Optical CCD photometry and Rosat X-ray spectral analysis of the shortest period CV EI Psc (1RXS J232953.9 + 062814)

New Astronomy 14 (2009) 160–165

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Optical CCD photometry and Rosat X-ray spectral analysis of the shortest period CV EI Psc (1RXS J232953.9 + 062814) _ F. Gök a,*, G. Ikis Gün b, E. Aktekin a, A. Sezer a, M. Altan c a

Faculty of Art and Sciences, Department of Physics, Akdeniz University, 07058 Camp., Antalya, Turkey Faculty of Art and Sciences, Department of Physics, Çanakkale Onsekiz Mart University, 17100 Terzioglu Camp., Çanakkale, Turkey c Faculty of Sciences, Physics Department, Anadolu University, 26470 Yunusemre Camp., Eskisßehir, Turkey b

a r t i c l e

i n f o

Article history: Received 13 July 2007 Received in revised form 22 June 2008 Accepted 7 July 2008 Available online 29 July 2008 Communicated by M. van der Klis PACS: 97.30.Qt Keywords: Close binary stars Cataclysmic variables Dwarf novae Accretion Accretion discs X-ray spectra Mass accretion rate EI Psc

a b s t r a c t We present here the optical observations of EI Psc (1RXS J232953.9 + 062814) through Turkish National Observatory (TUG) with RTT150 cm Russian–Turkish joint telescope and its X-ray observation using the ROSAT archival data. Our optical observations reveal a period of 0.0408 days (58.75 min) which is rather different from its early value of 0.046 days (66.24 min) as reported by [Schmeer, P., 2001. vsnetalert6830, ]. Also possible periodicities as well as any QPOs are studied without having any clear indication of it. Archival ROSAT RASS data are also analyzed for its X-ray spectra. The raw data were fitted with various spectral models and the best fit models are found to be that of Blackbody and Raymond–Smith with best fit temperatures of kT ¼ ð0:07  0:02Þ keV for blackbody model and kT ¼ ð0:13  0:04Þ keV for Raymond–Smith model while the column density fixed at 0:54  1021 cm2 . The estimated 0.1–2.4 keV flux is found to be in the range of between log F ¼ 13 and log F ¼ 14 erg cm2 s1 . The model dependent luminosity values were in the range of log L ¼ 29 erg s1 for Raymond–Smith model and log L ¼ 31 erg s1 for blackbody model. Using the well fitted temperature values, the mass of the primary value that is obtained by [Uemura, M., Kato, T., Ishioka, R., Yamaoka, H., Scgmeer, P., Starkey, D.R., Torii, K., Kawai, N., Urata, Y., Kohama, M., Yoshida, A., Ayani, K., Kawabata, T., Tanabe, K., Matsumoto, K., Kiyota, S., Pietz, J., Vanmunster, T., Krajci, T., Oksanen, A., Giambersio, A., 2002b. Superhump Evolution in the Ultrashort Period Dwarf Nova 1RXS J232953.9 + 062814. PASJ 54, 599–607] and the equations taken from literature, the mass accretion rate in the boundary layer is obtained to be ð1:58  0:14Þ  1021 gs1 for the blackbody model and ð2:2  0:052Þ  1019 gs1 for Raymond–Smith model. As a result of our study it seems that the system EI Psc has a very high mass accretion rate; and because of the observed soft X-ray photons and high mass accretion rates it has an optically thick boundary layer and M-type secondary star which can be a Brown Dwarf. Ó 2008 Elsevier B.V. All rights reserved.

1. Introduction Non-magnetic cataclysmic variables (hereafter CVs) are close binary systems and consist of a white dwarf (primary star), a late type main sequence star (secondary star) and an accretion disk. The secondary star fills its Roche lobe and starts to transfer matter on white dwarf via accretion disk. Inner regions of the accretion disks in CVs are referred to as ‘‘Boundary layers”. The basic theories express that the bulk of the X-rays are originated from this Boundary layer (Warner, 1995). Dwarf novae are one of the subclasses of CVs and eruptive variable stars. Amplitude of repetitive outburst is about 2–6 mag and recurrence time is about 20–300 days. The orbital period of dwarf * Corresponding author. Tel.: +90 2423102274; fax: +90 2422278911. E-mail address: [email protected] (F. Gök). 1384-1076/$ - see front matter Ó 2008 Elsevier B.V. All rights reserved. doi:10.1016/j.newast.2008.07.003

nova ranging from the shortest to the longest is about 2–12 h. The ones that have orbital periods of about 2 h are called SU UMa types (Warner, 1995). SU UMa types show two distinct outbursts: (i) a ‘‘normal” outburst with amplitude of about 2 mag, recurrence time of 20–30 days and a duration of a few days. (ii) A long ‘‘super outburst” with amplitude of about 6–8 mag, recurrence time of 300–? days and a duration of approximately two weeks. WZ Sge is an extreme case of SU UMa that gave its last super outburst after 23 years. The other extreme case, on the contrary, is RZ LMi stars or ER UMa stars which show super outburst cycles of only 20–50 days. Super outbursts often show photometric oscillations called super-hump with a period a few percent longer or shorter than orbital period of the system P orb . The reason for this is thought to be tidally driven precession of an eccentric disk. In other words, the eccentric disk is forced to precess by the perturbation from the secondary (Patterson, 2001).

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Mennickent et al. (2004) have pointed out that according to the current theories about the CVs which have hydrogen rich secondaries, if the secondary component becomes less and less massive because of the accretion process, the binary period of the system gets shorter. In the end, the secondary star becomes a Brown Dwarf which is the part of a system with 80 min orbital period. There is another scenario to explain the evolution of short period CVs. According to this scenario in some CVs which have a secondary star with intermediate mass, thermal-timescale mass transfer occurs and as a result ultra-short period systems are formed. These systems may have a secondary star which is hydrogen poor, low-mass and degenerate Brown dwarf (Mennickent et al., 2004 and the references therein). One of these ultra-short period systems is EI PSc. It was first discovered as an X-ray source with the ROSAT (RASS-BSC) and given catalog name of the system is 1RXS J232953.9 + 062814 (Voges et al., 1996). The coordinates are a2000 ¼ 23h 29m 54:3s and d2000 ¼ þ06o 280 1000 :9. The system was identified as a CV by Jingyao et al. (1998), Hu et al. (1998) and Wei et al. (1999). Optical identification has been made with V = 15.7 source. As a result of these observations EI PSc has been classified as a Dwarf Nova (DN). Light curve of the system was then described by two distinct states: one is its faint state: there are mostly hydrogen emission lines and TiO absorption bands in the optical spectrum shows that EI PSc is a hydrogen rich CV and has a M-type star as the secondary component. The other state is its bright state: there are absorption hydrogen lines in the optical spectrum. Schmeer (2001) found superhumps with amplitudes of 0.2–0.3 mag and a period of 0.046311 (12) days. These features point out that EI PSc is a SU UMa type dwarf nova. The orbital period of the system is below the period minimum of hydrogen rich CVs (about 1.3 h) because of this short superhump period (Uemura et al., 2001 and references therein). In this work, we studied EI Psc, since it is bright enough to be observed with TUGs 1.5 m and robotic telescopes, RTT150 and ROTSE, respectively. Its very short orbital period makes it possible to obtain more than two periods in one night with 1.5 m RTT150 telescope and long term light curve with ROTSE robotic telescope. The source is also observed in X-rays with ROSAT satellite with a rather low observation period. Detailed observational study of cataclysmic variables will give clues about the evolutionary scenarios.

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obtained by using IRAF aperture photometry package. We performed differential photometry by using nearby stars in the field. This technique is often used for studying the variations in the light curve of CVs. We used C1 (see Fig. 1) as a comparison star since its magnitude ðV  15:1 magÞ is close to the magnitude of the variable. The light curves of EI Psc for six night are shown in Fig. 2a–f. In Fig. 2g the light curves of variable (V), comparison (C1) and check star (C2) can be seen for one observation period to show the stability of the comparison star. Periodic variabilities in light curve was obtained by using Period04 time series analysis program (Lenz and Breger, 2005). This program is an extended version of Period98 by Sperl (1998) designed to search for and fit sinusoidal patterns within a time series of data in which one suspects periodic behavior. To extract all possible frequencies, this program performs Fourier analysis definition of

f ðtÞ ¼ Z þ

X

Ai sinð2pðXi t þ /i ÞÞ

ð1Þ

i

We observed EI Psc at TUG through the Robotic Telescope (ROTSE) for three and six months, the typical errors for these observations on the average are approximately 0.02 and 0.015 mag, respectively. Eight data points were obtained each night with clear filter. The time span between the data points was 30 min, see Fig. 3. The purpose of this observation was to obtain long term light curve of EI Psc, if exist, to find possible quasi periodic oscillations (QPO), outburst periods and duration of the outburst. Furthermore, spectral data of this source was obtained on July 26, 2006 with TUG’s low resolution spectrograph TFOSC equipped with 2048  2048 back illuminated camera with pixel size of 15l  15l. The spectra ranged between 6200 and 7850 Å with a resolution approximately 8 Å. Wavelength calibration of this data was performed by IRAF using standard method. The spectrum shows Ha and TiO bands as can be seen in Fig. 4. The X-ray data described in this paper are taken from the ROSAT Data Archive. The selected data are the ROSAT PSPC data

2. Observation Our photometric observations of EI Psc was carried out at TUG with 150 cm Russian–Turkish joint telescope between 2004 and 2006 for six nights about 2.35 h with V filter. The journal of observations are shown in Table 1. The data were obtained with the ANDOR CCD (2048  2048 pixels at 0.24 arcsec/pix resolution). The exposure time was 90 s for each frame and reduced using standard procedures. Bias, dark and flat field corrections were made for each frame. After cleaning the raw data, the V magnitude of the variable (15.7 mag) was

Table 1 The journal of observations for EI Psc UT date

Time of start (HJD-245300)

Run time (h)

Filter

Number of frames

040813 040814 040815 050924 060824 060825

231.5356 232.5092 233.4918 638.4760 972.372 973.292

1.67 2.35 2.3 2.5 2.16 3.11

V V V V V V

47 75 62 27 87 87

Fig. 1. The field of Ei Psc. The image indicates the comparison stars used for differential photometry. The field of view is 4.50  4.50 .

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Fig. 2. (a–f) Gives light curve of EI Psc in V filter for six nights. (g) Gives the light curves of variable (V), comparison (C1) and check star (C2), respectively, for one night.

which observed the source between November 29 and December 18, 1990. It was a RASS observation. The PSPC, which is a gas filled proportional counter in combination with the X-ray camera, is sensitive over the energy range of 0.1–2.4 keV. Its energy resolution is DE=E  44% at 0.93 keV. Detailed descriptions of the satellite, and PSPC detector can be found in early works of Trumper (1983) and Pfeffermann et al. (1986). The total exposure time was then 15,796 s. The shorter part of the observation (about 445 s) belongs to EI Pcs because this observation was made during ROSAT’s survey mode. The reduction of the ROSAT archival data was performed with the extended scientific analysis system (EXSAS) (Zimmermann et al., 1993).

3. Data analysis and results From the time series analysis especially prepared for finding variabilities in the optical light curve of variable stars, the orbital period of this source was determined to be 42.85 cycles per day and a second period, which may well be its double hump period, was approximately 24.48 cycles per day for two night (160) data points. For six night (450) data points these values were found to be 43.7 and 24.5 cycles per day as can be seen from Fig. 5a and b. The difference between these values may result from aliasing effect due to time gaps in our data set. Fig. 2a, b and d shows the light curve of the system at its quiescent state and on average shows 0:23  0:01 mag variations.

F. Gök et al. / New Astronomy 14 (2009) 160–165

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Fig. 3. The light curve of EI Psc obtained by ROTSE for (a) three months in 2004 September–2005 January (33 nights) and (b) for six months in 2005 August–2006 January (51 nights).

Fig. 4. Normalized flux versus wavelength spectrum of EI Psc.

Fig. 2c, e and f shows the system as it is probably approaching or just after the decline of outburst with a 0:30  0:02 mag variations. Early studies of the source reveals that the source is a SU Uma type CV. SU UMa type CVs undergo super outburst followed by succession of regular outbursts, and then another super outburst. Fig. 3a shows its three-month light curve while Fig. 3b shows its sixmonth light curve with clear filter of ROTSE. Our purpose was also to obtain QPOs if they existed (except the reported ones) but unfortunately, we could not reach such an opportunity because of the bad weather conditions. Therefore the gaps in light curves did not allow us to observe it. We applied period analysis program on these two data set, without any success. In the light curves however, one can see about 2.5 mag normal outburst. The recurrence time and duration of the outbursts unfortunately is not clear in the data. Fig. 4 gives the optical spectroscopy of EI Psc. It can be seen from the figure that at values of k > 7000 Å noise increases, as the star is near the faint end for the RTT150. Nevertheless at k < 7000 Å, we can still identify important lines like Ha, HeI in emission and TiO absorption lines expected at 6159 and 6178 Å cannot be identified separately because of the noise (the signal to noise ratio is 5). It is even more difficult to identify the TiO lines at 6322 Å, 6559 Å, 6651 Å,7053 Å in this spectrum since they are obscured by noise (the value of noise is 0.25 flux unit for k < 7000 Å and 0.75 flux unit for k > 7000 Å). TiO absorb-

Fig. 5. (a) The power spectrum of 171 data points for two nights, for the light curves given in Fig. 1a and b in V filter. (b) The power spectrum of 403 data points for six nights, for the data points given in Fig. 1(a–f), in V filter.

tion lines appearing in the quiescent spectrum imply that the spectral type of the secondary is M-type dwarf. The emission of Ha line implies the presence of a disc (Hellier, 2001). In this study we used the soft X-ray archival data of the EI Psc taken from the ROSAT. The count-rate of the system is 0:185 cts=s. Because there are only seven data points in the ROSAT observation of EI Psc, we could not bin the data. The observed count-rate value is very low because of the poor number of data points. Since there is no X-ray analysis of the system in the literature, in spite of this low count-rate value, we wanted to see what could be obtained as the result of spectral analysis of these X-ray

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data. Various spectral models were fitted for this X-ray data set. The best fit models were found to be the blackbody and the Raymond–Smith models. The statistical values were approximately the same for the same degrees of freedom ðm ¼ 5Þ: v2m ¼ 0:76 for Blackbody model and v2m ¼ 0:74 for Raymond–Smith model. Because of the poor statistics we could not resolve the best fitted model in these two models. Instead, we used both of them and their parameters in our calculations. Their spectrum can be seen in the Fig. 6 and parameters are given in the Table 2. Both models are thermal models and give ðð0:07 0:02Þ keV ð0:80  0:23Þ  106 KÞ and ðð0:13  0:04Þ keV ð1:50 0:46Þ  106 KÞ, respectively. The Galactic absorption value calculated using Colden Galactic Neutral Hydrogen Density Calculator in Chandra X-ray center (http://cxc.harvard.edu/toolkit/colden.jsp) with the given right ascension and the declination of the system and used in all analysis of this study to be 0:54  1021 cm2 . Energy flux values are 6:74  1013 ergcm2 s1 for blackbody and 2:96  1014 ergcm2 s1 for Raymond–Smith model. Uemura et al. (2002a) give an interval for the distance of the EI PSc to be 245–380 pc. Using this interval and the flux energies from our study, the model dependent luminosity interval can be found as 4:84  1030 —1:16  1031 erg1 for blackbody model and 2:13 1029 —5:11  1029 erg1 for Raymond–Smith model, respectively. Since the bulk of soft X-ray emission of a CV with an accretion disk comes from boundary layer (Warner, 1995), we expect the Xray photons of the source may have been well originated from the boundary layer of the system with a high probability. Mauche (2002) has given an equation for the blackbody temperature of the boundary layer if a boundary layer emits with blackbody mechanism

_ prR3 1=4 T bb ¼ ½GM1 M=8 1

ð2Þ

where G is the gravitational constant, M1 is the mass of the white _ is the mass accretion rate, r is the Stephan–Boltzmann dwarf, M constant and R1 is the radius of the white dwarf (Mauche, 2002). If the soft X-ray photons of the source comes from boundary layer as the theories suggest, this equation can be used for the calculation

of mass accretion rate of EI Psc. The blackbody temperature obtained with our spectral analysis is ð0:80  0:23Þ  106 K and M1 value is given by Uemura et al. (2002b) to be 0:7M  . However, the radius of the white dwarf is also needed in this equation. Since no value is reported for R1 of EI Psc in the literature, it is calculated by using the following equation taken from Patterson and Raymond (1985) and references therein

R1 =R ¼ 0:007ðM 1 =M  Þ0:8

The radius of the primary component is calculated to be approximately 6:48  103 km for EI Psc. Lyndel-Bell and O’Dwyer (2008) have given a relation for the radius of the Planets, White dwarfs and Neutron stars. Using this relation the authors have drawn a graph (Fig. 1 in Lyndel-Bell and O’Dwyer (2008)) between masses of the objects and the radius values, and the graph shows that the radius value must be in an interval between 6  103 and 7  103 km for the white dwarfs which have 0:7 M  ð1:39  1033 gÞ masses. This comparison gives us approximately the same value for the radius of the white dwarf of EI Psc (6:48  103 km as found from Eq. (3)). Using this radius value, the temperature value calculated in this study and the Eq. (2) the mass accretion rate is calcu_ ¼ ð1:71  0:012Þ  1021 gs1 for the boundary layer lated to be M of EI Psc. For soft X-ray blackbody photons there is another formulae that has been reported by Patterson and Raymond (1985). It is as follows

_ 0:18 T bl ¼ 2:16  105 M 0:86 0:7 M 18 K

Table 2 The result of the free parameter analysis of EI Psc for Black body model, normalization amplitude for Raymond–Smith in cm5 Models Parameters

Black body

Raymond–Smith model

T (keV) A N H ð1021 cm2 Þ Energy flux

0:07  0:02 ð4:6  2:8Þ  104 0.54 6:74  1013 0.76

0:13  0:04 ð1:62  1:10Þ  104 0.54 2:96  1014 0.74

v2m

T is radiation temperature for black body model, plasma temperature for Raymond– Smith model. Energy fluxes are given in erg cm2 s1 .

ð4Þ

_ is the where M0:7 is the mass of the white dwarf to be ðM1 =M Þ, M 18 _ Þ. Using the Eq. (4), the temperamass accretion rate to be ðM=10 ture value calculated in this study and the mass of the primary star _ can be calculated to be to be 0:7M  (Uemura et al., 2002b) M ð1:44  0:0014Þ  1021 gs1 for the boundary layer of EI Psc. Thus, the mean of these mass accretion values is found to be 1:575  1021 gs1 , so it can be adopted for the mass accretion rate of EI Psc to be approximately ð1:58  0:14Þ  1021 gs1 by using blackbody radiation temperature and taking previous two mass accretion values as the limits. Pringle (1977) has reported that for the soft X-ray boundary layer photons, which are originated by different thermal mechanisms from blackbody, there is a relation between the thermal plasma temperature, radius and the mass accretion that

_ 6=19 T bl ¼ 8:4  105 M 18

Fig. 6. The best spectral fits.

ð3Þ



M M

8=19

ðR8:7 Þ18=19 K

ð5Þ

_ value where R8:7 equals R1 =5  108 cm. Thus we can calculate the M using this Eq. (5) and temperature value obtained in this study from the best fitted Raymond–Smith model of ð1:5  0:46Þ  106 K. As a result, the mass accretion rate is obtained to be ð2:2  0:052Þ 1019 gs1 for EI Psc. It is much lower than the value that is calculated using blackbody radiation temperature. Since we obtained these temperatures and mass accretion rate values from our spectral analysis which depend upon only seven X-ray data points, all these values should be accepted as approximate values. The simple theoretical models suggest that the mass accretion _ ¼ 2  1016 gs1 ð3:17  1010 M  yr1 Þ is the critical value rate of M for the systems which have an accretion disks (Pringle and Savonije, 1979; Narayan and Popham, 1993; Patterson and Raymond, 1985 and Warner, 1995). When the mass accretion rate is larger than or equal to the critical value, the boundary layer becomes optically thick and as a result, only the thermalized soft X-rays (<10 KeV) are emitted (Silber et al., 1994). The temperature values obtained in this study ðð0:07  0:02Þ—ð0:13  0:04Þ keVÞ

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are less than 10 keV and the mass accretion rate values ðð1:58  0:14Þ  1021 —ð2:2  0:052Þ  1019 g s1 Þ are greater than the critical value. It can be said that EI Psc has an optically thick boundary layer. Uemura et al. (2002a,b) and Mennickent et al. (2004) have reported in their optical spectroscopic observations that EI Psc must have a relatively high mass accretion rate from secondary star to the accretion disc even in the quiescent state. Our X-ray results also indicate a similar trend for this source in the boundary layer. Because of the high mass accretion rate values that we have obtained in this work for the boundary layer of EI Psc, it can be concluded that the mass accretion rate from secondary star to the accretion disk must have a high value too. The main source of the mass accretion in the boundary layer is the mass accretion between the accretion disk and secondary star in the system if there is not any structure which blocks the mass stream in the accretion disk. Due to the high mass accretion on the accretion disk, the secondary component of the system becomes less and less massive along the accretion process. According to the current theories about CVs with hydrogen rich and less massive secondaries, the secondary can be eroded and lose most of its hydrogen depending on the high mass accretion rate. As a result, the orbital periods of the systems get shorter and shorter. The CVs with orbital periods shorter than the observed period minimum (78 min) will able to have Browns Dwarfs or start their life with a Brown Dwarf secondaries in their evolutions. The orbital periods for such systems can be as short as 46 min like V485 Cen (Pollitano, 2002, Mennickent et al., 2004 and references therein). As a result of this study we concluded from our optical analysis that EI Psc can have M-type secondary star and 58.75 min orbital period, and from X-ray analysis that EI Psc has very high mass accretion rate in the boundary layer. Depending on these results and the current theories on evolution of the shortest period CVs mentioned in the previous paragraph, we can conclude that EI Psc has a less massive secondary star which can be M-type star because there is TiO absorption band in the optical spectrum of EI Psc (one of the most important indicator of M-type stars) (Bessel, 1991). Because of the high mass accretion and the short orbital period in the system, according to the current theories of shortest orbital period CVs, the secondary star of EI Psc can be a Brown Dwarf. If it is a Brown Dwarf, the type of it must be late M-type. Because generally the late M-type stars hold Brown Dwarf stars (Bessel, 1991). Thorstensen et al. (2002) and Mennickent et al. (2004) reported from their optical analysis that EI Psc must have a K-type secondary. However, the last work have reported that the effective temperature they obtained is incompatible with a K-type star. Their temperature value is more suitable for M (as we resulted) or L-type stars. 4. Conclusion We report here an optical orbital period for EI Psc as 0.0408 days and a double humped wave with 0.023 days which is almost half the orbital period. This double hump period is thought to be arising from the ellipsoidal distortion of the secondary. The quiescent spectrum of EI Psc shows TiO bands which indicate that the companion has a spectrum like M-type. Furthermore, we determined the X-ray spectral parameters for EI Psc using ROSAT RASS archival data. The raw data were fitted to various spectral models

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and the best fit models were found to be that of Blackbody and Raymond–Smith. The best fit temperatures were estimated to be kT ¼ ð0:07  0:02Þ keV for blackbody model and kT ¼ ð0:13 0:04Þ keV for Raymond–Smith model. The calculated column density is 0:54  1021 cm2 . The estimated flux is in the range of F ¼ 1013 —1014 erg cm2 s1 . The model dependent luminosity values are found to be 4:84  1030 —1:16  1031 ergs1 for blackbody model and 2:13  1029 —5:1  1029 erg1 for Raymond–Smith model. Using the obtained temperature values and the mass of the primary value that is reported by Uemura et al. (2002b), mass accretion rate in the boundary layer was reported to be ð1:58  0:14Þ  1021 gs1 for the blackbody model and ð2:2 0:052Þ  1019 gs1 for Raymond–Smith model. As a result of this study we can argue that the system EI Psc has very high mass accretion rate ð1019 —1021 gs1 Þ, and because of the observed soft X-ray photons with ð0:07  0:02Þ—ð0:13  0:04Þ keV temperature and high mass accretion rate, it has an optically thick boundary layer. Optical and X-ray analysis in this study show that EI Psc can have a M-type secondary star which can be a Brown Dwarf. More detailed, long duration X-ray and optical observations and analysis must be performed to investigate the detailed emission mechanisms, boundary layer structure, period analysis and the overall system structure of the EI Psc. References Bessel, M.S., 1991. AJ. 101 (2), 662. Hellier, C., 2001. Cataclysmic Variable Stars. Springer-Verlag, Berlin. Hu, J., Wei, J., Xu, D., Cao, L., Dong, X., 1998. Ann. Shanghai Obs. Acad. Sin. 19, 235. Jingyao et al., 1998. Ann. Shanghai Obs. 19, 235. Lenz, P., Breger, M., 2005. Commun. Asteroseismol., 146. Lyndel-Bell, D., O’Dwyer, J.P., 2008. One Mass-Radius Relation form Planets, White Dawrfs and Neutron Stars. eprint . Mauche, C.W., 2002. EUVE observations of nonmagnetic cataclysmic variables. In: ASP Conference Series’ Continuing the Challenge of EUV Astronomy: Current Analysis and Prospects for the Future, 264, 75M. Mennickent, R.E., Diaz, M.P., Tappert, C., 2004. MNRAS 347, 1180M. Narayan, R., Popham, R., 1993. Nature 362, 820. Patterson, J., 2001. PASP 113, 736. Patterson, J., Raymond, J.C., 1985. AJ. 292, 550. Pfeffermann, E., Briel, U.G., Hippmann, H., Kettenring, G., Metzner, G., Predehl, P., Reger, G., Stephan, K.H., Zombeck, M.V., Chappell, J., Murray, S.S., 1986. SPIE 733, 519. Pollitano, M., 2002. ASp Conference Series 261, 293. Pringle, J.E., 1977. MNRAS 178, 195. Pringle, J.E., Savonije, G.J., 1979. MNRAS 187, 777. Schmeer, P., 2001. vsnet-alert6830 . Silber, A., Vrtilek, S.D., Raymond, J.C., 1994. AJ. 425, 829. Sperl, M., 1998. CoAst 111, 1. Thorstensen, J.R., Fenton, W.H., Patterson, J.O., Kemp, J., Krajci, T., Baraffe, I., 2002. ApJ. 567, L49. Trumper, J., 1983. Adv. Space Res. 2, 241. Uemura, M., Ishioka, R., Kato, T., 2001. IAU Circ., 7747. Uemura, M., Kato, T., Ishioka, R., Yamaoka, H., Scgmeer, P., Starkey, D.R., Torii, K., Kawai, N., Urata, Y., Kohama, M., Yoshida, A., Ayani, K., Kawabata, T., Tanabe, K., Matsumoto, K., Kiyota, S., Pietz, J., Vanmunster, T., Krajci, T., Oksanen, A., Giambersio, A., 2002a. PASJ 54, L15. Uemura, M., Kato, T., Ishioka, R., Yamaoka, H., Scgmeer, P., Starkey, D.R., Torii, K., Kawai, N., Urata, Y., Kohama, M., Yoshida, A., Ayani, K., Kawabata, T., Tanabe, K., Matsumoto, K., Kiyota, S., Pietz, J., Vanmunster, T., Krajci, T., Oksanen, A., Giambersio, A., 2002b. PASJ 54, 599. Voges, W., Aschenbach, B., Boller, T., Brauninger, H., Briel, U., Burkert, W., 1996. IAU Circ., 6420. Warner, B., 1995. Cataclysmic Variable Stars. Cambridge University Press, UK. Wei, J.Y., Xu, D.W., Dong, X.Y., Hu, J.Y., 1999. A&Astrophys. 139 (Suppl.), 575. Zimmermann, H.U., Belloni, T., Izzo, C., Kahabka, P., Schwentker, O., 1993. EXSAS User’s Guide, MPE Report 244.