1. Introduction
Among the big family of close binaries, W Ursae Majoris (W UMa)type contact binaries are the most frequently observed and intensively investigated type of eclipse binaries (Mochnacki Reference Mochnacki1981; Rucinski Reference Rucinski1998, Reference Rucinski2000). A contact binary consists of two components sharing a common convective envelope (Lucy Reference Lucy1968). The shape looks just like a dualcore peanut. Usually, this type of systems have the spectral type of F, G, or K (Rucinski Reference Rucinski1993), and the orbital period mainly distributes between 0.2 and 1.0 d (Paczyński et al. Reference Paczyński, Szczygieł, Pilecki and Pojmański2006). Because of mass and energy exchanges, two components of a contact binary often have almost equal surface temperatures although their masses may be typical unequal. The amplitudes of luminosity variations yielded by eclipse are, in general, less than 0.8 mag in the V band. In order to insight the basic observational features (Rucinski Reference Rucinski1985, Reference Rucinski1993), the properties of the common envelope (Lucy Reference Lucy1973), the physical processes of the mass and energy exchange between two components (Kähler Reference Kähler2002, Reference Kähler2004), and the internal structures of component stars (Webbink Reference Webbink2003; Li, Han, & Zhang Reference Li, Han and Zhang2004), have been also widely investigated over the past half century.
Although contact binaries are the largest part of stellar population, their formation, evolution, and final destinations are still not well understood. In formation process, a triple or multiple system seems to be a necessary stage, because it involves the Kozai mechanism with tidal frictions (Kozai Reference Kozai1962; Kiseleva, Eggleton, & Mikkola Reference Kiseleva, Eggleton and Mikkola1998; Fabrycky & Tremaine Reference Fabrycky and Tremaine2007). However, the statistical results (Tokovinin et al. Reference Tokovinin, Thomas, Sterzik and Udry2006; Eggleton & Tokovinin Reference Eggleton and Tokovinin2008) suggested that not all of contact binaries are indeed the members of triple or multiple star systems. In the contact phase, the W UMatype contact binaries have been further divided into two subtypes: A and Wsubtype (Binnendijk Reference Binnendijk1970). The physical parameters for two different subtypes show different statistical characteristics (Rucinski Reference Rucinski1974; Smith Reference Smith1984; Yakut & Eggleton Reference Yakut and Eggleton2005; Gazeas & Niarchos Reference Gazeas and Niarchos2006; Gazeas & Stępień Reference Gazeas and Stępień2008). Maceroni & van’t Veer (Reference Maceroni and van’t Veer1996) pointed out that these differences are not only related to the lightcurve morphology, but also have an underlying implication to both evolutionary state or origin. Weather do both A and Wsubtype contact binaries have the same progenitor, or not? Can do the one subtype evolve to the other? These problems are still not clear. Finally, both the evolutionary theories (Darwin Reference Darwin1908; Counselman Reference Counselman Charles1973; Webbink Reference Webbink1976; Hut Reference Hut1980), extensive simulations (Han et al. Reference Han, Podsiadlowski, Maxted, Marsh and Ivanova2002; Andronov, Pinsonneault, & Terndrup Reference Andronov, Pinsonneault and Terndrup2006; Li et al. Reference Li, Zhang, Han, Jiang and Jiang2008; Yıldız Reference Yıldız2014), and statistical analyses (Qian et al. Reference Qian, Yang, Zhu, He and Yuan2006; Gazeas & Stępień Reference Gazeas and Stępień2008; Yang & Qian Reference Yang and Qian2015) suggested a contact binary may evolve into the single and rapidly rotating star via merger. However, only one merging event of contact binary, i.e., the red nova event of V1309 Sco (Tylenda et al. Reference Tylenda2011; Zhu et al. Reference Zhu, Zhao and Zhou2016), has been observed so far. These open problems make contact binaries one of the most intriguing objects in stellar astrophysics (Eggleton Reference Eggleton2012), and could, perhaps, be solved by accumulating enough knowledge of the fundamental physical parameters for a large number of such samples.
In this work, we presented the first photometric and orbital period investigations for four W UMatype contact binaries: V473 And, V805 And, LQ Com, and EG CVn. Among them, V473 And (=GSC 0279000068) was discovered by Kuzmin (Reference Kuzmin2008) with the archive data of the Northern Sky Variability Survey (NSVS,Footnote a Woźniak et al. Reference Woźniak2004), where the orbital period of 0.40136 d and the Rband magnitude range from 13.8 to 14.3 were estimated. V805 And (=SvkV38 And, UCAC3 274028768) was discovered by Vrastak (Reference Vrastak2015) using the 0.24 m Newton reflector telescope at the Liptovská Štiavnica Observatory. It is a relatively dark source with the mean magnitude of $14^{\rm m}.716$ and the amplitude of $0^{\rm m}.75$ in V band. The orbital period is 0.44394 d (Vrastak Reference Vrastak2015). LQ Com (=ROTSE1 J123730.26 260451.8, GSC 0199001198) and EG CVn (=ROTSE1 J133726.05 373458.4, GSC 0302601046) were discovered by Akerlof et al. (Reference Akerlof2000) with the firstgeneration Robotic Optical Transient Search Experiment (ROTSE1) telescope. Based on the ROTSE1 survey’s data, Diethelm (Reference Diethelm2001) calculated the phasefolded light curve of LQ Com according to the ROTSE1 period of 0.35684 d, and derived two eclipse timings with the KweeVan Woerden (KW) method (Kwee & van Woerden Reference Kwee and van Woerden1956). EG CVn was later observed by Blattler & Diethelm (Reference Blattler and Diethelm2002) using private 0.15 m Starfire refractor telescope with SBIG ST7 CCD camera, where the orbital period was determined as 0.349271 d and a complete light curve without filter was presented. Together with the ROTSE1 observations, Blattler & Diethelm (Reference Blattler and Diethelm2002) also derived a total of 12 eclipse timings using the KW method.
Besides the photometric observations mentioned above, the four binary systems were scanned by several timedomain surveys, such as the Catalina RealTime Transient Survey (CRTS,Footnote b Drake et al. Reference Drake2009), the Wide Angle Search for Planets (WASP,Footnote c Butters et al. Reference Butters2010), the AllSky Automated Survey for SuperNovae (ASASSN,Footnote d Shappee et al. Reference Shappee2014; Jayasinghe et al. Reference Jayasinghe2019), the Transiting Exoplanet Survey Satellite (TESS, Ricker et al. Reference Ricker2015), and the Zwicky Transient Facility (ZTF,Footnote e Masci et al. Reference Masci2019; Bellm et al. Reference Bellm2019). In particular, based on the photometric data of ASASSN, Jayasinghe et al. (Reference Jayasinghe2018) further corrected their orbital periods using a comprehensive approach involving the Generalised LombeScargle, the MultiHarmonic Analysis of Variance, the Phase Dispersion Minimisation, and the Box Least Squares periodograms. Although these surveys’ photometric observations were discontinuous and of low precision, they distributed over a wide time interval and can provide some phasefolded light curves to estimate the preliminary lightcurve types and determine some eclipse timings. In addition, three targets (i.e., V473 And, LQ Com, and EG CVn) are recently observed in Sector 23 with a 30min cadence by TESS. TESS is a spacebased telescope launched in 2018 April and placed in a highly elliptical orbit with a period of 13.7 d. These three targets were continuously, but longcadence, monitored in 2019 for a cycle (two elliptical orbit periods, 27.4 d), with a break near the midpoint for downlink of data to Earth. These highprecision photometric data will be very helpful for deriving the lightcurve solutions more firmly.
On the side of spectroscopic observation, all of them have been scanned by the Gaia missionFootnote f (Gaia Collaboration et al. Reference Gaia Collaboration2016, Reference Gaia Collaboration2018), and some spectroscopic elements were well revealed. Combining these surveys’ information with our highprecision and multicolour photometric observations, we performed the detailed photometric and orbital period investigations for the four contact binaries, aiming to determine their absolute physical parameters and to uncover their photometric natures and evolutionary states.
2. Photometric observations
The four eclipse binaries were observed using the 60 and 85cm telescopes at Xinglong station of National Astronomical Observatories of China (NAOC). During the observations, the Andor DZ936 PI2048 CCD photometric system was equipped and the standard Johnsoncousins BVR filters were adopted. The observation information (e.g., date, exposure time, and number of images) are summarised in Table 1. All images are reduced to photometric data using the aperture photometry package from the Image Reduction and Analysis Facility (IRAF Footnote g) software library. From their respective fields of view, we selected two stars with the locations, magnitudes, and colours as close to those of the targets as possible, as their comparison and check stars. By checking the magnitude differences between the comparison and check stars, we confirmed that the photometric luminosity between them is stable throughout the observations. From the Set of Identifications, Measurements and Bibliography for Astronomical Data (SIMBAD) database and the latest data release (DR9) of AVVSO Photometric All Sky Survey (Henden et al. Reference Henden, Levine, Terrell and Welch2015), the coordinates and magnitudes in B and V band were collected. From the Gaia DR2, we extracted the parallaxes. All these information for the targets, comparison, and check stars are compiled in Table 2. The reduced data are presented in the form of different magnitude between the target and comparison stars, and the phases are calculated using the orbital periods (the second column of Table 2) determined by Jayasinghe et al. (Reference Jayasinghe2018). The photometric data, including the Heliocentric Julian Day (HJD), phases, and different magnitudes, are compiled as the Supplementary File (see Tables S1–S4 of Appendix A).
The phasefolded light curves for each of our targets are displayed in Figure 1. All of the four binary systems show the typical EWtype luminosity variation, indicating they are in contact or nearcontact configuration. In addition, the light curves of LQ Com show a very wide flatness around both the primary and secondary eclipses, especially for the secondary eclipse (the constant luminosity sustains about 0.08 phase, i.e., about 40 min). This implies that LQ Com is a totaleclipse binary and the radius of its hot primary component is significantly larger than that of the cool secondary one. The slight and symmetric distortion of the primaryeclipse flatness should mainly attribute to the limb darkening effect of the primary component. Combining the wider flatness around the primary eclipse with the fact that the primary eclipse is deeper than the secondary eclipse, we can infer that LQ Com is an Asubtype and totaleclipse binary. For V805 And, a narrow flatness appears on the deeper primary eclipse, but disappear for the somewhat shallow secondary eclipse. This means that V805 And should be also a totaleclipse binary, but belongs to the Wsubtype. The sharp light minima for both V473 And and EG CVn indicate that they should be two partialeclipse binaries.
3. Photometric solutions and absolute parameters
The WilsonDevinney (WD) code (Wilson & Devinney Reference Wilson and Devinney1971; Wilson Reference Wilson1979, Reference Wilson1990; Wilson & Van Hamme Reference Wilson and Van Hamme2014) is one of the most customary programmes to model the observed light and radial velocity curves of various binary systems. It has been revised and improved for many times. Here the latest 2015 version of WD code is adopted to analyse the multiband light curves of the four binaries: V473 And, V805 And, LQ Com, and EG CVn. Firstly, we estimate the effective temperatures of their primary components according to two available colour indices ( $B{}V$ and $J{}K$ ) and the Gaia spectroscopic observations. The $B{}V$ and $J{}K$ colour indices are collected from the AVVSO Photometric All Sky SurveyFootnote h (Henden et al. Reference Henden, Levine, Terrell and Welch2015) and the Two Micron All Sky SurveyFootnote i (Skrutskie et al. Reference Skrutskie2006), respectively. The corresponding temperatures, $T_{\rm BV}$ and $T_{\rm JK}$ , are estimated with the corrected calibrationFootnote j (Pecaut & Mamajek Reference Pecaut and Mamajek2013). The colour indices and the effective temperatures estimated from them and the Gaia spectroscopic observations, are summarised in Table 3. The average values of three temperatures are presented in the last column of Table 3, and the corresponding errors are denoted by the standard deviations. In fact, the average value could not represent exactly the temperature of the primary component, and should be a quadrature temperature of the two components. However, because the temperatures of two components, in general, do not significantly deviate from each other for a typical W UMatype contact binary, we adopted the average temperature as a preliminary one $T_{\rm pre}$ of the primary and the final temperatures for both primary and secondary would be adjusted according to the following equations (Kjurkchieva & Vasileva Reference Kjurkchieva and Vasileva2015),
where $c = l_2/l_1$ , denotes the relative luminosity ratio of two components, and $\Delta T = T_{\rm pre}T_2$ . The gravity darkening and bolometric albedos coefficients of the components were set to be $g_1=g_2=0.32$ (Lucy Reference Lucy1967) and $A_1=A_2=0.5$ (Ruciński Reference Ruciński1969), respectively. For the bolometric and bandpassspecific limb darkening law, we selected the square root law (DiazCordoves & Gimenez Reference DiazCordoves and Gimenez1992). The corresponding coefficients were interpolated from the table of Claret & Bloemen (Reference Claret and Bloemen2011). Of course, these coefficients will be modified when the temperatures are adjusted. The four targets show the typical EWtype light curves, suggesting they should be contact binaries. Thus, we selected the mode 3, i.e., the contact model, in the WD code. Some other parameters, e.g., the mass ratio q, the orbital inclination i, the secondary’s temperature $T_{2}$ , the primary’s monochromatic luminosity $L_{1}$ , and the dimensionless potential $\Omega_{1}=\Omega_{2}$ of two components, are set as the free ones. The secondary’s monochromatic luminosity $L_{2}$ is calculated according to the blackbody radiation model.
Because neither spectroscopic nor photometric mass ratios have been reported for these four binaries so far, we employed the $q$ search method to estimate their initial mass ratios. In the numerical scheme, we calculated Chi square $\chi^{2}=\frac {1} {n}\sum_{i}$ $w_{i}(O_iC_i)^2$ for the optimal fit of light curves at a series of fixed mass ratios q, with the values ranging from 0.20 to 5.0 in step of 0.1. For V473 And and LQ Com, finer grid search at increment of 0.01 for the value of q was conducted near the approximate value to obtain the more accurate values of their initial mass ratios. Figure 2 displayed the relation between Chi square $\chi^{2}$ and mass ratio $q=M_2/M_1$ . Clearly, there are two minima for LQ Com, and one minimum for the other three targets. The first minimum for LQ Com appears at $q=0.26$ , and the second minimum is at $q=4.2$ . As has been mentioned, we could determine that LQ Com is an Asubtype contact binary from the lightcurve morphology. Moreover, $\chi^{2}$ for the first minimum is smaller than that for the second minimum. Thus the mass ratio of LQ Com should be about 0.26. According to the criteria for the reliability of photometric mass ratio proposed by Zhang et al. (Reference Zhang, Qian, Han and Wu2017b), the sharper of the minimum in the $q$ search curve, the more reliable the photometric mass ratio. Both V473 And and LQ Com show a very sharper minimum in their $q$ search curves, but the minima for V805 And and EG CVn are relatively flat. This implies that the photometric mass ratios for V473 And and LQ Com are more reliable than those of V805 And and EG CVn. However, according to the statistical analysis of Pribulla et al. (Reference Pribulla, Kreiner and Tremko2003) and the numerical simulations of Terrell & Wilson (Reference Terrell and Wilson2005), photometric mass ratio should be reliable for totaleclipse contact binaries. Thus, in spite of the flat minimum of the $q$ search curve for V805 And, its mass ratio determined from the photometric light curves, would be still reliable. Comparatively, the mass ratio of EG CVn should be unreliable due to both the partial eclipse and the flat minimum in the qsearch curve.
Based on the above settings of parameters and the initial mass ratios, we run the WD code to simultaneously solve the multiband light curves. After achieving the optimal converging solutions via enough iterations, we modified the temperatures of the primary and secondary components according to Equation (1), and rerun the WD code to obtain the final and selfconsistent converging solutions. The final solutions of these four targets are listed in Table 4 and the corresponding theoretical light curves are displayed in Figure 1 as solid lines. The theoretical light curves could well fit the observation data for V473 And, but slightly deviate from the observational ones around the outofeclipse maxima for V805 And, LQ Com, and EG CVn.
It is well known that many W UMatype contact binaries show asymmetries in two outofeclipse maxima of their light curves, which is known as the O’Connell effect (O’Connell Reference O’Connell1951). Usually, it is very difficult to identify the slight asymmetry (i.e., the weak O’Connell effect) by a visual inspection. In order to justify weather the light curves of our four targets are symmetrical, or not, we employed two measurements to quantify their possible asymmetries in their light curves. The simplest and conventional measurement is the magnitude difference ( $\Delta m_{\rm max}$ ) between two outofeclipse maxima
Another one, called the O’Connell Effect Ratio (OER), was proposed by McCartney (Reference McCartney1997). It is the ratio of the areas beneath the two maxima of a phased light curve. In the numerical scheme, one may first divide the phased light curve into n equally wide bins, where n is chosen to sufficiently sample the light curve. Then, the weighted mean magnitude in each bin is calculated and normalised by subtracting the mean magnitude in the bin including the primary minimum. The OER is finally calculated according to the definition
where $m_{k}$ represents the weighted mean magnitude in the kth bin. Obviously, OER can more thoroughly detect the asymmetries of a light curve than $\Delta m_{\rm max}$ . If $\Delta m_{\rm max}>0$ or $OER>1$ , the O’Connell effect is positive, or it will be negative.
In view of the high precision of TESS photometry, the wideband (approximately 600–1000 nm) and longcadence (30 min) observations of three targets are also used to estimate the possible O’Connell effect and derive photometric solutions. The data product of TESS includes the Single Aperture Photometry (SAP) and the Presearch Data Conditioning light curves. The latter has been processed to make it more suitable for detecting shallow transits of extrasolar planets, and not intended for eclipse binaries which show large variability amplitudes. Thus, we adopted only the SAP data in the present work. The SAP flux data are downloaded from Mikulski Archive for Space Telescopes (MAST) portalFootnote k and converted into magnitude units. We excluded the data points with a nonzero flag for the QUALITY parameter.
Table 5 summarised both $\Delta m$ and OER of the light curve of each band for all four targets. From Table 5, we can conclude that the light curves of V473 And, obtained from both our observations and TESS photometry, are almost perfect symmetrical, while there are slight asymmetries (i.e., weak O’Connell effects) for the other three targets. It should be noted that although both two measurements for the groundbased observations of LQ Com, suggest the existence of positive O’Connell effect, we cannot decide it because the values of $\Delta m_{\rm max}$ are close and even smaller to the corresponding mean errors. But fortunately, the observation dates of TESS for LQ Com were just very close to those of our groundbased observations. Moreover, TESS light curve suggests a significant and the same positive O’Connell effect as our groundbased observations. Thus, we come to the conclusion that there exists a positive O’Connell effect for LQ Com.
${}^{\rm a}$ In this column, GBB, V and R denote the groundbased observation with B, V and R band, respectively. TESSW refers to the wideband observations of TESS.
${}^{\rm b}$ The errors are the mean error of the corresponding photometric data.
In general, the O’Connell effect could be interpreted as a result of the spots on the surface of the component stars. According to the above calculations, the light curves of both V805 And and LQ Com show a positive O’Connell effect, while the light curves of EG CVn presented a weak and negative O’Connell effect. The positive/negative O’Connell effect could be caused by any one of the following four spot models: (1) a hot spot on the primary star (HS1); (2) a cool spot on the primary star (CS1); (3) a hot spot on the secondary star (HS2); (4) a cool spot on the secondary star (CS2). With the four different spot models, we reapplied the WD code to fit the photometric data and obtained the photometric solutions with different spot models, which are compiled in Table 6. Among those spot models, the model with a cool spot on the primary component (CS1) seems to be the optimal one to reproduce the observational light curves for V805 And, LQ Com and EG CVn, because its chisquare value is smallest. However, it should be noted that the differences of $\chi^2$ among four different spot models are rather negligible for the three targets, especially for V805 And and EG CVn.
In order to determine the optimal solutions, we also fitted the TESS light curves with each of those spot models using the WD code. During the analyses of TESS light curves, the photometric solutions without spot derived from our groundbased observations are set as the initial parameters. The limb darkening coefficients for TESS light curves were taken from Claret (Reference Claret2017). The corresponding photometric solutions are summarised in Table 7. Although $\chi^2$ of four different spot models did not also significantly deviate from each other, the optimal lightcurves solutions tended to the CS1. The results are consistent with those derived from our groundbased observations. Thus, we think that the CS1 is the optimal one among the four possible spot models for there three targets. The theoretical light curves corresponding to the optimal photometric solutions for both TESS and our groundbased observations, were calculated and showed as the solid lines in Figure 3.
Because the O’Connell effect is positive for V805 And and LQ Com, but be negative for EG CVn, the cool spots for V805 And and LQ Com could be observed around phase 0.25, while the cool spot for EG CVn could be around phase 0.75 (see the three upper panels of Figure 4). In addition, except from the spot’s parameters, these photometric solutions with different spot models did not significantly deviate from each other. Thus, we would adopt the photometric solutions of CS1 for both V805 And, LQ Com, and EG CVn in the following analysis.
According to the photometric solutions, the orbital inclinations for V805 And and LQ Com are larger than $80^{\circ}$ , confirming the totaleclipse inference from their lightcurves morphology. For V473 And and EG CVn, the photometric solutions suggest that they are the partialeclipse binaries, but very close to the total eclipse. The total or partialeclipse characteristics for the four targets could be checked visually from the geometrical configurations shown in the bottom panels of Figure 4. In addition, we noted that in the photometric solutions with different spot models, the primary components of both V805 And and EG CVn have lower masses and higher temperatures than their secondary components, implying that they are two Wsubtype contact binaries.
Although the radialvelocity curves for the four binaries are absent so far, we can still estimate their absolute physical parameters according to the following two different methods: (1) the scheme based on the Gaia distance (Kjurkchieva et al. Reference Kjurkchieva, Popov, Eneva and Petrov2019), and (2) the empirical relations (Qian Reference Qian2003; Gazeas Reference Gazeas2009; Yu et al. Reference Yu, Li, Huang, Hu and Xiang2022). The standard procedure for the former was proposed by Kjurkchieva et al. (Reference Kjurkchieva, Popov, Eneva and Petrov2019) and developed by Liu et al. (Reference Liu, Qian, Li, He, Li, Zhao and Li2020) and Li et al. (Reference Li2021). In our calculations, the Vband magnitudes $V_{\rm ASASSN}$ at maximum brightness were determined from the phasefolded light curves of ASASSN observations, the interstellar extinctions $A_{\rm V}$ were derived with the S&F method (Schlafly & Finkbeiner Reference Schlafly and Finkbeiner2011) from the IRAS database.Footnote l The bolometric corrections $BC_{\rm V}$ were interpolated from the calibration proposed by Pecaut & Mamajek (Reference Pecaut and Mamajek2013). These parameters are summarised in Table 8. For the later, we adopted the update empirical relation between the semimajor axis and orbital period derived by Yu et al. (Reference Yu, Li, Huang, Hu and Xiang2022),
The values in parentheses represent the standard error which are calculated according to the error propagation rule.
Combining our photometric solutions with the semimajor axis determined from the two distinct methods, we calculated the absolute parameters of the four targets which are summarised in Table 9. For V473 And, LQ Com, and EG CVn, two sets of parameters calculated from two distinct methods are almost equal within the error range. However for V805 And, the absolute parameters calculated from the Gaia distance are significantly deviated from those determined from the empirical relation. As has been pointed out by Li et al. (Reference Li2021), the absolute parameters for most of the contact binaries (more than 80% of contact binaries) can be well estimated from the Gaia distance. But for a friction of contact binaries, the absolute parameters determined from the Gaia distance may be unreliable and even unrealistic due to many reasons, such as too far away, the thirdlight effect, inaccurate interstellar extinction, and inaccurate Vband magnitude at maximum brightness. Therefore, in the following analysis, we adopted the weighted mean parameters for V473 And, LQ Com, and EG CVn, and the parameters determined by the empirical relation for V805 And.
The values in parentheses represent the standard error which are calculated according to the error propagation rule.
4. Orbital period variations and their physical origins
4.1. Orbital period variations
The orbital period variations of close binaries have an underlying implication to various astrophysical processes and dynamical evolution of binary systems. In order to uncover the period variation of the four targets: V473 And, V805 And, LQ Com, and EG CVn, we performed an extensive search for all available eclipse timings, and obtained 107 data points from the literature and the eclipse timings database (i.e., the OC gatewayFootnote m). From our observations, we calculated 16 eclipse timings with the KW method (Kwee & van Woerden Reference Kwee and van Woerden1956). In addition, from several surveys (e.g., NSVS, WASP, and ASASSN, ZTF, TESS, etc.), we extracted the photometric data of the four targets and calculated the phasefolded light curves. With the method of Borkovits et al. (Reference Borkovits, Rappaport, Hajdu and Sztakovics2015), we derived 114 eclipse timings. Finally, a total of 237 eclipse timings, including those collected from the literature and calculated from our observations and surveys’ ones, are compiled in Table 10.
References: 1. This paper; 2. Kuzmin (Reference Kuzmin2008); 3. Banfi et al. (Reference Banfi2012); 4. Corfini et al. (Reference Corfini2014); 5. Diethelm (Reference Diethelm2013); 6. Hubscher (Reference Hubscher2014); 7. (Vrastak Reference Vrastak2015); 8. Yao et al. (Reference Yao2015); 9. GCVS; 10. Paschke (2019); 11. Diethelm (Reference Diethelm2001); 12. Paczyński et al. (Reference Paczyński, Szczygieł, Pilecki and Pojmański2006); 13. Nelson (Reference Nelson2008); 14. Brat et al. (Reference Brat2011); 15. Hubscher et al. (Reference Hubscher, Lehmann, Monninger, Steinbach and Walter2010); 16. Hubscher & Monninger (Reference Hubscher and Monninger2011); 17. Hoňková et al. (Reference Hoňková2013); 18. Hubscher, Lehmann, & Walter (Reference Hubscher, Lehmann and Walter2012); 19. Hubscher et al. (Reference Hubscher, Braune and Lehmann2013); 20. Honková et al. (Reference Honková2014); 21. Honkova et al. (Reference Honkova2015); 22. Hubscher (Reference Hubscher2016); 23. Juryšek et al. (Reference Juryšek2017); 24. Hubscher (Reference Hubscher2017); 25. Pagel (Reference Pagel2018); 26. Pagel (Reference Pagel2020); 27. Blattler & Diethelm (Reference Blattler and Diethelm2002); 28. Nelson (Reference Nelson2002); 29. Diethelm (Reference Diethelm2003); 30. Diethelm (Reference Diethelm2004); 31. Diethelm (Reference Diethelm2005); 32. Diethelm (Reference Diethelm2006); 33. Diethelm (Reference Diethelm2007); 34. Nelson (Reference Nelson2009); 35. Diethelm (Reference Diethelm2009); 36. Diethelm (Reference Diethelm2011); 37. Nelson (Reference Nelson2019).
${}^{\rm a}$ These times were derived from the photometric data of NSVS.
${}^{\rm b}$ These times were derived from the photometric data of WASP.
${}^{\rm c}$ These times were derived from the photometric data of ASASSN.
${}^{\rm d}$ These times were derived from our observations.
${}^{\rm e}$ These times were derived from the photometric data of ZTF.
${}^{\rm f}$ These times were derived from the photometric data of TESS.
${}^{\rm g}$ These times were derived from the photometric data of BRNO.
${}^{\rm h}$ These times were derived from the photometric data of ASAS DR3.
${}^{\rm i}$ These times were derived from the photometric data of AAVSO.
In order to build the $OC$ diagrams, we constructed the following linear ephemerides
where $T_{0}$ is the primary eclipse timing derived from our observations and $P_{0}$ is the orbital period determined by Jayasinghe et al. (Reference Jayasinghe2018). With the linear ephemerides, we calculated the $OC$ values for their eclipse timings which are displayed in Figure 5. Clearly, the $OC$ diagrams for V473 And, V805 And, and EG CVn, exhibit a complex variation. However, the variation of the $OC$ curve for LQ Com seems to be not observable due to the large scatters. In order to uncover the secular period changes, we first adopt the quadratic function, i.e.,
to fit their $OC$ values. Among the four targets, LQ Com should be not undergoing a secular period variation because the coefficient of its quadratic term is 3 orders of magnitude smaller than the typical values for contact binaries (Hu et al. Reference Hu, Jiang, Yu and Xiang2018). Thus, for LQ Com, we adopted a linear fit. After subtracting the linear or quadratic fit, the residuals $(OC)_{1}$ still exhibit the obvious periodic variation. By introducing the sinusoidal term into the quadratic function, i.e.,
we can achieve the sufficiently good fits.
By assigning the reciprocals of the normalised square errors as the weights and using the weighted leastsquares method, we derived the fitting parameters which were summarised in Table 11. As should be mentioned, although the errors for two eclipse timings of V805 And are absent, we performed a weighted fit by cutting the two data points. The quadratic term implied that V473 and, V805 And and EG CVn are undergoing a secular period change. The sinusoidal term revealed a cyclic EclipseTiming Variation (ETV) for all four targets. Based on the fitting parameters, we calculated the rates of secular period changes and the modulation periods of cyclic ETVs (see the third and fourth lines from the bottom of Table 11).
4.2. Physical origins of orbital period variations
4.2.1. Secular period changes
The above orbital period investigations suggested that V473 And, V805 And, and EG CVn shows the secular orbital period decrease and increase, respectively. The secular orbital period variations could be, in general, interpreted as a result of the mass transfer and/or the mass and angular momentum loss. According to the photometric solutions, these three binaries are not in the relatively deep contact configuration ( $f<50\%$ ), i.e., their surface potentials are significantly lower than those of their outer Lagrange points. Thus, it should be difficult and even impossible to loss their mass via Rochelobe overflow from the outer Lagrange point. In addition, the typical massloss rate due to stellar wind for an active dwarf star is about $10^{10}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ (Mullan et al. Reference Mullan, Doyle, Redman and Mathioudakis1992) or even lower (Lim & White Reference Lim and White1996). It contributes the decrease rate of orbital period at the level of ${\sim}10^{10}$ ${\rm d}\cdot {\rm yr}^{1}$ , which is two or three orders of magnitude smaller than the observed ones (typical observed value is $10^{7}$ ${\rm d}\cdot {\rm yr}^{1}$ , Hu et al. Reference Hu, Jiang, Yu and Xiang2018). Thus, the effect of mass loss on orbital period decrease may be neglected for V473 And, V805 And, and EG CVn.
When considering the mass conservation and the circular orbit of a contact binary, its orbit evolution can be dominated by
where $J=\sqrt[3]{\frac{G^{2}M_{1}M_{2}P}{2\pi (M_{1}+M_{2})}}$ is the orbital angular momentum. Equation (6) suggests explicitly that both the conservative mass transfer and the angular momentum loss could yield the secular period changes.
In regard of the angular momentum loss, the most plausible mechanism is the magnetic breaking due to stellar wind. From the standard magnetic breaking model (Weber & Davis Reference Weber and Davis1967), Guinan & Bradstreet (Reference Guinan and Bradstreet1988) derived the period decrease rate caused by the angular momentum loss via magnetic stellar wind
In the formula, the gyration constant $k^{2}$ is roughly 0.1 estimated from the main sequence stars. By putting the absolute physical parameters into Equation (7), we estimated the orbital period decrease rates of $1.59\times 10^{7}$ ${\rm d}\cdot {\rm yr}^{1}$ , $7.92\times 10^{8}$ ${\rm d}\cdot {\rm yr}^{1}$ and $1.18\times 10^{7}$ ${\rm d}\cdot {\rm yr}^{1}$ for V473 And, V805 And and EG CVn, respectively. For V473 And and V805 And, the orbital period decrease rates due to the angular momentum loss are smaller than the observed values, implying the angular momentum loss via magnetic stellar wind could not fully yield the observed orbital period decrease. For EG CVn, the orbital period is undergoing a secular increase at the rate of $2.69\times 10^{7}$ ${\rm d}\cdot {\rm yr}^{1}$ , but the angular momentum loss due to magnetic stellar wind can only lead to secular orbital period decrease. Thus, in order to yield the observed orbital period increase/decrease for V473 And, V805 And, and EG CVn, the conservative mass transfer seems to be required.
Combining Equation (8) with (9), we may write the masstransfer rate as
where $\dot P_{\rm obs}\dot P_{\rm aml}$ denotes the secular orbital period variation yielded by the conservative mass transfer. Thus, if the secular orbital period changes of V805 And and EG CVn are yielded by a combination of the angular momentum loss and the mass transfer, the conservative mass transfer rates would be $\dot M_{2}=\dot M_{1}=5.89\times 10^{9}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ for V473 And, $\dot M_{1}=\dot M_{2}=2.58\times 10^{7}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ for V805 And, and $\dot M_{2}=\dot M_{1}=1.73\times 10^{7}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ for EG CVn. If the secular orbital period changes are caused separately by the conservative mass transfer (i.e., $\dot P_{\rm aml}=0$ ), the required mass transfer rates are $\dot M_{2}=6.86\times 10^{8}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ for V473 And, $\dot M_{1}=3.11\times 10^{7}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ for V805 And, and $\dot M_{2}=1.20\times 10^{7}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ for EG CVn. We may estimate the ranges of masstransfer rate as $[0.59\sim 6.86]\times 10^{8}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ , $[2.58\sim 3.11]\times 10^{7}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ and $[1.20\sim 1.73]\times 10^{7}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ for V805 And and EG CVn, respectively.
According to the thermal timescale of the donor star ( $\tau_{\rm th}\sim \frac {GM_{\rm donor}^{2}} {R_{\rm donor}L_{\rm donor}}$ , Paczyński Reference Paczyński1971) of V473 And, V805 And and EG CVn, we estimated the mass transfer rates on the thermal timescale $\frac {M_{donor}} {\tau_{\rm th}}=1.05\times 10^{8}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ , $5.23\times 10^{8}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ and $4.89\times 10^{8}$ ${\rm M}_{\odot}\cdot {\rm yr}^{1}$ . For V473 and, the mass transfer rate on the thermal timescale is significantly larger than the upper limit of the mass transfer rate estimated from Equation (10), implying that its secular period decrease should be involved with some additional mechanisms, e.g., the mass/angular momentum loss and unstable mass transfer. For V805 And and EG CVn, the thermaltimescale mass transfer rates are significantly smaller than the lower limits of the mass transfer rates determined from Equation (10), implying that the conservative mass transfer could yield the observed period decrease and increase.
Theoretically, the continuous mass transfer could yield the impacting spots on the surface of the mass gaining component (Flannery Reference Flannery1975; Zhou & Leung Reference Zhou and Leung1990; Zhai & Fang Reference Zhai and Fang1995; Zhu et al. Reference Zhu, Qian, Zola and Kreiner2009; Hu et al. Reference Hu, Chen, Xiang, Yu and Zhao2019). For V805 And, the more massive secondary is slightly cooler than the primary. Moreover, due to the contact configuration, the surface potentials of two components are equal so that there is not accreting energy to heat the transferred material. Thus, the spot yielded by mass transfer should be a cool one. Under the effect of Coriolis force, this cool spot can move to a small longitude and be observed around phase 0.75. These theoretical predictions seem to be just consistent with the best photometric solutions (CS1), indicating that the cool spot might be indeed yielded by the mass transfer. However, for EG CVn, the optimal photometric solution (CS1) cannot match the theoretical prediction of impacting spot, which means that the cool spot of EG CVn may be originated from the magnetic activity of its primary component.
4.2.2. Cyclic period changes
The above analyses of $OC$ curves suggested that all four targets: V473 And, V805 And, LQ Com, and EG CVn, show periodic ETVs. The periodic ETVs are, generally, regarded as the cyclic period variation of an eclipsing binary. However, for some cases, they are not the embodiment of the cyclic period variations. This is dependent on the causes yielding the cyclic ETVs. So far, there are at least three different causes of cyclic ETVs: (1) the apsidal motion; (2) the lighttime effect due to the third body; (3) the cyclic magnetic activity. For the cyclic ETVs of the three targets, we can exclude the apsidal motion because the primary and secondary eclipse timings do not follow a reversedphase variation. Thus, we would focus on the lighttime effect due to a third body and the cyclic magnetic activity.
4.2.2.1. Lighttime effect due to third body
If the cyclic ETVs of V473 And, V805 And, LQ Com, and EG CVn, are caused by the lighttime effect of the third body (Irwin Reference Irwin1952), the orbits of their third body should be or at least close to circular one because of the good sinusoidal fits for the $OC$ curves. By inserting the modulation period $P_{\rm mod}$ and amplitude A of cyclic ETVs into the following equation
where $a_{12}\sin i'=A\times c$ ( $a_{12}$ and i ^{′} denote the orbital radius and inclination of eclipse binary pair or third body around the barycentre of triple system, and c is light velocity), we may determine the mass function $f(M_{3})$ of the third body. Based on the equation
we could then obtain the relation between the mass $M_{3}$ and orbital inclination i ^{′} of the third body. Figure 6 shows the relations between the mass and orbital inclination of the third bodies for these four targets. If these third bodies move just in a coplanar orbit around their corresponding host binaries, their masses are equal to $0.184$ , $0.336$ , $0.066$ and 307 ${\rm M}_{\odot}$ , respectively. When $i'=90^{\circ}$ , the third bodies of four binary systems would reach their minimum masses, i.e., $0.178$ , $0.334$ , $0.066$ and $0.298$ ${\rm M}_{\odot}$ , respectively. For LQ Com, the minimum mass is about 70 times of mass of Jupiter. If there really exists the third body orbiting LQ Com, it should be a brown dwarf. For V473 And, V805 And, and EG CVn, the minimum masses of third bodies are comparable to those of the less massive components of their hosting binaries. Their third bodies might be a stellar companion. Thus, it can be expected to detect the spectra and light contributions of the third bodies. Unfortunately, the spectroscopic observations for these targets are absent so far. For the light contributions of the third bodies, our photometric studies did not suggest any significant third lights. Of course, it should be very difficult, or even unrealistic to detect the third light from light curves.
The most direct way to confirm the existence of the third body orbiting a binary system is to detect the transiting circumbinary events (Carter et al. Reference Carter2011; Zhang, Qian, & He Reference Zhang, Qian and He2017a; Socia et al. Reference Socia2020) of the third body from the photometric light curves with high resolution, or the radialvelocity change of binary system (i.e., the socalled $\gamma$ velocity variation, Van Hamme & Wilson Reference Van Hamme and Wilson2007). However, the transiting circumbinary events are very rare because of both strict limitation to the orbital geometry and the low frequency of occurrence. For the $\gamma$ velocity variation, as has been mentioned, the spectroscopic observations for these targets are still absent. In addition, the gravitational interaction of the third body may typically result in precession of orbital plane of binary system, which could be observable as a change of eclipsing depth on a long timescale. But for our targets, both the lowprecision observations from early surveys and the disturbances of magnetic activity spots did not allow us to precisely detect any significant eclipsedepth variations on a long timescale. Of course, besides the Røemer delay (the lighttime effect), the third body may also yield the dynamical delay on eclipse timings (Soderhjelm Reference Soderhjelm1975; Borkovits et al. Reference Borkovits, Érdi, ForgácsDajka and Kovács2003, Reference Borkovits, Csizmadia, ForgácsDajka and Hegedüs2011, Reference Borkovits, Rappaport, Hajdu and Sztakovics2015, Reference Borkovits, Hajdu, Sztakovics, Rappaport, Levine, Bró and Klagyivik2016). Because of the short orbital periods ( $P<1$ d) of the targeting binaries and the relatively long modulation periods ( $P_{\rm mod}\sim 10$ yr) of the third bodies, the classical Røemer delay will dominate the ETVs and the dynamical delay may be negligible (Rappaport et al. Reference Rappaport, Deck, Levine, Borkovits, Carter, El Mellah, SanchisOjeda and Kalomeni2013). In view of these facts, the followup spectroscopic and photometric observations for these binaries should be performed to confirm the existence of their third bodies.
4.2.2.2. Cyclic magnetic activity
Regarding the explanation of the cyclic magnetic activity, Applegate (Reference Applegate1992) proposed a theoretical model to explain the cyclic period oscillations for contact binaries. It has been developed lately by Lanza, Rodono, & Rosner (Reference Lanza, Rodono and Rosner1998) and Lanza & Rodonò (Reference Lanza and Rodonò2002). According to this theory, strong magnetic activities below the surface of one or both components play major roles in producing the cyclic period variations. When the magnetic field distribution of an active component changes during its magnetic cycle, the distribution of angular momentum will be also adjusted. This sequentially yields variations of the quadrupole moment, resulting in variations of the star’s oblateness and radial differential rotation. Variations in shape of the star would be gravitationally coupled to the orbit, modulating the orbital period. Meanwhile, the luminosity of the active component would be variable with variations of its shape.
Using the modulation period and amplitude of cyclic period variation, we may firstly calculate the amplitude $\Delta P$ and relative amplitude $\Delta P/P$ of orbital period modulation with the formula
The other model parameters, i.e., the change of quadruple moment $\Delta Q$ , the transferred angular momentum $\Delta J$ , the relative variation of angular velocity $\Delta \Omega/\Omega$ , the required energy $\Delta E$ , the relative luminosity change $\Delta L/L$ , and the required strength of magnetic field B, can be obtained from the following formulas derived by Applegate (Reference Applegate1992)
In these formulas, G is the gravitational constant, $M_{\rm shell}$ and $I_{\rm shell}$ denote the mass and the moment of inertia of the shell for the active component, respectively. M, R, and L are the mass, radius, and luminosity of the active component.
By insetting the absolute physical parameters and modulation period of three targets into the above formulas, we calculated the model parameters needed to yield the observed cyclic period variations for the three targets, which are listed in Table 12. In the calculations, the mass of convective shell was assumed to be typically equal to onetenth of the mass of the corresponding active component, i.e., $M_{\rm shell}= 0.1M_{\rm active}$ . Also, our calculations were performed for both the primary and secondary components because both two components of these binaries can, in principle, satisfy the condition to yield strong magnetic field. In addition, the typical values of those model parameters required by Applegate’s mechanism are presented in the penultimate column of Table 12 for a reference. Clearly, our four binaries are observed to undergo orbital period modulation with amplitude $\Delta P/P\sim 10^{6}$ , over time scales of $P_{\rm mod}\sim 10$ yr. These modulations can be yielded by the gravitational coupling of the orbit to variations in the shape of a magnetically active component in these systems. In order to yield the variable deformation of their active stars, the distribution of angular momentum needed to be variable at a level of $\Delta J\sim 10^{47}$ g $\cdot{\rm cm}^{2}\cdot{\rm s}^{1}$ , to produce the variation of the quadrupole moment at a level of $\Delta Q\sim 10^{49}$ g $\cdot{\rm cm}^{2}$ . Meanwhile, due to the variation of shape of the active component, the relative luminosity variation should be at the $\Delta L/L\simeq 0.01$ level, and the active component would be differentially rotating at the $\Delta \Omega/\Omega\simeq 0.01$ level. The torque needed to redistribute the angular momentum should be exerted by a mean subsurface magnetic field of several kilogauss. Obviously, the variations of those model parameters and the mean subsurface magnetic field, well followed the typical values required by Applegate’s mechanism. These results indicated that the cyclic period changes for all four targets may be interpreted as a result of the cyclic magnetic activity.
For the energy required by Applegate’s model, Brinkworth et al. (Reference Brinkworth, Marsh, Dhillon and Knigge2006) proposed another quantised criterion, i.e., whether the active component can provide enough energy to support Applegate’s mechanism, or not. According to the method of Brinkworth et al. (Reference Brinkworth, Marsh, Dhillon and Knigge2006), we calculated the available energiesFootnote n generated by the components for the four targets. Because the energy required by Applegate’s mechanism is very sensitive to the assumed shell mass, the required energies with various different shell masses are calculated and showed in Figure 7. From Figure 7, it can be seen that except for the less massive primary components of V805 And, the available energies provided by the components for all four targets may significantly exceed the required ones for a certain range of shell mass, i.e., these components can generate enough energy to drive Applegate’s mechanism. In another word, the observed cyclic period variations for the four targets may be caused by the cyclic magnetic activity of at least one component of these binaries. Of course, some observable evidences should be further investigated to confirm the explanation of cyclic magnetic activity. For instance, the luminosity variation and any other indicator of magnetic activity (starspot activity, coronal Xray luminosity, emission cores in CaII or MgII lines, etc.) should also show the same period as the orbital period modulation (Kim et al. Reference Kim, Jeong, Demircan, Muyesseroglu and Budding1997; Hu et al. Reference Hu, Yu, Zhang and Xiang2020). However, these evidences require a longterm (the timescale of several decades, or at least one period of cyclic period variation) photometric and/or spectroscopic observations.
5. Evolution of contact binaries
Based on the absolute physical parameters, we located the four binaries at the massradius and massluminosity diagrams (Figure 8) to discuss their evolutionary states. For comparison, 179 contact binaries collected by Yu et al. (Reference Yu, Li, Huang, Hu and Xiang2022) were also added in Figure 8. In the diagrams, the ZeroAge Main Sequence (ZAMS) and TerminalAge Main Sequence (TAMS) lines were calculated based on the PARSEC modelsFootnote o (Bressan et al. Reference Bressan, Marigo, Girardi, Salasnich, Dal Cero, Rubele and Nanni2012). All the contact binaries are classified into two subtypes: A and W subtype. Similar to other contact binaries, the more massive components of our four targets are mainly located on the main sequence band and fairly close to the ZAMS line with relatively high metallicity. The less massive components exhibit the oversized and overluminous characteristics relative to the single star with the same mass. They have radii 34 times larger than expected for the corresponding ZAMS masses. Both Lucy (Reference Lucy1968) and Moses (Reference Moses1976) suggested the oversize and overluminous phenomena for the less massive components are caused by the energy transfer from the more massive component to the less massive one. However, it seems physically hard to inflate a ZAMS star to such size purely by the energy transfer from its more massive companion (Stępień Reference Stępień2004). So, the less massive components should be more evolved with hydrogen depleted in their centres.
Although both A and Wsubtype systems belong to the same family of lowmass contact binaries, the evolutionary relation between two subtypes is still a matter of controversy and debate. Following the original idea of Lucy (Reference Lucy1976), Hilditch King & McFarlane (Reference Hilditch, King and McFarlane1988) and Hilditch (Reference Hilditch1989) suggested that the Asubtype contact binaries are more evolved than the Wsubtype systems from the massradius and colourluminosity diagrams. However, by further analysing the massluminosity and periodangular momentum diagrams, both Maceroni & van’t Veer (Reference Maceroni and van’t Veer1996) and Yakut & Eggleton (Reference Yakut and Eggleton2005) pointed out that the conclusion of Hilditch et al. (Reference Hilditch, King and McFarlane1988) might be incorrect. Moreover, from a statistical standpoint, because the total mass and angular momentum of Asubtype systems are significantly larger than those of Wsubtype systems, the Asubtype contact binaries should be less evolved than the Wsubtype (Gazeas & Niarchos Reference Gazeas and Niarchos2006).
Combining the catalog of Yu et al. (Reference Yu, Li, Huang, Hu and Xiang2022) with our four targets, we here systematically calculated the statistical averages of physical parameters for two subtypes of contact binaries (Table 13). Clearly, the Wsubtype contact binaries have shorter orbital periods and lower temperatures than the Asubtype systems. The two statistical features are similar with those of Rucinski (Reference Rucinski1974) and Smith (Reference Smith1984), indicating that the Wsubtype systems are more evolved than Asubtype ones. Moreover, similar to the statistical results of Gazeas & Niarchos (Reference Gazeas and Niarchos2006), our calculations also show that both the total mass and orbital angular momentum of Wsubtype systems are lower than those of Asubtype ones. In addition, we found that the temperature difference between two components of a Wsubtype system is significantly smaller than that of an Asubtype one, implying that Wsubtype systems should be closer to the thermal equilibrium state than Asubtype systems. These statistical evidences show formation of Wsubtype systems from Asubtype systems is possible, but the opposite direction is not strongly supported. Usually, the evolution from Asubtype to Wsubtype systems may be accompanied by a simultaneous mass and angular momentum loss, as well as mass and energy exchanges between two components of a contact binary. Thus, the opposite evolutionary direction seems to be impossible, because it requires an increase of both total mass, orbital angular momentum, and temperature difference.
The units of orbital period and temperature are day and Kelvin, respectively. The orbital angular momentum is in a unit of $\times 10^{51}$ ${\rm cm}^{2}\cdot {\rm g}\cdot {\rm s}^{1}$ . The other parameters are in the solar unit. The values in the third and fifth lines denote the standard derivations of the corresponding parameters.
An additional evidence for the evolutionary direction is that binary systems with shorter orbital periods are, in general, kinematically older (age 8 Gyr) than those with longer periods (age 2 Gyr) (Bilir et al. Reference Bilir, Karataş, Demircan and Eker2005). Of course, there is an alternative evolutionary path for the two different subtypes. For instance, the close detached binaries with different initial states may evolve independently into either A or Wsubtype contact binaries (Zhang, Qian, & Liao Reference Zhang, Qian and Liao2020). Also, the close detached binaries evolve in a similar way but with different mass/energy transfer rate and/or different mass/angular momentum loss rate, leading to different evolutionary outcomes (Gazeas & Stępień Reference Gazeas and Stępień2008). Moreover, we cannot exclude the possibility that some Asubtype systems with small angular momentum have evolved from Wsubtype systems, while others (with relatively large angular momentum) have evolved directly from nearcontact binaries (Yakut & Eggleton Reference Yakut and Eggleton2005).
6. Conclusion
In this paper, we performed the first photometric and orbital period investigations for four W UMatype contact binaries: V473 And, V805 And, LQ Com, and EG CVn. Also, we analysed their current states of evolution and made a simple statistical study based on the physical parameters of nearly 200 contact binaries. From these investigations, we could draw out several following conclusions:

(1) V805 And and LQ Com are the total eclipsing binaries, while V473 And and EG CVn is the partial eclipsing one. V473 And and LQ Com belong to the A subtype of W UMatype contact binaries, while V805 And and EG CVn belong to the Wsubtype.

(2) The light curves of V805 And, LQ Com, and EG CVn show a weak O’Connell effect which may be interpreted as a result of the cool spot on their primary components.

(3) With two distinct methods, the absolute parameters for four targets have been well estimated.

(4) All four targets show cyclic period variations which may be caused by either the lighttime effect of the third body or the cyclic magnetic activity. In addition, due to mass transfer, three systems: V473 And, V805 And, and EG CVn, are undergoing a secular period decrease and increase, respectively.

(5) A statistical study suggested that the Wsubtype contact binaries should be more evolved than the Asubtype ones. Furthermore, it is possible that some Wsubtype systems have been evolved from the Asubtype systems, but the opposite direction seems to be physically difficult.
Finally, it should be mentioned that two conclusions: the W subtype of EG CVn and the possible evolution direction of Asubtype contact binaries to Wsubtype systems, are rather dubious. Although our photometric solutions for EG CVn derived from both TESS and our groundbased observations, suggested that EG CVn is a Wsubtype contact binary, we should note that the difference between two eclipse depths of its light curves is very small, even smaller than the error of photometric data, indicating an almost equal temperatures of the primary and secondary components. Thus, it would be not very reliable that the primary star is somewhat cooler than the secondary star. Moreover, according to the statistical conclusion of Qian (Reference Qian2001), the Wsubtype contact binaries with a high mass ratioFootnote p ( $q'>0.4$ ) usually show a secular period increase, while the orbital periods of the lowmassratio ( $q'<0.4$ ) systems are undergoing a secular decrease. Among our two Wsubtype systems, V805 And seems follow to the conclusion, while EG CVn breaks this rule because it has a low mass ratio ( $q'=1/q=0.27$ ), and shows a secular period increase. However, another statistical analysis for secular variations in the orbital periods of 73 contact binaries suggested that these systems were divided roughly equally between the secular increase and decrease of their orbital periods, regardless of the mass ratio (Dryomova & Svechnikov Reference Dryomova and Svechnikov2006). This result does not support the conclusion of Qian (Reference Qian2001). Moreover, the equal division plausibly indicated mass might be transferred alternately from one component to the other, just as predicted by the thermal relaxation oscillation theory (Lucy Reference Lucy1976; Flannery Reference Flannery1976). For the latter, the statistical differences of physical parameters between two subtypes, perhaps, are originated from the insufficient samples of two subtypes. In a word, the two conclusions should be open to question and needed to be further studied observationally and theoretically.
Acknowledgments
This work is supported by the Joint Research Funds in Astronomy (U1931115, U2031114 and U1731110) under cooperative agreement between the National Natural Science Foundation of China and the Chinese Academy of Sciences. We acknowledge the support of the staff of the Xinglong 60cm and 85 telescopes. This work was partially supported by the Open Project Program of the Key Laboratory of Optical Astronomy, National Astronomical Observatories, Chinese Academy of Sciences.
We acknowledge the following surveys for their valuable data: (1) European Space Agency (ESA) mission Gaia Footnote q (Gaia Collaboration et al. Reference Gaia Collaboration2016, Reference Gaia Collaboration2018), processed by the Gaia Data Processing and Analysis Consortium (DPACFootnote r). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement; (2) the Northern Sky Variability Survey (Woźniak et al. Reference Woźniak2004), whose data were collected by the first generation Robotic Optical Transient Search Experiment; (3) the Wide Angle Search for Planets, where the data we used is from the DR1 of the WASP data (Butters et al. Reference Butters2010) as provided by the WASP consortium, and the computing and storage facilities at the CERIT Scientific Cloud, reg. no. CZ.1.05/3.2.00/08.0144 which is operated by Masaryk University, Czech Republic; (4) the AllSky Automated Survey for SuperNovae (Shappee et al. Reference Shappee2014; Jayasinghe et al. Reference Jayasinghe2019), which consists of two telescopes on a common mount which is hosted by Las Cumbres Observatory Global Telescope Network in the Faulkes Telescope North enclosure on Mount Haleakala, Hawaii; (5) the AAVSO Photometric All Sky Survey (Henden et al. Reference Henden, Levine, Terrell and Welch2015), which was funded by the Robert Martin Ayers Sciences Fund to perform a widefield, allsky photometric survey; (6) the Brno Regional Network of Observers project (BRNO, Hajek Reference Hajek2006; Zejda Reference Zejda2006), a project administered by the Variable Star and Exoplanet Section of Czech Astronomical Society; (7) the Zwicky Transient Facility (ZTF, Masci et al. Reference Masci2019; Bellm et al. Reference Bellm2019) a publicprivate partnership aimed at a systematic study of the optical night sky. ZTF is funded in equal part by the US National Science Foundation and an international consortium of universities and institutions; (8) the Transiting Exoplanet Survey Satellite (TESS, Ricker et al. Reference Ricker2015) mission. Funding for the TESS mission is provided by the NASA’s Science Mission Directorate; and (9) the Two Micron All Sky Survey (Skrutskie et al. Reference Skrutskie2006), a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation.
Authors are grateful to two anonymous referees for constructive comments and suggestions, which have proved to be very helpful for improving the manuscript.
Supplementary material
To view supplementary material for this article, please visit http://doi.org/10.1017/pasa.2022.53.