# A Study of Twelve Potential Merger Candidate Contact Binary Systems

Surjit S. Wadhwa School of Science, Western Sydney University,
Locked Bag 1797, Penrith, NSW 2751, Australia.
Bojan Arbutina Department of Astronomy, Faculty of Mathematics, University of
Belgrade, Studentski trg 16, 11000 Belgrade, Serbia.
Nick F. H. Tothill School of Science, Western Sydney University,
Locked Bag 1797, Penrith, NSW 2751, Australia.
Miroslav D. Filipović School of Science, Western Sydney University,
Locked Bag 1797, Penrith, NSW 2751, Australia.
Ain Y. De Horta School of Science, Western Sydney University,
Locked Bag 1797, Penrith, NSW 2751, Australia.
Jelena Petrović Astronomical Observatory, Volgina 7, 11060 Belgrade, Serbia
Gojko Djurašević Astronomical Observatory, Volgina 7, 11060 Belgrade, Serbia

###### Abstract

Photometric observations and analysis of twelve previously poorly studied
contact binary systems is presented. All show total eclipses and have
extremely low mass ratios ranging from 0.072 to 0.15. Also, all show
characteristics of orbital instability with mass ratios within the theoretical
orbital instability range. Although none demonstrate a significant O’Connell
effect at least nine of the systems have other indicators of increased
chromospheric and magnetic activity.

Contact Binary, Low Mass Ratio, Red Nova

##  1 Introduction

As the number of identified contact binary systems grows with new discoveries
from an ever growing number of sky surveys, the interest in their
investigation has also increased. The confirmation of Nova 2008 Sco (V1390
Sco) as a red nova resulting from the merger of contact binary components
(Tylenda et al., 2011) has also intensified interest in the investigation of
orbital stability of contact binary systems (Wadhwa et al., 2021; Gazeas et
al., 2021; Christopoulou et al., 2022; Liu et al., 2023). Theoretical models
of contact binary evolution predict a likely merger event leading to a rapidly
rotating cool giant star as the final outcome (Soker & Tylenda, 2003;
Eggleton, 2012; Stȩpień, 2011; Robertson & Eggleton, 1977). Although the
predicted rate of Galactic merger events is high at once every 2 to 3 years,
brighter events suitable for study are likely to be much rarer at once every
10 or more years (Kochanek et al., 2014).

Investigators such Rasio & Shapiro (1995), Li et al. (2007) and Arbutina
(2007, 2009) have clearly shown that merger events are likely to occur when
the mass ratio of the components is low. More recently Wadhwa et al. (2021)
linked the instability parameters to the mass of the primary ($M_{1})$ and
concluded that there is not one global minimum mass ratio at which merger will
take place, rather the instability mass ratio can range from below 0.05 to
above 0.2 for systems where $0.6M_{\sun}<M_{1}<1.4M_{\sun}$. Although a number
of studies, each with a large (10 or greater) number of contact binaries with
extremely low mass ratios, have recently been reported they did not identify
any system that meet the instability criteria (Gazeas et al., 2021;
Christopoulou et al., 2022; Li et al., 2022; Liu et al., 2023). Wadhwa et al.
(2022a) recently linked the amplitude of totally eclipsing contact binaries
with the mass ratio which when used in conjunction with estimates of the mass
of the primary can result in the rapid identification of potentially unstable
systems. They identified over 40 new potential merger candidates based on the
analysis of survey photometric data. In this study we perform follow up ground
based observations on nine systems identified by (Wadhwa et al., 2022a) and
three other systems identified from the All Sky Automated Survey - Super Nova
(ASAS-SN) (Shappee et al., 2014; Jayasinghe et al., 2020) using the criteria
described in Wadhwa et al. (2022a). Basic identification data for each system
is presented in Table 1. Almost all ground based surveys have been carried out
using either small aperture telescopes or telephoto lenses. The angular
resolution is therefore quite poor ranging from 9” to over 60” (Drake et al.,
2017; Shappee et al., 2014; Jayasinghe et al., 2020; Thiemann et al., 2021;
Pojmanski, 2002). The risk of potential blending of the survey light curves
was highlighted by Wadhwa et al. (2022b) where a system thought to be of
extreme low mass ratio and a potential merger candidate is in fact a high mass
ratio system with no indications of orbital instability. Each system reported
here was checked against $Gaia$ EDR 3 (Anders et al., 2022; Gaia Collaboration
et al., 2022) to ensure there was no other star system brighter than 17th
magnitude within 10” of the target. We perform light curve analysis and
compare our results with reported data from survey analysis, in addition, we
show that all systems described demonstrate features of orbital instability,
have secondaries that are denser and most show evidence of chromospheric
activity in the absence of star spots.

Name | Abbreviation | Comparison Star | Check Star
---|---|---|---
ASAS J045814+0643.1 | A0458 | TYC 97-681-1 | TYC 97-273-1
ASAS J051459-7356.3 | A0514 | TYC 9174-356-1 | 2MASS 05140657-7351334
ASAS J100101-7958.6 | A1001 | TYC 9404-105-1 | 2MASS 09593687-7954508
V396 Lup | V396 Lup | TYC 7851-983-1 | 2MASS16031759-3751375
ASAS J170715-5118.7 | A1707 | TYC 8340-637-1 | 2MASS 17072965-5118000
ASAS J184644-2736.4 | A1846 | TYC 6867-2234-1 | TYC 6867-2250-1
ASAS J202231-4452.5 | A2022 | TYC 7961-1064-1 | TYC 7961-842-1
ASAS J204452+0622.6 | A2044 | 2MASS 20445329+0619549 | 2MASS 20444317+0619507
ASAS J213219-5351.6 | A2132 | UCAC2 8680692 | UCAC4 181-218564
SSS-J221327.1-445401 | S2213 | UCAC2 12815832 | 2MASS 22140108-4453567
ASAS J225826-2603.6 | A2258 | TYC 6974-1080-1 | TYC 6974-880-1
ASAS J234823-4054.7 | A2348 | TYC 8018-55-1 | 2MASS 23480833-4046082
Table 1: Name, Abbreviations (used in all other tables and text) along with
comparison and check stars for 12 contact binary systems

##  2 Photometric Observations, Absolute magnitudes and Mass of the Primary
Component

###  2.1 Photometric Observations

We acquired dual band (V and R) images of 12 contact binary systems from April
2020 to January 2023 using either the 0.6m Western Sydney University (WSU)
telescope equipped with a cooled SBIG 8300 CCD camera and standard Johnson
$BVR$ filters or the 0.4m telescopes of the Las Cumbres Observatory (LCO)
network equipped with SBIG STL-6303 CCD camera and Bessel V,B and Sloan r’
filters. In addition to V and R band images, B band images were obtained only
during eclipses to document $B-V$ magnitude for each system. The LCO network
automatically calibrates all images while the images obtained at the WSU were
calibrated with multiple flat, dark and bias frames. Differential photometry
was performed for each system using the AstroImageJ (Collins et al., 2017)
software package. All observations with software reported error of greater
than 0.01 magnitude were excluded. Observation dates, number of observations,
exposure times, maximum V band brightness, V Band amplitude and $B-V$
magnitude are summarised in Table 2.

Name | Obs Date | Obs (V,R) | Exp Times (V,R) | Max (V) | Ampl (V) | B-V | $M_{V1}$ | $M_{1}/M_{\sun}$
---|---|---|---|---|---|---|---|---
A0458 | 11/21 - 01/23 | 285, 350 | 45s,45s | 11.76 | 0.26 | 0.46 | $3.88\pm 0.01$ | $1.17\pm 0.02$
A0514 | 01/22 - 10/22 | 320, 300 | 45s,45s | 11.71 | 0.37 | 0.71 | $4.90\pm 0.02$ | $0.98\pm 0.01$
A1001 | 04/20 - 04/20 | 373, 375 | 40s,40s | 11.47 | 0.33 | 0.91 | $5.78\pm 0.01$ | $0.83\pm 0.01$
V396 Lup | 06/21 - 07/21 | 670, 240 | 30s,25s | 10.88 | 0.36 | 0.83 | $5.20\pm 0.02$ | $0.95\pm 0.01$
A1707 | 07/21 - 08/21 | 510, 500 | 40s,35s | 11.19 | 0.44 | 0.55 | $3.67\pm 0.02$ | $1.22\pm 0.01$
A1846 | 08/21 - 08/21 | 300, 410 | 40s,40s | 11.78 | 0.39 | 0.78 | $5.43\pm 0.02$ | $0.90\pm 0.01$
A2022 | 08/21 - 08/21 | 550, 300 | 45s,45s | 12.09 | 0.31 | 0.53 | $4.38\pm 0.01$ | $1.11\pm 0.01$
A2044 | 08/22 - 10/22 | 440, 280 | 50s,45s | 12.67 | 0.30 | 0.62 | $4.18\pm 0.02$ | $1.13\pm 0.01$
A2132 | 10/21 - 09/22 | 255, 250 | 55s,55s | 12.89 | 0.28 | 0.55 | $4.48\pm 0.02$ | $1.04\pm 0.01$
S2213 | 09/21 - 09/21 | 250, 300 | 45s,45s | 12.49 | 0.27 | 0.55 | $4.05\pm 0.02$ | $1.15\pm 0.02$
A2258 | 09/21 - 10/21 | 255, 330 | 35s,35s | 11.49 | 0.26 | 0.60 | $4.21\pm 0.02$ | $1.12\pm 0.01$
A2348 | 11/21 - 08/22 | 350, 350 | 50s,50s | 12.83 | 0.30 | 0.54 | $4.16\pm 0.03$ | $1.11\pm 0.02$
Table 2: Observation dates, number of observations, exposure times, light
curve parameters and spectral classification of the twelve reported contact
binary systems. A1001, V396 Lup, A1707 and A1846 were observed with the WSU
telescope others through the LCO telescope network. ∗ = Spectral
classification from published literature.

Based on the times of minima observed and available V band photometric data
from the ASAS-SN and All Sky Automated Survey (ASAS) (Pojmanski, 2002) we
updated the orbital elements as summerised in Table 3.

Name | Epoch (HJD) | Period (d)
---|---|---
A0458 | $2459543.900671\pm 0.000728$ | $0.333482\pm 0.000012$
A0514 | $2459594.108439\pm 0.000520$ | $0.345729\pm 0.000085$
A1001 | $2458961.971160\pm 0.000920$ | $0.279133\pm 0.000085$
V396 Lup | $2459372.012959\pm 0.000420$ | $0.363248\pm 0.000060$
A1707 | $2459430.984062\pm 0.000511$ | $0.525901\pm 0.000040$
A1846 | $2459439.005458\pm 0.000225$ | $0.302853\pm 0.000030$
A2022 | $2459429.668598\pm 0.000986$ | $0.345011\pm 0.000006$
A2044 | $2459818.566466\pm 0.000564$ | $0.370536\pm 0.000029$
A2132 | $2459509.962001\pm 0.000488$ | $0.316348\pm 0.000058$
S2213 | $2459459.704234\pm 0.000621$ | $0.369894\pm 0.000005$
A2258 | $2459436.659418\pm 0.000267$ | $0.327648\pm 0.000060$
A2348 | $2459545.581557\pm 0.000223$ | $0.347195\pm 0.000052$
Table 3: Updated orbital elements

###  2.2 Absolute Magnitude of the Primary Component

As all systems described have a very low mass ratio (see below) and all
demonstrate total eclipses we are presented with a prospect of being able to
estimate the absolute magnitude of the primary component using direct
observations. The apparent magnitude of the secondary eclipse represents the
apparent magnitude of the primary. If the distance is known, the absolute
magnitude of the primary ($M_{V1}$) can be estimated as follows:

$M=m-5log_{10}d+5$ (1)

Where M is the absolute magnitude, m is the apparent magnitude and d is the
distance in parsecs.

The $GAIA$ mission (Anders et al., 2022) provides highly accurate estimates of
the distance particularly of nearby stars and as such makes distance based
estimate of the absolute magnitude a reality. Interstellar extinction however
must be taken into consideration when determining the apparent magnitude. As
most dust maps report line of sight extinction to infinity and all our systems
have distance estimations well below 1 kpc we corrected the extinction for
distance as follows: Using Schlafly & Finkbeiner (2011) dust maps we
determined the line of sight reddening at infinity $E(B-V)_{\infty}$. We used
the $Gaia$ distance to scale this value $E(B-V)_{d}$ using the equation (Bilir
et al., 2008):

$E(B-V)_{d}=E(B-V)_{\infty}\Bigg{[}1-\mathrm{exp}\bigg{(}-\frac{|d{\cdot\mathrm{sin}(b)}|}{h}\bigg{)}\bigg{]}$
(2)

In the equation $b$ is the galactic latitude of the system, $d$ is the $Gaia$
distance (in parsecs) and $h$ is the galactic scale height, taken as $h=125pc$
as per Bilir et al. (2008). The total extinction ($A_{V}$) was then estimated
as $A_{V}=E(B-V)_{d}\times 3.1$.

The estimated absolute magnitude for each system is summarised in Table 2. The
largest contributor to the estimated error in the absolute magnitude was the
reported error in the distance estimate.

###  2.3 Mass of the Primary Component

One of the main aims of this report is to assess potential orbital instability
of the reported contact binary systems. Recently, Wadhwa et al. (2021)
developed new relationships linking the mass of the primary ($M_{1}$), the
degree of contact and the instability mass ratio ($q_{inst}$). One can
estimate the range of mass ratios over which a system may become unstable if
the mass of the primary is known. Direct estimation of the mass of the primary
is not possible, one is therefore reliant on secondary observations and
calibrations. Primary components of contact binary systems follow a main
sequence profile (Yildiz & Doğan, 2013). For this study we use two methods to
estimate the mass of the primary. The Two Micron All Sky Survey (2MASS)
(Skrutskie et al., 2006) acquired simultaneous photometry in multiple infrared
bands. As our first estimate of the mass of the primary we used the 2MASS
$J-H$ magnitudes and the April 2022 update tables of Pecaut & Mamajek (2013)
for low mass ($0.6M_{\sun}<M_{1}<1.4M_{\sun}$) main sequence stars to
interpolate the mass of the primary. Our distance based estimate is linked to
the absolute magnitude of the primary determined above. Again we use the April
2022 update tables of Pecaut & Mamajek (2013) for low mass
($0.6M_{\sun}<M_{1}<1.4M_{\sun}$) main sequence stars to interpolate the mass
of the primary. We use the mean of the colour and distance based estimates as
our adopted value for the mass of the primary. The recorded errors were
propagated when calculating the instability mass ratio range (see below).
Summary results for the mass of the primary are recorded in Table 2.

##  3 Light Curve Analysis and the Mass Ratio

Having established an estimate for the mass of the primary the next (and only)
other parameter required to determine potential orbital stability is the mass
ratio of the system. In the absence of radial velocity measurements light
curve analysis of contact binary systems can be successfully carried out if
total eclipses are present (Terrell & Wilson, 2005). One of the input
parameters for such analysis is the temperature of the primary component
($T_{1}$). There is no standard method for assigning the value for the
temperature of the primary. Colour based estimations have been the mainstream,
however these have proven troublesome, for example the ViziR database reports
a range of over 3000K for one of our targets (V396 Lup). Given the lack of a
standard approach many investigators (see e.g. Zhou et al., 2018; Chang et
al., 2022; Guo et al., 2022) are now moving to low resolution spectra (where
available) to aid in determining the primary’s temperature. To this end for
this study we acquired spectra using the LCO network of 2m telescopes equipped
with FLOYDS spectrograph, a cross-dispersed, spectrograph with variable
resolution of R=400 to R=700 for 10 of the systems. The reduction pipeline for
the FLOYDS spectral data is fully automated and described in detail on the LCO
website (https://lco.global/documentation/data/floyds-pipeline/ - accessed
14/1/2023). For A0514 we used the published spectral data (Cannon & Pickering,
1993). We were unable to obtain a spectrum for A2044 due to technical reasons
during the allotted time. All FLOYDS spectra were visually compared to
standard library spectra of main sequence stars (Jacoby et al., 1984; Pickles,
1998) to determine the spectral class for each system. Representative FLOYDS
spectra and matching library spectra are shown in Figure 1.

For this study we use the mean ($\pm{SD}$) of four estimates for the
temperature of the primary. We estimate effective temperature of the primary
based on our extinction corrected value of $B-V$, the observed spectral class,
the 2MASS $J-H$ colour and the reported value from the GAIA third data release
all adjusted to the nearest 10. For consistency we use the April 2022 update
tables of Pecaut & Mamajek (2013) for colour and spectral calibrations.
Summary of the spectral class, temperatures and the adopted mean value of the
four calibrations is presented in Table 4. In the case of A2044 where no
spectral observations were available we adopt the mean of the two colour
calibrations and the GAIA estimate.

Name | $B-V(K)$ | $J-H(K)$ | $Sp(K)$ | $GAIA(K)$ | $Mean\pm SD$ | Sp Class
---|---|---|---|---|---|---
A0458 | 6650 | 6060 | 7020 | 7110 | $6710\pm 410$ | F1
A0514 | 5810 | 5600 | 5930 | 5310 | $5660\pm 240$ | G0∗
A1001 | 5140 | 4870 | 5380 | 4990 | $5100\pm 190$ | G9
V396 Lup | 5520 | 5550 | 6050 | 5930 | $5760\pm 230$ | F9
A1707 | 6610 | 6190 | 6750 | 6550 | $6530\pm 210$ | F3
A1846 | 5530 | 5370 | 5550 | 5080 | $5380\pm 190$ | G7
A2022 | 6280 | 6060 | 5930 | 6240 | $6130\pm 140$ | G0
A2044 | 6070 | 6060 | - | 5900 | $6010\pm 80$ | -
A2132 | 6190 | 5710 | 6180 | 6020 | $6030\pm 190$ | F8
S2213 | 6190 | 6060 | 6280 | 5970 | $6130\pm 120$ | F7
A2258 | 6000 | 5990 | 6050 | 5910 | $5990\pm 50$ | F9
A2348 | 6180 | 5930 | 6550 | 6110 | $6190\pm 230$ | F5
Table 4: Effective temperatures based on $B-V$, $J-H$, spectral class and
GAIA. Mean value $\pm(SD)$ was the adopted value for light curve analysis. In
the last column we record the spectral class for each system. ∗ In the case of
A0514 we use the published data (Cannon & Pickering, 1993).

The 2013 version of the Wilson-Devinney (WD) code with Kurucz atmospheres
(Wilson, 1990; Kallrath et al., 1998; Nelson, 2021) was used to obtain the
photometric solution for each system. As all systems show total eclipses
accurate photometric solutions are possible and as there was no significant
asymmetry at maximum brightness (i.e no significant O’Connell Effect) only
unspotted solutions were obtained. As the effective temperature of the primary
is less than 7200K in all cases, the gravity coefficients were set as
$g_{1}=g_{2}=0.32$, and bolometric albedoes as $A_{1}=A_{2}=0.5$. Logarithmic
(Nelson & Robb, 2015) limb darkening coefficients were interpolated from the
2019 update of van Hamme (1993).

Simultaneous solutions were obtained for both the V and R band in all cases.
Mass ratio ($q$) search grid method was used to find the best solution for a
range of fixed mass ratios from 0.05 to 20. The initial search was performed
with mass ratio increments of 0.1 up to $q=1$ and then 0.2 up to $q=10$ and
0.5 increments up to $q=20$. The search was then refined near the best
solution in increments of 0.05 and finally in increments of 0.01. During the
search procedure the orbital inclination ($i$), the surface potential,
temperature of the secondary component ($T_{2}$) and the dimensionless
luminosity of the primary ($L_{1}$) were the adjustable parameters and all of
which were adjusted between each iteration. The iterations were carried out
until the suggested adjustment for all adjustable parameters was less than the
reported standard deviation. During the last iteration the mass ratio was also
made a adjustable parameter.

To determine the effects of the variable temperature of the primary each
system was further modelled with mass ratio increments of 0.01 around the mean
temperature solution with the temperature of the primary adjusted both to the
upper and lower limits reported in Table 4. As has been well established the
light curve shape of contact binaries is almost completely due to the system
geometry with the degree of contact, mass ratio and inclination being the
major contributors (Rucinski, 1993, 2001). The shape of contact binary light
curves places a tight constrain on the component temperature ratio
($T_{2}/T_{1}$) but not on individual component temperatures (Rucinski, 1993;
Sun et al., 2020). With change in $T_{1}$ we expected change in $T_{2}$ but no
significant change in the geometric parameters. This proved to be the case
with the light curve solution converging to the same mass ratio regardless of
the input value of $T_{1}$ except in the case of A1846 where the mean $T_{1}$
solution converged at $q=0.151$ and the higher $T_{1}$ solution converged at
$q=0.150$. Similarly the inclination variation between high, mean and low
values for $T_{1}$ were within 0.5∘ and within the reported standard
deviations reported by the WD software. The dimensionless potential and the
geometric mean of the fractional radii again varied minimally. Although the
errors in geometric parameters were small with change in temperature of the
primary they were propagated to the mean solution reported. Apart from the
mass ratio and fractional radii (reported to two decimal places) no other
parameter resulting from modelling of the light curve is further used in this
study. We report light curve solution parameters pertinent to this study in
Table 5 with errors propagated relative to the input errors in $T_{1}$. Fitted
and observed light curves are illustrated in Figure 2.

Figure 1: The low resolution FLOYDS spectra were visually matched to standard
main sequence star spectra to determine the spectral class for 10 of the
systems.

##  4 Orbital stability

Orbital stability and potential merger of contact binary systems has received
considerable attention recently (Wadhwa et al., 2021; Christopoulou et al.,
2022; Liu et al., 2023). Wadhwa et al. (2021) developed new relationships
linking the mass of the primary, the degree of contact and the instability
mass ratio ($q_{inst}$). Following their methodology, and adopting the values
of the gyration radii of the components as described, we calculated $q_{inst}$
for each system for fill-out factor values of 0 and 1 (Table 5). This provides
a range of mass ratios where orbital instability is likely. As noted by
Christopoulou et al. (2022) a small uncertainty in the estimation of the mass
of the primary propagates to a significant uncertainty in the estimation of
orbital instability parameters. This is confirmed by this study where the mean
mass of the primary components of the systems presented in this study is
$1.06M_{\sun}$. At this mass the instability mass ratio range for fill-out
factor 0 to 1 extends from 0.094 to 0.109. A 10% change in the mass of the
primary (either above or below) results in a greater than 17% change in the
instability mass ratio range. Thus for the purpose of this study we consider
any system where the mass ratio is within 17% of the maximum instability mass
ratio to be potentially unstable. Based on this criteria ten of the twelve
systems lie within the instability mass ratio range taking into account
reported errors. One (A1846) is within 7% and one (A2348) 17% of the maximum
instability mass ratio. We note that the distance estimation of A2348 has a
reported error near 10% so we consider it reasonable to include the system as
potentially unstable. The instability mass ratio ranges are summarised in
Table 5.

Manual analysis of the survey photometric data for 9 of the systems presented
here was reported by Wadhwa et al. (2022a). The survey data mass ratio
estimate was within 10% of the mass ratio based on dedicated observations
except in the case of A2044 and A2132 where the survey data estimates of the
mass ratio were 20% and 12%, respectively, below the values reported here.The
findings confirm that generally survey photometric data provides a good
starting point in the selection of potential low mass ratio systems and merger
candidates. Wadhwa et al. (2022a) found that all 9 of the common systems could
be regarded as potential merger candidates as their estimation of the mass of
the primary was based on $J-H$ calibration. Looking at the two candidates in
our report A1846 and A2348 which are just outside the instability range, have
$J-H$ calibrated mass of the primary, 0.89 and 1.07, respectively, both lower
than the combined with distance estimation used in this study. If the single
colour estimation was used than the 2 would fall within the instability range.

Figure 2: Observed and modelled light curves for the 12 systems. The green and red light curves represent the V and R(r’) bands while the purple curve represents the check star. The black line represents the WD modelled curve. The vertical axis labelled nrFlux represents the normalised flux which has been arbitrarily shifted vertically for clarity Name | $T_{1}(K)$ | $T_{2}(K)$ | Incl $(^{\circ})$ | Mass Ratio ($q$) | Fr Radii ($r_{1,2}$) |  $q_{inst}$ Range
---|---|---|---|---|---|---
A0458 | $6710\pm 410$ | $6190\pm 380$ | $88.4\pm 0.8$ | $0.089\pm 0.001$ |  $0.60\pm 0.01$, $0.21\pm 0.01$ |  $0.079\pm 0.003$ \- $0.09\pm 0.003$
A0514 | $5660\pm 240$ | $5820\pm 240$ | $73.5\pm 0.6$ | $0.12\pm 0.001$ |  $0.59\pm 0.01$, $0.26\pm 0.01$ |  $0.105\pm 0.002$ \- $0.123\pm 0.002$
A1001 | $5100\pm 190$ | $4960\pm 180$ | $72.4\pm 0.7$ | $0.138\pm 0.003$ |  $0.56\pm 0.01$, $0.24\pm 0.01$ |  $0.129\pm 0.002$ \- $0.155\pm 0.002$
V396 Lup | $5760\pm 230$ | $5670\pm 230$ | $78.0\pm 0.4$ | $0.133\pm 0.002$ |  $0.58\pm 0.01$, $0.26\pm 0.01$ |  $0.110\pm 0.002$ \- $0.130\pm 0.002$
A1707 | $6530\pm 210$ | $6080\pm 190$ | $80.8\pm 0.6$ | $0.072\pm 0.002$ |  $0.62\pm 0.01$, $0.21\pm 0.01$ |  $0.072\pm 0.002$ \- $0.082\pm 0.002$
A1846 | $5380\pm 190$ | $5720\pm 210$ | $87.0\pm 0.7$ | $0.150\pm 0.002$ |  $0.56\pm 0.01$, $0.25\pm 0.01$ |  $0.117\pm 0.001$ \- $0.140\pm 0.002$
A2022 | $6130\pm 140$ | $6260\pm 150$ | $80.1\pm 0.7$ | $0.100\pm 0.003$ |  $0.60\pm 0.01$, $0.24\pm 0.01$ |  $0.087\pm 0.002$ \- $0.100\pm 0.002$
A2044 | $6010\pm 80$ | $5940\pm 80$ | $87.2\pm 0.6$ | $0.088\pm 0.001$ |  $0.61\pm 0.01$, $0.22\pm 0.01$ |  $0.084\pm 0.003$ \- $0.097\pm 0.004$
A2132 | $6030\pm 190$ | $6040\pm 190$ | $73.4\pm 0.5$ | $0.113\pm 0.002$ |  $0.58\pm 0.01$, $0.23\pm 0.01$ |  $0.096\pm 0.001$ \- $0.112\pm 0.002$
S2213 | $6130\pm 120$ | $5960\pm 120$ | $90.00^{+0.00}_{-1.5}$ | $0.088\pm 0.003$ |  $0.61\pm 0.01$, $0.22\pm 0.01$ |  $0.081\pm 0.003$ \- $0.093\pm 0.003$
A2258 | $5990\pm 50$ | $5970\pm 50$ | $90.00^{+0.00}_{-0.6}$ | $0.087\pm 0.002$ |  $0.57\pm 0.01$, $0.22\pm 0.01$ |  $0.081\pm 0.003$ \- $0.098\pm 0.002$
A2348 | $6190\pm 230$ | $6070\pm 220$ | $73.9\pm 0.5$ | $0.117\pm 0.002$ |  $0.59\pm 0.01$, $0.24\pm 0.01$ |  $0.087\pm 0.003$ \- $0.100\pm 0.003$
Table 5: Light curve solution summary of the main pertinent parameters.
($r_{1,2}$) = geometric fractional radii and ($q_{inst}$) = the instability
mass ratio range.

##  5 Density and Chromospheric Activity

###  5.1 Component Density

It is well known that the secondary components are larger than their main
sequence counterparts. In addition to change in the radius some researchers
(Yildiz & Doğan, 2013) postulate that distortions of the secondary may result
from it actually being initially a larger primary that has lost mass to the
current primary and due to retention of higher mass core constituents is
likely to be denser than the primary. Kähler (2004) argues that the densities
of the components will always be different and that the difference between
density of the primary ($\rho_{1}$) and that of the secondary ($\rho_{2}$)
will always be negative.

As noted by (Mochnacki, 1981) the density of the components in ($\rm
gcm^{-3}$) can be expressed as a function of the period, relative radii and
mass ratio. The density difference ($\Delta\rho$) between the components can
be expressed as:

$\Delta\rho=\frac{0.0189}{r_{1}^{3}(1+q)P^{2}}-\frac{0.0189q}{r_{2}^{3}(1+q)P^{2}}$
(3)

All our system confirm the finding that the secondary is much denser and that
$\Delta\rho$ is negative. The findings (to two decimal places) are summerised
in Table 6

###  5.2 Chromospheric Activity

Contact binary systems typically have synchronised rotation with short periods
of less than 1 day. Such rapid rotation rates are thought to increase magnetic
activity related to contact binary systems (Gharami et al., 2019). Increased
magnetic activity may manifest itself as magnetic stellar wind and magnetic
breaking leading to loss of angular momentum from the system (Li et al.,
2004). Although it is difficult to directly measure angular momentum loss
there are potential secondary indicators of increased magnetic activity, such
as strong ultraviolet emissions and/or increased chromospheric activity (Li et
al., 2004; Rucinski & Vilhu, 1983; Vilhu, 1983). There are some chromospheric
and magnetic signals associated with contact binaries with the most readily
observed light curve feature being the asymmetry in the maxima (O’Connell
effect). Apart from the O’Connell effect and incorporation of starspots, the
analysis of light curves provides little indication of chromospheric activity.
None of the systems described in this report show significant variation of
maxima, however, that does not exclude the presence of significant magnetic
and chromospheric activity. High energy spectral emissions provide a much
clearer indicators of enhanced chromospheric/magnetic activity. In low mass
dwarfs the chromospheric emissions are by photospheric light in the visual
band. Higher energy emissions particularly in the far-ultraviolet region
provide a less obscured alternative (Smith & Redenbaugh, 2010). The GALEX
(Galaxy Evolution Explorer) satellite imaged the sky in the far-ultraviolet
band (FUV) centered on 1539 Å and this can be employed to explore
chromospheric activity of contact binaries(Smith & Redenbaugh, 2010).

The $R_{\rm HK}^{\prime}$ index (Noyes et al., 1984) is an excepted measure of
chromospheric emission strength with $\log R_{\rm HK}^{\prime}\geq-4.75$
characteristic of a more active star (Henry et al., 1996). Smith & Redenbaugh
(2010) matched GALEX FUV magnitudes ($m_{\rm FUV}$) to the $\log R_{\rm
HK}^{\prime}$ for dwarf stars to derive the $\Delta(m_{\rm FUV-B})$ colour
excess:

$\Delta(m_{\rm FUV-B})=(m_{\rm FUV}-B)-(m_{\rm FUV}-B)_{\rm base}$ (4)

where

$(m_{\rm FUV}-B)_{\rm base}=6.73(B-V)+7.43,$ (5)

They deduced that for active stars the colour excess was always below -0.5 and
often less than -1.0. Less active stars usually had colour excess well above
-0.5.

Six of our twelve systems were observed by the GALEX mission with recorded FUV
magnitudes. We calculated the colour excess for all six using the above
relationships and show that all have an ultraviolet colour excess well below
-0.5 indicative of significant chromospheric/magnetic activity in the absence
of light curve features. Magnetically active features, sometimes referred to
as plages (Hall, 2008), are responsible for the chromospheric emissions while
starspots are regions of intense magnetic activity leading to suppression of
convection in the photosphere (Solanki, 2003). Although starspots are
accompanied by plague regions, the reverse is not always true such that
chromospheric activity is possible without photospheric starspots (Mandal et
al., 2017). The UV colour excesses are summarised in Table 6.

Coronal and chromospheric activity combined with fast synchronous rotation of
the common envelopes may also result in X-Ray emissions from contact binaries
(Kandulapati et al., 2015; Gondoin, 2004). Seven of our systems (A0458, A0514,
V396 Lup, A1846, A2044, S2213, A2348) were reported as potential X-Ray sources
(Liu et al., 2022) including three (A0514, V396 Lup and A2044) without FUV
observations bringing to a total of nine systems with signs suggestive of
increased chromospheric activity without characteristic photospheric changes.

Name | $\rho_{1}(\rm gcm^{-3})$ | $\rho_{2}(\rm gcm^{-3})$ | $\Delta\rho$ | $\Delta(m_{\rm{FUV}}-B)$
---|---|---|---|---
A0458 | $0.72\pm 0.01$ | $1.49\pm 0.01$ | $-0.77$ | $-3.34$
A0514 | $0.69\pm 0.01$ | $0.96\pm 0.01$ | $-0.27$ | $--$
A1001 | $1.21\pm 0.02$ | $2.13\pm 0.02$ | $-0.92$ | $--$
V396 Lup | $0.65\pm 0.03$ | $0.96\pm 0.02$ | $-0.31$ | $--$
A1707 | $0.27\pm 0.01$ | $0.50\pm 0.02$ | $-0.23$ | $--$
A1846 | $1.02\pm 0.01$ | $1.72\pm 0.02$ | $-0.70$ | $--$
A2022 | $0.67\pm 0.02$ | $1.04\pm 0.02$ | $-0.37$ | $-2.30$
A2044 | $0.56\pm 0.01$ | $1.05\pm 0.01$ | $-0.49$ | $--$
A2132 | $0.87\pm 0.02$ | $1.58\pm 0.02$ | $-0.71$ | $-2.80$
S2213 | $0.56\pm 0.01$ | $1.05\pm 0.02$ | $-0.49$ | $-2.32$
A2258 | $0.87\pm 0.01$ | $1.32\pm 0.02$ | $-0.45$ | $-2.56$
A2348 | $0.68\pm 0.01$ | $1.19\pm 0.02$ | $-0.51$ | $-3.23$
Table 6: Density and ultraviolet colour excess.

##  6 Discussion and Conclusion

Photometric observations and light curve solution of twelve low mass contact
binary systems is presented. All were found to be of extreme low mass ratio
ranging from 0.072 to 0.15. Compared to theoretical parameters all have a mass
ratios indicative of likely orbital instability and are likely merger (red
nova) candidates. Regular high cadence future observations are encouraged to
accurately determine period variations and brightness changes which maybe an
indicator of impending merger (Tylenda et al., 2011). A thorough study of
orbital behavior of contact binaries have been performed recently by Zhang &
Qian (2020); Poro et al. (2022); Loukaidou et al. (2022). The first and the
last paper discuss in particular the systems close to the orbital period cut-
of. Ultra-short period systems seem to have generally stable orbits and do not
show evidence of significant period decrease. From the viewpoint of Darwin
instability, this is understandable, since these are not such a low-mass ratio
systems and it may take a long evolutionary time for ultra-short contact
binaries to reach an extremely low mass ratio and become dynamically unstable
(Loukaidou et al., 2022).

Review of some astrophysical characteristics confirm similarity with other
contact binary systems with the secondary considerably brighter and larger
than main sequence stars of similar mass. In addition, the secondary in all
cases is significantly denser than the current primary. Chromospheric activity
in contact binary is usually detected as a variation in the two maxima from
the light curve due to starspots. As we have shown in this study non-
photospheric markers probably can detect such activity without typical light
curve features. Another high energy feature of excess chromospheric/magnetic
activity is X-Ray emissions (Gondoin, 2004). Only a small fraction of contact
binary star systems are X-ray sources (Liu et al., 2022) such that presence of
X-Ray emissions can be taken as a marker of increased magnetic activity and
potential indicator for an increased loss of angular momentum and orbital
instability. Seven of twelve systems reported have been catalogued as having
X-Ray emissions (Liu et al., 2022). Also, as noted above the FUV magnitude can
be used as a marker for chromospheric activity. Another interesting point of
note is how much brighter relative to main sequence stars are the contact
binaries in the FUV band. Findeisen & Hillenbrand (2010) have estimated the
absolute FUV magnitudes ($M_{\rm FMS}$) for single main sequence stars from B8
to M2 spectral types. For the six systems from our sample with FUV magnitude
we find that all systems are between 0.5 and over 3.0 magnitude brighter in
the FUV band relative to their main sequence counterparts.

Progress in the detection and monitoring of potential contact binary merger
candidates is rapidly progressing. There now exists a theoretical framework to
determine the potential for merger from the estimate of the mass of the
primary and the mass ratio of the system (Wadhwa et al., 2021). Models and
techniques have been developed to rapidly select potential merger candidates
from survey photometry (Wadhwa et al., 2022a). However, follow up dedicated
observations are required due to the low resolution of most survey data
(Wadhwa et al., 2022b).

Acknowledgements. Based on data acquired on the Western Sydney University,
Penrith Observatory Telescope. We acknowledge the traditional custodians of
the land on which the Observatory stands, the Dharug people, and pay our
respects to elders past and present.
This research has made use of the SIMBAD database, operated at CDS,
Strasbourg, France.
This research has made use of the VizieR catalogue access tool, CDS,
Strasbourg, France (DOI : 10.26093/cds/vizier).
B. Arbutina, G. Djurašević and J. Petrović acknowledge the funding provided by
the Ministry of Science, Technological Development and Innovation of the
Republic of Serbia through the contracts 451-03-47/2023-01/200104 (BA) and
451-03-47/2023-01/200002 (GDj, JP).

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