cyclically with the orbital period.  However, in V495 Cen we observe almost always V < R. This can be interpreted as evidence of material escaping from the system, i.e. a wind emerging probably from the hotspot region. We also notice here that the strength of the Hα emission is large compared with other cases of DPVs. The equivalent width (EW) of the Hα line shows a clear periodic behavior and varies with the orbital period. A least square fit of sine type, with the orbital period fixed at 1.0 P_o (solid line in Fig. 11), represents the observed data quite well, and shows a mean EW of 9.517 ± 0.063 Å and an amplitude of 1.985 ± 0.663 Å. The phase of observations is such that the minima occur at orbital phase 0.15. A probable interpretation is that there is an asymmetric distribution of Hα emission on the surface of the star or material escaping from the hotspot, which is assumed to be synchronously rotating. Also we note the spectral resolution difference for each equivalent width obtained by fiber mode (R∼25000, blue dots) and slicer mode (R∼80000, black dots). 












0.5cm
4.5 Gainer, mass ratio and circumstellar matter


In some spectroscopic binaries stars, spectral lines from both stars are visible and the lines are alternately double or single. These systems are known as double-lined spectroscopic binaries, denoted as SB2. But in other systems, the spectrum of only one component is seen and the lines in the spectrum shift periodically towards the blue and red. Such stars are known as single-lined spectroscopic binaries or SB1. V495 Cen is probably an SB2 system and we have  assumed that the velocity obtained from the C1 He I 5015 line represents the gainer orbital motion hence we have inferred a mass ratio q=0.158 ± 0.008.

Now let's investigate if this value is compatible with synchronous rotation for the secondary star filling the Roche lobe; we use the following equation valid for semi detached binaries:



    v_rotsini/K≈ (1+q) 0.49 q^2/3/0.6q^2/3+ln(1+q^1/3)




(<cit.>, eq. 3.5 and 3.9). Therefore with K_2=106.880 km s^-1 and v_rot2sin i= 26 km s^-1 with their respective errors and considering the 1% of accuracy of equation 8, we get q= 0.119 ± 0.037, which is very close to the q value derived from the RV half-amplitude of the helium components (Fig. 12); therefore we can say that the 5015 HeI line comes from the gainer star. 
Since synchronism is expected for such a short orbital period binary due to the influence of dynamical tides (<cit.>),  from now on we assume the donor rotating synchronously. 

Another important point is the relation between the mean density ρ (in solar units) of a star that just fills its Roche lobe and the effective radius R_L (<cit.>) we quantified as the product between the critical orbital period (days) and mean density (g cm^-3) as:
 



    P_cr√(ρ)= (3π/G)^1/2(q/1+q)^1/2 x_L^-3/2




where x_L is in units of the orbital separation. Assuming q=0.158 from the above RV study, we have obtained the ratio between the effective radius and the orbital separation R_L/a=0.236 ± 0.155 (<cit.>, eq. 3.5), and estimated the critical period P_cr√(ρ)= 0.44 ± 0.06, which is the shortest period possible for a binary of given mass ratio into which a star of given  mean density ρ can be fitted without overflowing its Roche lobe (<cit.>, eq. 3). Both values are significant figures to understand the evolution of the Double Periodic Variable V495 Cen.






0.5cm
4.6 Mass constraints from spectroscopy


The mass function represents the minimum possible mass for the unseen star, for a system with unknown orbital inclination i and elliptical orbit, it may be expressed as:



    f= 1.0361 × 10^-7 (1-e^2)^3/2( K_2/km s^-1)^3P_o/dayM_




(Hilditch 2001, eq. 2.53)
Using q = 0.158 (donor rotating synchronously) and an elliptical orbit with e= 0.007 from the above RV study, we obtained a mass function f=4.235 ± 0.015 M_. Assuming an orbital inclination i = 84.8 (see the next section), the respective masses are M_d= 0.911 ± 0.420 M_ and M_g= 5.753 ± 0.451 M_. These are preliminary results. Both are compatible with results obtained, independently, through LC analysis in the next section.

 



0.5cm 
5 LIGHT-CURVE MODEL AND SYSTEM PARAMETERS


0.5cm
5.1 The fitting procedure


We find the main physical parameters for the stellar components fitting the V-band light curve with the aid of an inverse-problem solving method based in the algorithm developed by <cit.>. We have assumed a semi-detached system with the donor filling the Roche lobe, and we adopted a configuration that included an optically thick accretion disc around the gainer star.

Our model considers a hot spot located on the edge side of the disc, and it is described by the ratio of the hot spot temperature and the unperturbed local temperature of the disc, the angular dimension and longitude (in arc degrees). Synchronous rotation is assumed for both stellar components. The model used follows results of hydrodynamical simulation of gas dynamics in interacting close binary stars by <cit.> and has been tested in studies of DPVs (e.g. <cit.>,<cit.>,<cit.>). 

We fixed q=0.158 and T_c=6000 K based on the spectroscopic study previously presented. In addition, we set the gravity darkening exponent and the albedo of the gainer to β_h= 0.25 and A_h=1.0 in accordance with von Zeipel's law for radiative shells and complete re-radiation; for the donor we set β_c=0.08 and A_c=0.5, as is appropriate for stars with convective envelopes according to <cit.>,<cit.>,<cit.>.









0.5cm	
5.2 The best light-curve model



The best fit model for V495 Cen contains an optically and geometrically thick disc around the gainer star; the stellar, orbital and disc parameters are given in Table 2.

The system inclination angle is  i=84^∘.8 ± 0.6 and the disc radius  R_d=40.2 ± 1.3 R_, which is 9 times larger than the radius of the main star  R_h=4.5. The disc has central vertical thickness d_c=0.6 R_ and edge thickness d_e=3.8 R_, i.e. it has a concave form. The temperature of the disc increases from T_d=4040 K at the outer edge to T_h= 16960 K at the inner edge where it is in thermal and physical contact with the gainer. The surface gravity of the secondary component (donor) is logg_c=1.83± 0.1  i.e. almost 1σ below the spectroscopic value 2.5 ± 0.5. The temperature of the hotspot is T_hs≈ 4484 K i.e. is 11 percent higher than the disc edge temperature.

The light curve and the fit, the O-C curve, residuals, and the individual flux contributions of the donor, disc and the gainer, are shown in Fig. 13. Also we show a view of the model at orbital phases 0.05, 0.55 and 0.80. 

The obtained solution can be treated as unique with a model of the system with  fixed initial parameters (mass ratio and temperature of the donor). The uniqueness of the obtained optimal solution is checked by the solving the inverse problem of the light curve interpretation by applying the Simplex algorithm. Even, if we vary the initial system parameters in the initial simplex in the interval of 5-10% of the optimal solution, the optimization process converges to the obtained solution within a given errors bars. Of course, if the initial fixed parameters are used with a larger uncertainty, the estimated errors bars of the free model parameters increase too.






0.5cm 
6 REDDENING, DISTANCE  AND  SPECTRAL ENERGY DISTRIBUTION



0.5cm
6.1 Distance determination


In order to determine the distance to the system we have applied a standard method based on the distance modulus  to both binary components observed in the V band (Clausen, 2004):


    (m_d,g-M_d,g)_0    =      5log(R_d,g/R_)+(m_d,g-A_V)
          - M_bol+ 10log(T_d,g/T_)
           + BC_d,g



where A_V is interstellar absorption and BC the bolometric correction; the apparent magnitudes, radii and effective temperatures of gainer and donor are represented with sub-indexes g and d, respectively. 

Maps of Galactic dust and extinction in the region of V495 Cen are available on-line and give  E(B-V)_S&F= 0.2820± 0.0069 (<cit.>) and E(B-V)_SFD= 0.3279±0.0080 (<cit.>).  Assuming a visual extinction to reddening ratio A_V/E(B-V)= 3.1 we get  A_V_S&F=0.8460  ± 0.006 and A_V_SFD=0.9837  ± 0.008, where we preferred the most recent determination of the reddening in the direction of the target.

The model of the V-band light curve (Fig. 13) shows that the flux of the system at the V band at quadratures ϕ_o=0.25 and ϕ_o=0.75 is f_t=1.77 ± 0.05 while the individual flux contribution of the donor is f_d=1.0 ± 0.05, of the gainer is f_g=0.57± 0.05 and the disc is f_disc=0.2± 0.05. This means that at quadrature the donor, gainer and disc contribute 
57, 32 and 11 % to the total flux at the V band, respectively.

The apparent magnitude of the donor and gainer are derived from:


    m_d-m_t= -2.5 log(f_d/f_t)




Considering that the observed apparent magnitude of the system is m_t(V)= 9.95 ± 0.03, and the flux fractions given above, we find m_d=10.569 ± 0.044 mag and m_g=11.180 ± 0.037 mag. 

The bolometric corrections were taken from <cit.> and the solar bolometric magnitude was taken from <cit.>. We use logT_d= 3.778 ± 0.042 and  BC_d= -0.045 ± 0.110 for the donor and logT_g =4.229 ± 0.024 and BC_g= -1.525 ± 0.079 for the gainer, and the individual distance are d_d=1923.092 ± 137.538  pc d_g= 2401.043 ± 187.169  pc. The difference in the distance for donor and gainer can be interpreted as an estimate of the intrinsic error of the method due to the gainer star is hidden by the disc, causing an extra reddening that is not the same reddening caused by the total absorption across the line of sight to V495 Cen. Therefore we have averaged the values obtaining for the system distance:


    d(pc)   =     2162.068± 324.707   pc


Since the maximum absorption through the galactic gas column along the system line of sight has been considered for this calculation, this distance must be considered as a lower limit only.





0.5cm
6.2 Position in the H-R Diagram


From the above sections it is clear that the less massive star is the more evolved star of the binary pair, as usual in Algol systems. This can be explained by the mass transfer that has inverted the system mass ratio in such a way that the less massive star started as the more massive one some time ago and evolved first
until filling its Roche lobe just before start mass transfer onto the now more massive stellar component.
  
In order to study the position of V 495 Cen in the luminosity-temperature diagram, we have chosen as comparison models based on a grid that is part of the large database of Geneva stellar  models[https://obswww.unige.ch/Recherche/evoldb/index/]. We consider single star models with metallicity Z= 0.014 (<cit.>) as a first approximation to understand the evolutionary stage of V 495 Cen (Fig. 14).

We observe that the gainer (5.76 M_⊙) is located outside the main sequence, showing a luminosity similar to a single star of that mass. It is located near the start of the stellar contraction stage for the evolutionary track of 6.0 M_⊙. The core fractions of hydrogen and helium are 0.72 and 0.266, respectively. On the other hand, the donor (0.91 M_⊙) has a huge luminosity compared with a star of  similar mass and fits the evolutionary track of a 4.35 M_⊙ single star in the phase of hydrogen shell narrowing (Iben 1967, Fig.14). The high luminosity donor can be understood in terms of an evolved giant filling its Roche lobe that has transferred part of its mass onto the gainer forming the accretion disc.


Also we have inferred the color of each star as function of effective temperature (<cit.>). (B-V)_o_g=-0.200 ± 0.001, (B-V)_o_d=0.571 ± 0.001 and compared the intrinsic color of both obtaining the spectral classes of each star. Hence, considering all above we find that the gainer is an early B3 V type, and the donor corresponds to late F9 IV, based on the intrinsic color analysis (<cit.>).









0.5cm
6.3 Spectral Energy Distribution (SED).



The determination of the physical parameters of astronomical objects from observational data is frequently linked with the use of theoretical models as templates. In order to obtain the broad-band photometric fluxes we compiled all the information available for V495 Cen and built the SEDs with the aid of the Spanish Virtual Observatory SED Analyzer [http://svo2.cab.inta-csic.es/theory/vosa4/] (<cit.>). We performed a statistical test to decide which synthetic model best reproduces the observed data. The provided “best” fitting model is the one that minimizes the value of reduced χ^2, considering the composite flux as:



    f_λ=f_λ,010^-0.4 E(B-V)[k(λ- V)+R(V)]




where



    f_λ,0=(R_2/d)^2[(R_1/R_2)^2f_1,λ+f_2,λ]




and  f_1 and f_2 are the fluxes of the  star and secondary star respectively, R(V)≡ A(λ)/E(B-V) is the ratio of reddening to extinction at V, d is the distance to the system and  R_1/R_2 is the ratio of the primary star radius to the secondary star radius, k(λ-V)≡ E(λ-V)/E(B-V) is the normalized extinction curve and was calculated in two steps:





    k={[            ελ^-β-R_v                   if                                         x                < 0.3; R_v(a(x)+b(x)/R_v-1)                   if                0.3 ≤                    x                ≤ 8.0;                      ].




Where x≡ 1/λ(μ m^-1), ε=1.19, β=1.84 and R_v=3.05, with formal errors of about 1% (<cit.>), however, this value of β is not applicable at wavelengths beyond 5 μ m. The parameters a(x) and b(x) are the parametrization coefficients for the extinction law from <cit.>.

The theoretical spectra  were taken from the library of ATLAS9 Kurucz ODFNEW /NOVER models. We used fluxes calculated with metallicity index [M/H]= 0.0,  the micro turbulence velocity for both stars was of 0.0 kms^-1. We considered the reddening produced by galactic dust as discussed in Section 3.1. 
 The parametric space of the models was reduced to consider:
donor temperature between 4250 and 8000 K with step of 250 K and surface gravity between 0.5 and 3.0 with steps 0.5, gainer temperature between 16000 and 17000 K with step of 1000 K and surface gravity 
between 3.5 and 4.5 with step of 0.5. We fixed the radii and considered all values of the photometry obtained with VOSA from 4280 Å until 220883 Å. 


We implemented a chi-square adjustment considering a distance with a 5% error, for 7 degrees of freedom, i.e. x_critic  0.95,7^2. We found a distance of d = 2092 pc with a χ^2= 1.447. Due to χ^2 < χ_critic  0.95,7^2, then the solution is accepted and the solution obtained in Section 6.1 is consistent with the current results (Fig. 15). 

In addition, we observe a slight infrared excess at wavelengths longer than 10000 Å, which cannot come from the main star itself but from cooler emitting material. This excess is probably formed in the same circumstellar material responsible of the strong Hα line emission. 






0.5cm 

7 CONCLUSIONS


In this work we have presented a detailed spectroscopic and photometric analysis of V 495 Cen, 
the hitherto longest orbital period Galactic DPV. We conclude:




  * From the photometric study we have determined that V495 shows an orbital period of P_o= 33.492 ± 0.002 d and a long period of P_l= 1283 d. 



  * We have modeled the spectrum and light curve and the best system and orbital parameters we found are given in Table 2. For instance, we determined the mass of the secondary star  M_2= 0.91 ± 0.2 M_, its temperature T_2=6000 ± 250 K and radius R_2=19.3 ± 0.5, the mass of the primary  M_1=5.76 ± 0.3 M_, its the temperature T_1=16960 ± 400K and  radius R_1=4.5 ± 0.2 R_. 


  * The early B type dwarf is surrounded by an optical and geometrically thick  accretion disc of radial extension R_d=40.2± 1.3 R_. 


  * At quadratures the donor, gainer and disc contribute 57, 32 and 11%  to the total flux at the V band, respectively.


  * The best model shows a hot spot located in the outer edge of the disc, 11% hotter than the surrounding disc and displaced 18.2 degree for the line joining the star centers in the direction of the orbital motion.


  * The spectral energy distribution shows infrared excess indicating the presence of a circumstellar material. 


  * We found a lower limit to the distance  to V495 Cen of 2092 pc with a χ^2=1.447 for 7 degrees of freedom and 95% of confidence level, i.e. the distance for V495 Cen is 2092 ± 104.6 pc. 



0.5cm

ACKNOWLEDGEMENTS


This investigation is based on observations conduced under CNTAC proposal CN2015A-123. This publication makes use of VOSA, developed under Spanish Virtual Observatory project supported from the Spanish MICINN/MINECO trough grant AyA2008-0256, Ay2011-24052. Also we made use of the SIMBAD database, operated at CDS, Strasbourg, France. This research was funded in part by scholarship from the Faculty  of Physical Sciences and Mathematics of the Universidad de Concepcin. J.R. and R.E.M. gratefully acknowledge support from the Chilean BASAL Centro de Excelencia en Astrofsica y Tecnologas Afines (CATA) grant PFB-06/200. J.R. thanks the SOCHIAS grant through Gemini-Conicyt Project 32140015. R.E.M. thanks the grant VRID 216.016.002-1.0. I. A. acknowledges support from Fondo Institucional de Becas FIB-UV and Gemini-Conicyt 32120033. G.D. acknowledges the financial support of the Ministry of Education and Science of the Republic of Serbia through the project 176004 "Stellar physics, also we acknowledge the anonymous referee whose comments helped to improve a first version of this manuscript.



0.5cm
99
	
[Barría et al.(2013)]2013A   A...552A..63B Barría, D., Mennickent, R. E., Schmidtobreick, L., et al. 2013, A&A, 552, A63 

[Barría et al.(2014)]2014A   A...567A.140B Barría, D., Mennickent, R. E., Graczyk, D., & Kołaczkowski, Z. 2014, A&A, 567, A140 
[Bayo et al.(2008)]2008A   A...492..277B Bayo, A., Rodrigo, C., Barrado Y Navascués, D., et al. 2008, A&A, 492, 277 
[Bisikalo et al.(1998)]1998MNRAS.300...39B Bisikalo, D. V., Boyarchuk, A. A., Chechetkin, V. M., Kuznetsov, O. A., & Molteni, D. 1998, mnras, 300, 39 
[Bisikalo et al.(1999)]1999ARep...43..797B Bisikalo, D. V., Boyarchuk, A. A., Chechetkin, V. M., Kuznetsov, O. A., & Molteni, D. 1999, Astronomy Reports, 43, 797 
[Bisikalo et al.(2003)]2003ARep...47..809B Bisikalo, D. V., Boyarchuk, A. A., Kaigorodov, P. V., & Kuznetsov, O. A. 2003, Astronomy Reports, 47, 809 
[Budaj(2011)]2011AJ....141...59B Budaj, J. 2011, aj, 141, 59 
[Cardelli et al.(1989)]1989IAUS..135P...5C Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, Interstellar Dust, 135, 5 
[Castelli & Kurucz(2004)]2004astro.ph..5087C Castelli, F., & Kurucz, R. L. 2004, arXiv:astro-ph/0405087 
[Charbonneau(1995)]1995ApJS..101..309C Charbonneau, P. 1995, ApJS, 101, 309 
[Clausen(2004)]2004NewAR..48..679C Clausen, J. V. 2004, nar, 48, 679 
[Deschamps et al.(2015)]2015A   A...577A..55D Deschamps, R., Braun, K., Jorissen, A., et al. 2015, A&A, 577, A55 
[Djurašević(1996)]1996Ap   SS.243..413D Djurašević, G. 1996, Ap&SS, 243, 413 
[Djurašević(1992)]1992Ap   SS.197...17D Djurašević, G. 1992, Ap&SS, 197, 17 
[de Mink et al.(2014)]2014ApJ...782....7D de Mink, S. E., Sana, H., Langer, N., Izzard, R. G., & Schneider, F. R. N. 2014, apj, 782, 7
[Eggleton(1983)]1983ApJ...268..368E Eggleton, P. P. 1983, apj, 268, 368 
[Eggleton(2006)]2006epbm.book.....E Eggleton, P. 2006, Evolutionary Processes in Binary and Multiple Stars, by Peter Eggleton, pp. . ISBN 0521855578. Cambridge, UK: Cambridge University Press,  2006.
[Ekström et al.(2012)]2012A   A...537A.146E Ekström, S., Georgy, C., Eggenberger, P., et al. 2012, A&A, 537, A146 
[Fitzgerald(1970)]1970A   A.....4..234F Fitzgerald, M. P. 1970, A&A, 4, 234 
[Flower(1996)]1996ApJ...469..355F Flower, P. J. 1996, apj, 469, 355 
[Garrido et al.(2013)]2013MNRAS.428.1594G Garrido, H. E., Mennickent, R. E., Djurašević, G., et al. 2013, mnras, 428, 1594 
[González & Levato(2006)]2006A   A...448..283G González, J. F., & Levato, H. 2006, A&A, 448, 283 
[Hilditch(2001)]2001icbs.book.....H Hilditch, R. W. 2001, An Introduction to Close Binary Stars, by R. W. Hilditch, pp. 392. ISBN 0521241065. Cambridge, UK: Cambridge University Press, March 2001., 392 
[Iben(1967)]1967ARA   A...5..571I Iben, I., Jr. 1967, ARA&A, 5, 571 
[Lucy & Sweeney(1971)]1971AJ.....76..544L Lucy, L. B., & Sweeney, M. A. 1971, AJ, 76, 544 

[Lucy(1967)]1967ZA.....65...89L Lucy, L. B. 1967, zap, 65, 89

[Martin & Whittet(1990)]1990ApJ...357..113M Martin, P. G., & Whittet, D. C. B. 1990, apj, 357, 113 
[Mennickent(2017)]2017SerAJ.194....1M Mennickent, R. E. 2017, Serbian Astronomical Journal, 194, 1 
[Mennickent et al.(2016)]2016MNRAS.461.1674M Mennickent, R. E., Zharikov, S., Cabezas, M., & Djurašević, G. 2016, mnras, 461, 1674 
[Mennickent et al.(2015)]2015ASPC..496..175M Mennickent, R. E., Djurašević, G., Cabezas, M., et al. 2015, Living Together: Planets, Host Stars and Binaries, 496, 175 
[Mennickent et al.(2015)]2015MNRAS.448.1137M Mennickent, R. E., Djurašević, G., Cabezas, M., et al. 2015, mnras, 448, 1137 
[Mennickent et al.(2015)]2015IAUS..307..125M Mennickent, R. E., Djurašević, G., Cabezas, M., et al. 2015, New Windows on Massive Stars, 307, 125 
[Mennickent & Rosales(2014)]2014IBVS.6116....1M Mennickent, R. E., & Rosales, J. 2014, Information Bulletin on Variable Stars, 6116, 1 
[Mennickent et al.(2012)]2012MNRAS.421..862M Mennickent, R. E., Djurašević, G., Kołaczkowski, Z., & Michalska, G. 2012, mnras, 421, 862 
[Mennickent et al.(2012)]2012IAUS..282..317M Mennickent, R. E., Kołaczkowski, Z., Djurašević, G., Diaz, M., & Niemczura, E. 2012, From Interacting Binaries to Exoplanets: Essential Modeling Tools, 282, 317 
[Mennickent et al.(2010)]2010MNRAS.405.1947M Mennickent, R. E., Kołaczkowski, Z., Graczyk, D., & Ojeda, J. 2010, mnras, 405, 1947 
[Mennickent et al.(2008)]2008MNRAS.389.1605M Mennickent, R. E., Kołaczkowski, Z., Michalska, G., et al. 2008, mnras, 389, 1605 
[Mennickent et al.(2003)]2003A   A...399L..47M Mennickent, R. E., Pietrzyński, G., Diaz, M., & Gieren, W. 2003, A&A, 399, L47 
[Pojmanski(1997)]1997AcA....47..467P Pojmanski, G. 1997, AcA, 47, 467 
[Poleski et al.(2010)]2010AcA....60..179P Poleski, R., Soszyński, I., Udalski, A., et al. 2010, AcA, 60, 179 

[Rafert & Twigg(1980)]1980MNRAS.193...79R Rafert, J. B., & Twigg, L. W. 1980, mnras, 193, 79 

[Ruciński(1969)]1969AcA....19..245R Ruciński, S. M. 1969, actaa, 19, 245


[Schleicher & Mennickent(2017)]2017A   A...602A.109S Schleicher, D. R. G., & Mennickent, R. E. 2017, A&A, 602, A109 
Astronomy & Astrophysics, Volume 602, id.A109, 10 pp.

[Schaller et al.(1992)]1992A   AS...96..269S Schaller, G., Schaerer, D., Meynet, G., & Maeder, A. 1992, A&AS, 96, 269 
[Schlafly & Finkbeiner(2011)]2011ApJ...737..103S Schlafly, E. F., & Finkbeiner, D. P. 2011, Apj, 737, 103 
[Schlegel et al.(1998)]1998ApJ...500..525S Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, Apj, 500, 525 
[Torres(2010)]2010AJ....140.1158T Torres, G. 2010, aj, 140, 1158 
[van Rensbergen et al.(2008)]2008A   A...487.1129V van Rensbergen, W., De Greve, J. P., De Loore, C., & Mennekens, N. 2008, A&A, 487, 1129 
[von Zeipel(1924)]1924MNRAS..84..665V von Zeipel, H. 1924, mnras, 84, 665 
[Wright et al.(2010)]2010AJ....140.1868W Wright, E. L., Eisenhardt, P. R. M., Mainzer, A. K., et al. 2010, aj, 140, 1868-1881 
[Zahn(1977)]1977A   A....57..383Z Zahn, J.-P. 1977, A&A, 57, 383 
[Zahn(1975)]1975A   A....41..329Z Zahn, J.-P. 1975, A&A, 41, 329 









































