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tev @xmath0-ray binaries are a subclass of binaries with a compact object , established in the 2000 s ( e.g. , * ? ? ? * for a recent review ) . these systems have a spectral energy distribution with a peak beyond 1 mev and are variable at multi - wavelengths up to tev energies . there are 5 known binaries of this kind , all of which have either an o - type main - sequence star or a be star with a circumstellar disk as the optical counterpart . the nature of the compact object is unknown in all systems but one ( ) . for such systems two competing scenarios have been proposed , based on the different nature of the compact object , and hence the different mechanism of high energy emissions . one scenario assumes that the collision between a relativistic pulsar wind and a stellar wind and/or circumstellar disk produces strong shocks , where the high energy emission arises ( e.g. , * ? ? ? this pulsar wind scenario has successfully been applied to . on the other hand , the other scenario assumes the presence of an accreting black hole ( or neutron star ) . in this microquasar or accretion / ejection scenario , a large amount of mass transferred from a companion star powers relativistic jets , where @xmath0-ray emission originates ( e.g. , * ? ? ? * ) . ( @xmath6 ) is a recently established tev @xmath0-ray binary @xcite comprised of a b0ve star and a compact object of unknown nature . the orbit is wide and highly eccentric ( @xmath7 ; * ? ? ? the system shows correlated variability in the x - ray and tev energy bands @xcite , with two peaks in one orbital cycle : the primary outburst prior to apastron ( orbital phase @xmath8 = 0.30.4 ) and the secondary outburst after it ( @xmath9 0.60.9 ) , with an x - ray dip in between . this is puzzling , given that these outbursts and the dip occur when the compact object is orbiting very far from the be star . understanding the mechanism for these phenomena is one of the key issues of this system . the optical counterpart of , , is established to be a classical be star , i.e. , a massive star with a geometrically thin , circumstellar disk , where balmer and other emission lines originate @xcite . @xcite observed the h@xmath2 emission line after apastron ( @xmath9 0.70.8 ) . their profiles over continuous 35 days showed an `` s - shaped '' variability with a period of @xmath4 60 days . @xcite reported the orbital modulation in the h@xmath2 line profile parameters , likely caused by the interaction between the be disk and the compact object . however , the lack of observations in @xmath8=0.60.8 makes it difficult to constrain the modulation after apastron .    in order to constrain the interaction between the be disk and the compact object , we have been monitoring  using various methods such as high - dispersion spectroscopy , photometry , and polarimetry . in this paper we report the initial result of the line profile variabilities over 160 days from 2013 october to 2014 april , covering the second half of the orbital period . in sect . 2 , we summarize our observations . the results are described in sect . 3 . in sect . 4 we discuss the be disk region affected by the compact object and the nature of the interaction . optical high - dispersion spectroscopic observations of  were carried out on 12 nights from 2013 october 31 ( phase 0.555 ) to 2014 march 11 ( phase 0.969 ) at the okayama astrophysical observatory with a 188 cm telescope and hides with fiber - feed system @xcite . here , we calculated the orbital phase , @xmath8 , taking the orbital period of 315 days from @xcite and the origin ( jd 2454857.5 ) from @xcite , and setting the periastron phase to be 0 ( in @xcite , it is set to 0.967 ) . note that , although @xcite determined the ephemeris using the orbital period of 321 days , the orbital parameters remain approximately unchanged @xcite . the observed spectra covered 42007400   wavelength range , with s / n of @xmath10 around h@xmath2 . the typical wavelength resolution @xmath11 is @xmath12 . the data were reduced in the standard way , using the iraf echelle package . high - dispersion spectra were also taken on two nights ( 2014 february 05 , phase 0.861 , and april 10 , phase 0.064 ) using the canada france hawaii telescope / espadons @xcite in spectropolarimetric mode . the spectra covered a wavelength range of 370010500   with a resolving power of @xmath13 . we obtained reduced data using the libre - esprit / upena pipeline , provided by the instrument team . we rectified the normalized intensity by re - determining the continuum level around each line , in order to compare with hides data . in this paper , we focus on the spectroscopic variability . the polarimetric data will be discussed in a forthcoming paper .    in addition to spectroscopic data , @xmath14- , @xmath15- , and @xmath16-band photometric data were obtained from 2013 september to 2014 may using howpol @xcite and honir @xcite attached to kanata 1.5 m telescope at higashi - hiroshima observatory in hiroshima , japan . we observed  on 73 , 71 and 10 nights in @xmath14 , @xmath15 , and @xmath16 bands , respectively , using howpol . we used honir to observe the source on 8 nights , in these three filters . the logs of the photometric and spectroscopic observations are given in table [ tbl : log ] . figure [ fig : profave ] shows the averaged h@xmath2 , h@xmath3 , and h@xmath0 profiles . in these profiles , the vertical scale of the h@xmath2 line profile is half that of the other line profiles , because the flux is much stronger in the h@xmath2 . the h@xmath2 line is very strong ( @xmath17 ) and significantly asymmetric . the profile exhibits a double peak at @xmath18 and @xmath19 , with the latter being brighter than the former , and a hump at @xmath20 on the blue wing . on the other hand , the averaged h@xmath3 line profile ( @xmath21 ) is rather symmetric , with a slightly stronger blue peak . the h@xmath0 line , which is on a broad absorption component , also exhibits a double - peaked profile with stronger blue peak . balmer lines exhibited complicated line profile variabilities during the monitoring period . figure [ fig : prof_ind ] displays all observed profiles of the h@xmath2 , h@xmath3 , and h@xmath0 lines , while fig . [ fig : prof_diff ] presents the time sequence of the residual spectra of these lines from the average . the variation in the equivalent width , ew , the full width at half maximum , fwhm , and the centroid velocity , @xmath22 , of the h@xmath2 line profile are shown in fig . [ fig : prof_param_ha ] . note that fwhm and @xmath22 are measured by fitting the whole profile with a single gaussian and the line wings with a voigt profile , respectively .    in the long term ( @xmath23 days , from @xmath24 ) , the h@xmath2 line changed from a red - enhanced profile to a rather symmetric profile , while the stronger peak of the h@xmath3 and h@xmath0 lines changed from the red side to the blue side . in the residual spectra of these lines , s - shaped variability is visible ; the bright part at @xmath25 migrates from the blue side to the red side at @xmath26 , and returns to the blue side at @xmath27 . simultaneously , the faint part at @xmath28 migrates in the opposite direction . at the same time , the maximum and minimum , respectively at @xmath29 and @xmath30 in the residual spectra , migrate in a similar fashion . the period of these variations is estimated at @xmath31 days using fourier analysis , which is about twice the @xmath32-day variation period reported by @xcite . we will discuss this in more detail in the next section . on the other hand , in the short term ( @xmath33 days ) , a bright hump appears at @xmath29 in balmer line profiles at @xmath34 . this hump disappears by @xmath35 ( see the third , forth , and fifth profiles from the bottom in figs . [ fig : prof_ind ] and [ fig : prof_diff ] ) . when the hump appears , ew(h@xmath2 ) increases by @xmath36 and fwhm(h@xmath2 ) decreases by @xmath37 . these long- and short - term variabilities respectively seem to occur in phase in different lines , although time lags between lines can not be ruled out because of the low cadence of the observations . recently , @xcite reported the orbital modulation of the profile parameters of the h@xmath2 line . their data cover the phase intervals of 00.6 and 0.81 , using the orbital period of 321 days @xcite . if we take the same orbital period , our data spans the orbital phase 0.41 , which overlaps with @xcite in the phase range 0.40.6 and 0.81 and provides for the first time the modulation data in the phase interval of 0.60.8 . in fig . [ fig : prof_param_ha ] , ew(h@xmath2 ) seems to increase ( phase 0.40.5 ) and then decrease back to the previous value in about 30 days ( phase 0.50.6 ) , as reported by @xcite . afterwards , it seems to fluctuate for about 60 days ( phase 0.60.8 ) , which could be associated with the variabilities mentioned above . @xmath22(h@xmath2 ) , whose variation is thought to reflect the orbital motion of the be star , shows a similar pattern to @xcite except just after the apastron and just before the periastron . this difference in @xmath22 is possibly caused by the profile variabilities . fwhm(h@xmath2 ) , on the other hand , shows a different variation pattern from @xcite . it stays constant except around apastron ( phase 0.450.55 for the period of 321 days ) in this work , whereas in @xcite it showed a sinusoidal pattern . although many lines are contaminated by neighboring lines , there are six less - contaminated lines ( @xmath385018 , 5316 , 5363 , 5535 , 6433 , and 6456 ) that enable us to analyze the line profile variabilities . in our analysis , we have used these six lines . the averaged @xmath395363 line profile in fig . [ fig : profave ] is shown as a representative profile . in contrast to the balmer lines , the emission lines show symmetric double - peaked profiles , as predicted for a rotating , axi - symmetric disk . figures [ fig : prof_ind ] and [ fig : prof_diff ] also display the time sequence of the observed and residual @xmath395363 profiles , respectively . as seen in these figures , lines exhibited no variation . @xcite analyzed the spectra of several be stars and confirmed that emission lines arise from a disk region of radius of @xmath40 . they also found empirical relationships between the projected rotational velocity of the central star ( @xmath41 ) and the profile parameters . applying their relationships , we have derived @xmath41 to be @xmath42 . finally , we estimate the radii of emitting regions of balmer lines . for the h@xmath2 line , we use the mean peak separation of the last six profiles ( @xmath43 ) , because the averaged profile in fig . [ fig : profave ] , in which all our observed epochs were used , has too complicated features to determine the peak velocities . for the h@xmath3 and h@xmath0 lines , we use the averaged profiles . as a result , we have the peak separations of the h@xmath2 , h@xmath3 , and h@xmath0 lines to be @xmath44 , @xmath45 and @xmath46 , respectively . adopting @xmath41 derived above , we obtain the radii of the emitting region of these lines as @xmath47 @xmath5 ( h@xmath2 ) , 11 @xmath5 ( h@xmath3 ) , and 7 @xmath5 ( h@xmath0 ) , where @xmath48 is the radius of the be star .    throughout our monitoring period ( from apastron to periastron ) , optical brightness of  stayed constant in the range of 9.059.32 ( @xmath14 ) mag , 8.538.71 ( @xmath49 ) mag , and 8.588.67 ( @xmath50 ) mag , within the typical 1-@xmath51 error of 0.05 mag . the average values are @xmath52 mag , @xmath53 mag and @xmath54 mag . although the monitoring was performed mostly after apastron , the line profiles exhibited remarkable variabilities . wide wavelength coverage revealed that the variations are seen not only in the h@xmath2 line but also in the h@xmath3 and h@xmath0 lines . no variation in lines indicates the inner part of the be disk was kept undisturbed during the observing period . this fact is in agreement with little variability in the optical brightness , which is thought to originate from the disk region within 23 @xmath5 @xcite . the slight , but significant , variation in the h@xmath0 line profile implies that the interaction with the compact object affects the be disk down to , at least , a radius of 7 @xmath5 . the dynamical residual spectra of the balmer lines showed an s - shaped variation ( fig . [ fig : prof_diff ] ) , as @xcite reported . high - dispersion residual spectra clearly show that there are two pairs of peaks at @xmath55 and @xmath56 . the lower velocity peaks are the same as those of @xcite ( see their fig . the period of the variation is @xmath4150 days , which is about twice the @xmath4 60-day variation period @xcite . the s - shaped variation seen in many be stars is thought to be caused by global disk oscillations . they are mostly low - frequency @xmath57 oscillations , where @xmath58 is the azimuthal wave number @xcite , but in eccentric binaries , it is also possible that @xmath59 oscillations are excited by the corresponding fourier components of the tidal potential ( e.g. , * ? ? ? if the period of observed variation is @xmath4 150 days , about half the orbital period , it might be due to an @xmath60 oscillation mode in the be disk , which could be excited by the @xmath60 fourier component of the tidal potential of the compact object . the difference of the periods between this work and @xcite might be caused by the difference of the length and frequency of the observations . because we observed  for about half an orbital cycle ( @xmath61 days ) , variabilities of this timescale can be detected . with low cadence of observation , typically once every 1020 days , however , profile variabilities of shorter timescales are invisible . on the contrary , the short monitoring epoch of @xcite ( 35 days ) would have hidden the variability with longer timescales ( @xmath62 days ) . moreover , observations of this work and @xcite have an interval of @xmath41800 days . therefore , there are possibilities that the oscillation mode has changed or that both oscillation modes exist . in addition to the s - shaped variation , a remarkable variability was seen in the balmer lines after apastron ( @xmath8=0.600.65 ) . the short lifetime of this variability ( @xmath63 days ) indicates that it is temporarily caused by an external force such as the tidal force and the ram pressure of the pulsar wind , because internal waves / oscillations have much longer lifetimes ( @xmath64 days ; * ? ? ? * and references therein ) . it is surprising , however , that such a remarkable variability appeared after apastron , when the interaction is expected to be weak because of the large distance between the be disk and the compact object . in order to understand this phenomenon , it is essential to constrain the starting phase and repeatability of the variability . comparison of the ew(h@xmath2 ) between this work and @xcite implies the presence of a regular orbital modulation , but different absolute values due to different spectral resolutions make it difficult to quantitatively analyze it . given the orbital period close to one year , it is necessary to monitor over two or three more successive cycles , within the timescale of be disk variability ( @xmath65 days ) , in order to cover the full orbital phase with the same spectral resolution . because the compact object is located far away from the be star around apastron ( @xmath66 ) , the gas in the be disk at the radius of 7 @xmath5 can not be significantly affected by the tidal force of the compact object , whose strength is less than @xmath67 of the gravity of the be central star at this radius . this leaves the interaction with the pulsar wind as the most likely mechanism for the balmer line variability .    at the periastron passage , on the other hand , the compact object is thought to pass through the be disk at @xmath68 . nonetheless , the h@xmath3 line , emitted from the disk region of the similar radius , showed no remarkable variation at @xmath690.06 ( see the top profile in the upper right panel of fig . [ fig : prof_ind ] ) . this fact suggests that either the tidal interaction is very weak in this system or the orbital period is longer than 315 days , so that the compact object had not passed the disk yet at the time of observation . the former possibility could be realized if the be disk rotates in the retrograde direction or is tilted by a large angle with respect to the orbital plane , where the disk gas and the compact object interact on a very short timescale . future observations at significantly later phases will distinguish between these two possibilities . as described in sects . [ sec : short ] and [ sec : imp ] , our observations leave the pulsar wind model as the sole candidate for . if the compact object is a pulsar , little line profile variation around periastron ( the top two profiles in each panel of fig . [ fig : prof_diff ] ) suggests that the pulsar wind is too weak to significantly affect the be disk during these phases . @xcite proposed a flip - flop scenario for another gamma - ray binary , , where the pulsar is in a rotationally powered regime in the apastron , while it is in a propeller regime in the periastron ( see also * ? ? ? * ) . in a flip - flop system , if the gas pressure of the be disk overcomes the pulsar - wind ram pressure , the pulsar wind is quenched . because the be disk of is estimated to be about three times larger than the binary separation at periastron , the compact object crosses a dense region of the disk near the periastron . in such a situation , the strong gas pressure is likely to quench the pulsar wind and hence suppress high - energy emissions . in the framework of the pulsar wind model , there are a few mechanisms that might explain the short - term episodic variability discussed above . the be disk is likely as large as the binary orbit because of no tidal truncation in highly eccentric , large orbit @xcite . because the disk density rapidly decreases with radius ( e.g. , * ? ? ? * ) , the wind from the pulsar close to the outer part of the disk effectively changes its structure , giving rise to remarkable variability in emission lines arising from the disk outer part . if the be disk is misaligned with the orbital plane and a node happens to be in the direction corresponding to the phase of the variation ( @xmath70 ) , the pulsar - wind effect on the be disk will be the strongest in this phase interval . after @xmath71 the pulsar comes close to the denser disk region , which is not easily affected by the pulsar wind . this might be a cause of the short - term episodic variation after aspastron . it is not clear , however , how the pulsar wind can affect the inner disk at the radius of 7 @xmath5 . alternatively , the short - term post - apastron variations might be explained as the emission from the gas captured by the pulsar , if the pulsar wind is not strong enough to expel the surrounding gas . this picture seems to fit well with the flip - flop model . at any rate , observational investigation in the first half of the orbital cycle is needed in order to further test the pulsar - wind scenario . the next periastron passage of  will take place in 2015 december , according to the ephemeris of @xcite ( 2016 january if the ephemeris is taken from @xcite ) . observations covering this period will provide more clues to the complex interaction and the nature of the compact object in this puzzling @xmath0-ray binary . this paper is based on the observations taken at the okayama astrophysical observatory . we are grateful for dr . eiji kambe to kindly observe . this work was supported by research fellowships for the promotion of science for young scientists ( ym , kt ) . ato acknowledges support by the jsps grant - in - aid for scientific research ( 24540235 ) and a research grant from hokkai - gakuen educational foundation . acc acknowledges support from cnpq ( grant 307076/2012 - 1 ) . aharonian , f. a. et al . 2007 , , 469 , l1 aliu , e. , archambault , s. , aune , t. et al . 2014 , , 780 , 168 aragona , c. , mcswain , m. v. & de becker , m. 2010 , , 724 , 306 arias , m. l. , zorec , j. , cidale , l. , ringuelet , a. e. , morrell , n. i. & ballereau , d. 2006 , , 460 , 821 artymowicz , p. , & lubow , s. h. 1994 , , 421 , 651 bongiorno , s. d. , falcone , a. d. , stroh , m. , holder , j. , skilton , j. l. , hinton , j. a. , gehrels , n. & grube , j. 2011 , , 737 , 11 carciofi , a. c. & bjorkman , j. e. 2006 , , 639 , 1081 casares , j. , rib , m. , ribas , i. , paredes , j. m. , vilardell , f. & negueruela , i. 2012 , , 421 , 1103 dubus , g. , 2013 , , 21 , 64 kambe , e. , yoshida , m. , izumiura , h. at al . 2013 , 65 , 15 kawabata , k. s. , nagae , o. , chiyonobu , s. , et al . 2008 , proc . spie , 7014 , 70144l levinson , a. & blandford , r. 1996 , , 456 , l29 manset , n. & donati , j .- f . 2003 , proc . spie , 4843 , 425 maraschi , l. & treves , a. 1981 , , 194 , 1 okazaki , a. t. 1991 , , 43 , 75 okazaki , a. t. & negueruela , i. 2001 , , 377 , 161 papaloizou , j. c. , savonije , g. j. & henrichs , h. f. 1992 , , 265 , l45 papitto , a. , torres , d. f. & rea , n. 2012 , , 756 , 188 rivinius , th . , carciofi , a. c. , & martayan , c. 2013 , , 21 , 69 sakimoto , k. , akitaya , h. , yamashita , t. et al . 2012 , proc . spie , 8446 , 844673 torres , d. f. , rea , n. , esposito , p. , li , j. , chen , y. , zhang , s. 2012 , , 744 , 106    lccc + & date & jd & phase + oao / hides + & 2013.10.31 & 2456597.278 & 0.555 + & 2013.11.08 & 2456605.275 & 0.580 + & 2013.11.11 & 2456608.274 & 0.589 + & 2013.11.19 & 2456616.199 & 0.615 + & 2013.12.01 & 2456628.219 & 0.653 + & 2013.12.19 & 2456646.201 & 0.710 + & 2013.12.29 & 2456656.007 & 0.741 + & 2014.01.02 & 2456660.047 & 0.754 + & 2014.01.15 & 2456673.095 & 0.795 + & 2014.01.23 & 2456681.094 & 0.821 + & 2014.02.08 & 2456696.968 & 0.871 + & 2014.03.11 & 2456728.008 & 0.969 + cfht / espadons + & 2014.02.05 & 2456693.843 & 0.861 + & 2014.04.10 & 2456757.772 & 0.064 +   + & filter & jd & nights + kanata / howpol + & @xmath14 & 24565512456795 & 73 + & @xmath15 & 24565512456795 & 71 + & @xmath16 & 24567222456795 & 10 + kanata / honir + & @xmath14 , @xmath15 , @xmath16 & 24567432456769 & 8 +
we report on monitoring observations of the tev @xmath0-ray binary , which were carried out to constrain the interaction between the be circumstellar disk and the compact object of unknown nature , and provide for the first time high - dispersion ( r@xmath1 ) optical spectra in the second half of the orbital cycle , from apastron through periastron . 
 the h@xmath2 , h@xmath3 , and h@xmath0 line profiles are found to exhibit remarkable short - term variability for @xmath41 month after the apastron ( phase 0.60.7 ) , whereas they show little variation near the periastron . 
 these emission lines show `` s - shaped '' variations with timescale of @xmath4150 days , which is about twice that reported previously . 
 in contrast to the balmer lines , no profile variability is seen in any emission line . 
 we estimate the radii of emitting regions of the h@xmath2 , h@xmath3 , h@xmath0 , and emission lines to be @xmath430 , 11 , 7 , and 2 stellar radii ( @xmath5 ) , respectively . 
 the amplitudes of the line profile variations in different lines indicate that the interaction with the compact object affects the be disk down to , at least , the radius of 7 @xmath5 after the apastron . 
 this fact , together with little profile variability near the periastron , rules out the tidal force as the major cause of disk variability . although this leaves the pulsar wind as the most likely candidate mechanism for disk variations , understanding the details of the interaction , particularly the mechanism for causing a large - scale disk disturbance after the apastron , remains an open question .