the apparent synchronization of v1500...

The Apparent Synchronization of V1500 Cygni 1

The Apparent Synchronization of V1500 Cygni1

Thomas E. Harrison3? Ryan K. Campbell4†3Department of Astronomy, New Mexico State University, Box 30001, MSC 4500, Las Cruces, NM 88003-80014Department of Physics & Astronomy, Humboldt State University, 1 Harpst S t., Arcata, CA 95521

Accepted XXX. Received YYY; in original form ZZZ

ABSTRACTWe have used XMM − Newton and ground-based optical/near-infrared photometry to explorethe old classical nova, and asynchronous polar V1500 Cyg. The X-ray light curve shows a singlebright phase once per orbit, associated with the main accretion region. Analysis of the X-ray lightcurve indicates that the white dwarf spin period is now similar to the orbital period. When thisinference is combined with the ground-based photometry, we find that the most probable explanationfor the observed behavior is a system that has become fully synchronized. The X-ray spectrum andluminosity of V1500 Cyg are consistent with a high state polar. The optical and near-IR light curvescan partially be explained as a heavily irradiated secondary star, but exhibit strong deviations fromthat model. We find that the most likely explanation for the observed excesses in these light curvesis cyclotron emission from distributed accretion regions, or from a multipole geometry.

Key words: X-ray: stars — infrared: stars — stars: cataclysmic variables — stars:individual (V1500 Cygni)

1Partially based on observations obtained with the ApachePoint Observatory 3.5-meter telescope, which is owned andoperated by the Astrophysical Research Consortium.

1 INTRODUCTION

V1500 Cyg, Nova Cygni 1975, vies with CP Pup for thetitle of the most luminous classical nova ever observed, hav-ing MVmax ' −10.7. V1500 Cyg is remarkable for a numberof other attributes rarely seen in classical novae (CNe). Forexample, during the weeks before its eruption, it showed adramatic ∼ 7 mag brightening (Collazzi et al. 2009). Fol-lowing eruption, the dominant photometric period declinedfrom 0.140 d in 1975 (Semeniuk et al. 1977), to the currentstable period of 0.13962 d in 1978 (Patterson 1979). Overthis interval the amplitude of the variations in its visual lightcurve increased to ∆m ∼ 0.65 mag. Even more interesting,however, were the observations by Stockman, Schmidt &Lamb (1988) that V1500 Cyg exhibited circularly polarizedlight, and that the polarized light curve had a period thatwas 2% shorter than the photometric period. V1500 Cygbecame the first “asynchronous polar,” where a highly mag-netic white dwarf spins at rate that is slightly different fromits orbital period. Unlike the Intermediate Polars, where themagnetic white dwarf primary also spins at a different rate

? E-mail: [email protected]† E-mail:[email protected]

than the orbital period, asynchronous polars do not appearto have accretion discs.

Continued polarimetric study of V1500 Cyg showedthat the white dwarf spin period was slowly increasing, andthe system would attain synchrony within a few hundredyears. Initial explorations of the processes that might gener-ate the torques necessary to create this rapid spin-down rateshowed that they were at least an order of magnitude toosmall to match observations (Schmidt & Stockman 1991,and references therein). In contrast, Campbell & Schwope(1999) found that the torque resulting from induced elec-tric fields and current flow in the secondary star from anasynchronous magnetic white dwarf could explain the spin-down time scale for V1500 Cyg. In their modeling, if thewhite dwarf had a mass of 0.5 M�, the spin-down time scalewas of order 50 yr. If the white dwarf mass was higher thanthis, the spin-down time scale would be longer. Given theviolence of the CNe eruption of V1500 Cyg, and the pres-ence of a modest enhancement of neon in its ejecta, it isassumed that the mass of the white dwarf in the system islarge (Politano et al. 1995, Lance, McCall, & Uomoto 1988).

Without adequate knowledge of many important sys-tem parameters, it is difficult to assess the viability of anymagnetohydrodynamic torque model that attempts to ex-plain the spin-down rate of V1500 Cyg. To enable more ro-bust models requires measurements of the current spin rate,the current mass accretion rate, an estimate of the orienta-tion of the magnetic field with respect to the binary orbit,and limits on the magnetic field strength of the primary.Since the accretion regions of magnetic CVs are sources of

MNRAS Advance Access published April 22, 2016 at N

ew M

exico State University L

ibrary on April 25, 2016

http://mnras.oxfordjournals.org/

Dow

nloaded from


2 Harrison & Campbell

high energy photons, X-ray observations of such objects canbe used to directly measure the white dwarf spin periodwith little ambiguity. Such data can also provide limits onthe accretion rate, and often allow some constraints on theaccretion geometry. When these data are combined with co-temporal optical and infrared light curves, there is the possi-bility of additional insight into the properties of the system.

Harrison et al. (2013a) showed that the partial XMMlight curve of V1500 Cyg was indicative of a brighteningthat was consistent with the behavior expected of a polaraccretion region. We have obtained new XMM observationsthat cover four orbital cycles of V1500 Cyg and these reveala strong, periodic brightening once per orbit. Period estima-tion techniques applied to the X-ray data suggest that thecurrent white dwarf spin period is very close to the orbitalperiod. When we combine this result with new optical andnear-IR photometry, we conclude that the white dwarf inV1500 Cyg appears to have achieved synchronization, andis again a true polar. In the next section we describe thedata reduction process, in section 3 we model the X-ray andground-based data; we discuss the results and our conclu-sions in section 4.

2 DATA

We have obtained both X-ray and U -band data for V1500Cyg using XMM. We have also used the New Mexico StateUniversity (NMSU) 1 m telescope to obtain BV RI lightcurves, and NICFPS on the Astrophysical Research Consor-tium (ARC) 3.5 m telescope at Apache Point Observatoryto obtain near-infrared photometry. We elaborate on thesedata sets here.

2.1 XMM Data

V1500 Cyg was observed with XMM −Newton beginningat 2014 11 11 13:35:11 UT, and was monitored continuouslyfor 53.3 ksec. XMM (Lumb et al. 2012) consists of threeco-aligned grazing-incidence mirror assemblies that feed theEuropean Photon Imaging Camera (“EPIC”), and two re-flection grating spectrometers (the “RGS”). XMM is alsoequipped with a co-aligned optical monitor (“OM”) for ob-taining UV/optical photometry. The EPIC camera consistsof three detector arrays, one in each of the focal planes ofthe telescopes: the back-illuminated CCD array called the“PN’, and the two front-illuminated CCD arrays called the“MOS”. For our observations, the PN camera was set-upin “small window” mode where just the central CCD chipis read. The MOS detectors were used in full frame mode.Due to the visual/UV faintness of V1500 Cyg, the “thin”blocking window could be used so as to increase the S/N ofthis source. We used the U -band filter on the OM to obtaina light curve of V1500 Cyg. Given its expected brightness,along with the significant reddening to this object, 800 s ex-posure times were employed to insure sufficient S/N ratiosfor the U -band observations.

We used the standard reduction process for these dataemploying the “Scientific Analysis System” (SAS1) devel-

1 http://xmm.esac.esa.int/external/xmm data analysis/

oped specifically for XMM . To enable proper period search-ing, the data were first barycentric corrected using the SAStask barycen. Due to its orbit, XMM data can be compro-mised by a strong soft proton background (see Carter &Read 2007). To remove these events from our data set, theprocess outlined in the “XMM-Newton ABC Guide”2 wasfollowed. Briefly, this process consists of using regions awayfrom X-ray sources to identify time intervals when the dataare compromised by proton flaring events, and then extract-ing data that exclude those time intervals. To produce the X-ray spectra presented below, only events from the screeneddata set were used. For producing the X-ray light curves,however, the data were simply background-subtracted us-ing an annulus around the point source extraction aperture.Fortunately, the soft proton background was relatively lowthroughout these observations.

2.2 Ground-based Data

We used the NMSU 1.0 m telescope to obtain BV RI lightcurves for V1500 Cyg on 2014 August 6, and 31, and V -band light curves on 2007 October 20 and 21. The NMSU1.0 m telescope at Apache Point is equipped with an E2V2048x2048 CCD with standard Johnson-Cousins UBVRI fil-ters. These data were processed in the normal way, and cal-ibrated light curves for the program object were generatedusing known UBV RI photometry for several nearby fieldstars listed in Kaluzny & Semeniuk (1987). We obtainedJHK photometry of V1500 Cyg using the Near-InfraredCamera and Fabry-Perot Spectrometer (“NIC-FPS”) on theARC 3.5 m telescope on 2014 August 31 and October 30.Calibrated differential photometry was obtained relative toseveral nearby stars that are in the 2MASS catalog. An ob-servation log is presented in Table 1.

3 RESULTS

3.1 The X-ray and U-band Light Curves of V1500Cyg

The X-ray and U -band light curves from XMM are shownin Fig. 1. The top panel has the MOS1 (black) and MOS2(red) data, while the middle panel has the PN light curve(the light curve model in green will be described below).The presence of reflection gratings in the optical paths ofthe MOS detectors, to feed the RGS, are partially respon-sible for the lower observed count rates from these deviceswhen compared to that for the PN. The U -band data arepresented in the bottom panel (the light curve model plot-ted in this panel will be described in §3.3). The long dura-tion of the observations allowed the partial coverage of fiveX-ray maxima. The X-ray maxima are not symmetric, nordoes their morphology repeat from one cycle to the next.The fourth X-ray maximum was much briefer, and had afaster rise to, and fall from, maximum. The “missing” sec-tion from the early part of this brightening corresponds tothe faintest part of the U -band light curve. Perhaps therewas a pause in the mass accretion at this particular instant.Another such minimum at 49,000 s in the U -band light curve

2 https://heasarc.gsfc.nasa.gov/docs/xmm/abc/abc.html

at New

Mexico State U

niversity Library on A

pril 25, 2016http://m

nras.oxfordjournals.org/D

ownloaded from



Table 1. Observation Log

Resource Start Time Stop Time Wavelengths

XMM 2014-11-11 13:35:11 2014-11-12 04:48:31 X-ray/U -bandNMSU 1m 2007-10-20 04:17:09 2007-10-20 06:58:39 V

” 2007-10-21 04:12:13 2007-10-21 08:10:18 V

” 2014-08-06 06:00:59 2014-08-06 10:10:14 BV RI” 2014-08-31 02:42:22 2014-08-31 09:34:41 BV RI

ARC 3.5 m 2014-08-31 03:42:13 2014-08-31 06:53:11 JHK

” 2014-10-30 00:47:17 2014-10-30 04:31:15 JHK

also corresponds to a dip in the X-ray flux. Note that themaxima in the U -band light curves occur during the X-rayminima, and the large excursions seen in these data cannotbe directly attributed to thermal hotspot emission from themain accretion region.

We extracted a light curve from the PN event list binnedinto 1 s intervals and searched it for the best fitting pe-riod using the XRONOS3 task efsearch. The ephemeris ofSchmidt, Liebert, & Stockman (1995) would predict the spinperiod at the time of these observations to be Pspin = 11880± 4 s. Obviously, with temporally limited observations on asingle epoch, it is difficult to determine a precise spin period.Using a range of input values for the binning inside efsearch,leads to ambiguous results. The period with the strongestratio of the measured to expected variance (Var/expVar =58) is Pspin = 12031 s, and occurs when we use 25 phasebins per period. However, there are a wide range of statis-tically significant periods ranging from 12017 s, to 12048 s.The limited time series, the low S/N of the X-ray data, andasymmetries of the maxima, do not allow for a more accu-rate period measurement with efsearch. As an alternativeto efsearch, we used the period searching, light curve fittingsoftware “Period04” (Lenz & Breger 2005). When input intoPeriod04, the PN light curves binned by 10 s and 60 s, yieldperiods of 12062 s and 12058 s, respectively. These are es-sentially identical to the orbital period of V1500 Cyg: Porb

= 12062.5 s. Even without a robust value for the spin pe-riod, it is clear that the white dwarf is now much closer tosynchronization than would have been predicted.

Given these results, and when analysed in conjunctionwith the ground-based data presented below, we believe thespin period of the white dwarf is now synchronous with theorbital period. If we fold the (1 s binned) light curve onthe orbital period, we can attempt to extract other usefulinformation. The folded PN light curve is shown in Fig. 2,where we have set phase φspin = 0 to be equal to the cen-troid of the primary maximum. Once we have folded thesedata, we find that there is little evidence for a secondarymaximum near phase 0.5. In their polarization light curve,Schmidt & Stockman (1991) found direct evidence for twopole accretion. A weaker X-ray maximum could arise from alower accretion rate, or due to a poor viewing geometry forthis secondary pole. From the X-ray light curve data alone,the reality of this pole cannot be established. In this way,V1500 Cyg resembles the polar BL Hyi, an object that hasa very weak/rarely detected secondary X-ray pole, but one

3 https://heasarc.gsfc.nasa.gov/docs/xanadu/xronos/

0 10000 20000 30000 40000 50000-10

0

10

20

30

0 10000 20000 30000 40000 50000-10

0

10

20

30

40

0 10000 20000 30000 40000 5000019

18.5

18

17.5

Figure 1. The light curves of V1500 Cyg derived from the XMM

observations. The top panel has the MOS1 (black) and MOS2

(red) data, while the middle panel shows the PN light curve.The green curve over-plotted on the PN data set is a light curve

model that was generated to explain the folded light curve (see

Fig. 2) as due to two bright spots on the white dwarf primary.The changing morphology of the X-ray maxima make absolute

spin period determination difficult. The bottom panel are the U -band data from the OM. The solid cyan line in this panel is the

light curve model described in the text in section §3.3.

that is prominent in polarimetric light curves (see Schwopeet al. 1998).

We can use the Wilson-Divinney light curve modelingcode WD20104 to explore possible geometries that can re-produce the profile of the folded light curve. Given the largenumber of possibilities for the input parameters in any lightcurve modeling project, we need to supply some constraintsto derive anything useful. As shown in Schmidt & Stock-man (1991), there are two accreting poles that appear to belocated on opposite sides of the white dwarf. In addition,the switch between positive and negative branches is welldefined, and suggest that these regions are not visible for

4 ftp://ftp.astro.ufl.edu/pub/wilson/lcdc2010/

at New

Mexico State U




ownloaded from



0 0.5 1 1.5 2

0

1

2

3

4

Figure 2. The PN light curve folded on a period of 12062.5 s. The

red line is a light curve model that reproduces the morphology ofthe X-ray data set. In this model the binary orbital inclination

is 60◦, and the two spots are opposite each other on the white

dwarf: separated by 180◦ in longitude, and having co-latitudes ofβ1 = 120◦, and β2 = 60◦. Phase 0 here is defined as the center

of the X-ray maximum.

much more than one half of a rotation period. The spec-tropolarimetry presented in Schmidt et al. (1995) suggeststhat the duration of the negative polarization may be brieferthan that for the positive polarization. The polarization dataalso suggest that neither pole is face-on at transit. Schmidtet al. derived an orbital inclination of i = 60◦ that we acceptfor the remainder of this paper.

We start by separating the two spots by 180◦ in longi-tude. A spot with a co-latitude of β = 90◦ will be visiblefor one half of a rotation period (or orbit here). Larger co-latitudes mean briefer transits, smaller co-latitudes producelonger transits. While there are a large number of light curvemodels that play the co-latitude of the spot off against thespot radius, there are basically two families of solutions: thespots are small in size and in the same upper hemisphere(above the line-of-sight), or the spots are in both upper andlower hemispheres, and are much larger. If both spots wereto have β < 60◦, they would remain visible at all times, pro-ducing broad maxima. The optical/IR light curves discussedbelow appear have a sharp minimum near φ = 0, suggestingthat at least one of the poles disappears from view. Thisis also suggested by the polarization results, and thus weprefer a model where we locate one of the poles below ourline-of-sight. In Fig. 2, we plot a model where we have twovery large hotspots (radius = 45◦) separated by 180◦ in lon-gitude, and that have co-latitudes of β1 = 120◦, and β2 =60◦. This large spot radius is required to generate the broadprimary maximum. While the Wilson-Divinney code modelsthese spots as circles, a long arc of X-ray emission orientedin longitude could also reproduce a broad maximum. Note

that we have set β2 = 60◦, and the secondary spot would beface-on at transit. Unfortunately, the location of this spot isimpossible to constrain given the low count rate.

3.2 The X-ray Spectra of V1500 Cyg

To extract X-ray spectra for V1500 Cyg, we rigorouslyscreened the X-ray data for proton events by constructing“good time” intervals free of such contamination. We thenfurther delineated regions containing the X-ray maxima andminima from this screened event list to extract spectra forboth phase intervals. The spectra for the X-ray maxima areshown in Fig. 3, and that at light curve minimum in Fig.4. Due to the very low count rate during minimum, onlythe PN data were extracted, and they had to be heavilybinned to achieve a spectrum that could be modeled. Weconstructed model spectra using XSPEC5 with three emis-sion components consisting of a soft blackbody, a thermalbremsstrahlung source, and an emission line, all absorbed bya hydrogen column. As discussed in Harrison et al. (2013a),the visual extinction to V1500 Cyg is AV = 1.5 mag, lead-ing to NH = 2.5 × 1021 cm−2. The hydrogen column waskept “frozen” during all XSPEC model fitting. It has beenshown previously that there is often evidence for additionalabsorption in the X-ray spectra of polars, in step with the or-bital period (c.f., Wolff, Wood, & Imamura 1999), and oftenattributed to the accretion stream and the changing view-ing geometry. We do not believe our data are of sufficientquality to allow this parameter to be unconstrained duringmodeling, nor to search for possible orbital modulations.

We chose a simple three component model as it a typicalspectrum for high state polars (see Beuermann, El Kholy,& Reinsch 2008). The best fitting model has a thermalbremsstrahlung source with kTtb = 100 keV, and a softblackbody source with kTbb = 60 eV, resulting in χ2

red =1.03. With kTtb fixed, both cooler and hotter values for theblackbody component result in poorer fits. For example, set-ting kTbb = 40 eV, gives χ2

red = 1.25, while kTbb = 80 eVresults in χ2

red = 1.13. The very high temperature of thethermal bremsstrahlung component is not well constraineddue to the lack of sufficient hard photons in the X-ray data.For example, setting kTtb = 30 keV (with kTbb = 60 eV),leads to χ2

red = 1.09, while kTtb = 40 keV gives χ2red =

1.06. These lower values for the temperature of the thermalbremsstrahlung component in V1500 Cyg are more in linewith those observed for the typical polar in a high state (seeWarner 1995, Landi et al. 2009, Christian 2000). The finalmodel spectrum, listed in Table 2 and plotted in Fig. 3, haskTtb = 40 keV and kTbb = 60 eV. We conclude that dur-ing the bright phases, the X-ray spectrum of V1500 Cyg issimilar to, but possibly hotter than, the typical high statepolar.

The X-ray flux during the minimum light phasesis very low, making its spectrum harder to constrain.With kTbb fixed at 60 eV, we find the best fitting ther-mal bremsstrahlung component has kTtb = 2.1 keV (χ2

red

= 0.96). If we let both temperatures be free, the bestfit XSPEC solution has kTbb = 132 eV, and kTtb =11.7 keV (χ2

red = 0.88). In either scenario, the thermal

5 http://heasarc.gsfc.nasa.gov/xanadu/xspec

at New

Mexico State U




ownloaded from



Table 2. X-Ray Spectral Parameters

Component Primary Maximum Secondary Maximum

NH (fixed) 2.45 × 1021 cm−2 2.45 × 1021 cm−2

TBB 60 ± 1.3 eV 50 ± 6.7 eV

Normalization (BB) 1.4 ± 0.1 × 10−5 4.5 ± 0.3 × 10−6

TTB 40 ± 9 keV 1.85 ± 0.7 keVNormalization (TB) 3.2 ± 0.2 × 10−5 7.9 ± 0.2 × 10−6

Gaussian Line Energy 6.6 ± 0.2 keV 6.6 ± 0.4 keV

Gaussian Line Width 0.4 ± 0.3 keV 0.4 ± 0.4 keVNormalization (Line) 2.4 ± 0.6 × 10−6 1.0 ± 0.6 × 10−6

HJD(TDB) of Maximum 2456073.3731 ± 0.0010 · · ·

Figure 3. The maximum light X-ray spectra of V1500 Cyg.

Green data points are from the PN, the black from the MOS1,and the red from MOS2. The solid line through each data set is

the best fitting XSPEC model for that detector, consisting of a

thermal bremsstrahlung source, a blackbody, and a Gaussian linewith the parameters listed in Table 2.

bremsstrahlung source is much cooler during minimum lightthan during the X-ray maxima. Lower temperatures for theblackbody source directly lead to lower temperatures for thethermal bremsstrahlung source. A blackbody temperature ofkTbb = 50 eV requires kTtb = 1.9 eV to achieve χ2

red = 1.08.This is the model listed in Table 2, and plotted in Fig. 4.Even though not tightly constrained, this result does indi-cate the presence of a second accretion region.

Over the energy interval 0.5 to 10 keV, the peak brightphase X-ray flux is 6.3 ± 0.5 × 10−13 erg cm−2 s−1. Fora distance of 1 kpc, the X-ray luminosity is 7.5 ± 0.6 ×1031 erg s−1. This should be compared with the value of1.6 × 1032 erg s−1 for the hard component of AM Her ina high state (Ishida et al. 1997), or 1.3 × 1031 erg s−1 forBL Hyi (Beuermann & Schwope 1989). Using equation 2.90in Warner (1995), and a shock temperature of 40 keV, thelower limit on the mass of the white dwarf in V1500 Cygis 1.13 M�. This estimate for the mass of the white dwarfis in line with expectations, given the rapid evolution of itsCNe eruption. Assuming Lacc = GMwdM/Rwd, we find M∼ 6.6 × 1014 gm s−1. Thus, the spectral parameters, masstransfer rate (c.f., Ramsay & Cropper 2003), and the X-rayluminosity of V1500 Cyg are consistent with a normal polarin a high state.

Figure 4. The minimum light PN spectrum of V1500 Cyg. TheXSPEC model fitted to these data has a blackbody temperature

of kTbb = 50 eV and kTtb = 1.9 eV, and a Gaussian line (see

Table 2).

3.3 The Optical/Near-IR Light Curves

We obtained four epochs of ground-based light curve data:V -band photometry in 2007 October, BV RI photometry on2014 August 06, BV RI and JHK data on 2014 August 31,and JHK only on 30 October. The BV RI light curves areshown in Figs. 5 and 6, while the JHK data are shown inFigs. 7 and 8. The JHK light curves shown here are es-sentially identical to those we obtained in 2006 (Harrisonet al. 2013a). The 2007 V -band data are plotted in Fig. 9.We have extracted the V -band data from Semeniuk, Olech,& Nalezyty (1995), and present it as Fig. 10. The morphol-ogy of the V -band light curve appears to have changed littlesince 1995, though the source has faded by nearly a magni-tude. In all figures, the photometry has been phased usingthe minimum light ephemeris in Semeniuk et al. (1995). Thisprecise ephemeris was established using visual data spanningnearly a decade, and at the epoch of our observations theerror bars on the phasing of the minima are < 0.5%. In thebottom panel of Fig. 5 we combine all of the light curve datato enhance the phase resolution by applying vertical offsetsto the individual light curves to overlay them. For the 31August data, we repeat the V -band panel, but with the in-clusion of a light curve model that has two large hot spotsat the longitudes and latitudes of the two accretion regionsfound above (see Fig. 2). We will discuss this panel in §3.4.

at New

Mexico State U




ownloaded from



0 0.5 1 1.5 2

19.5

19

18.5

18

Figure 5. The BV RI light curves of V1500 Cyg for 2014 August

06. In the bottom panel we plot all four data sets with arbitraryvertical offsets to improve the phase resolution. The solid line in

each panel is the light curve model described in the text. The

solid vertical lines in each panel are the locations of the transit ofthe primary X-ray maximum, while the dashed lines is the transit

phase of the secondary pole. These data have been repeated for

clarity.

As elaborated by Schmidt et al. (1995, and referencestherein), the large optical modulations are believed to be dueto a hot white dwarf heating a cool secondary star. They de-rived Twd = 90,000 K, T2 = 3000 K, and i = 60◦. As shownin Harrison et al. (2013a), such a model generates a flat-bottomed minimum. None of our light curves truly has sucha flat bottom, though the I-band light curve on 31 Augusthas the flattest minimum. We cannot produce a sharp min-imum in light curve models, regardless of the inclination,assuming an irradiated secondary star. The complexity ofthe V -band light curves suggest additional emission compo-nents in the binary system. The two epochs of JHK datawere similar, though the 31 August data set was plaguedby poor seeing near φ = 0. Morphologically, they are essen-tially identical, though the large excess from φ =0 to φ =0.5 seen in the 30 October data appears to be somewhatsmaller on 31 August. We focused our modeling effort ontrying to simultaneously match both the I-band data for 31August and the 30 October JHK light curves with only anirradiated secondary star model.

We ran WD2010 over a wide range of tempera-ture and inclination (including the appropriate wavelength-dependent light curve modeling parameters required by theWilson-Divinney code, see Harrison et al. 2013c). We settledon a model with Twd = 52,000 K, T2 = 3100 K, and i = 60◦.As can be seen in the plots, this model reasonably fits theI-band data for 31 August, and appears to perfectly explainone half of the JHK light curves (over 0.5 ≤ φ ≤ 1.0). While

Figure 6. The BV RI light curves for 31 August, as in Fig. 5,

except that the V -band data has been repeated in the middlepanel and overlaid with a two hotspot model to show the differing

transit durations of the magnetic poles, see §3.4.

0 0.5 1 1.5 2

18.5

18

17.5

17

16.5

0 0.5 1 1.5 2

17.5

17

16.5

16

Figure 7. The JHK light curves for 2014 August 31. The solidline in each panel is the irradiated secondary star light curvemodel described in the text. As in the previous figures, the phasesof accretion region transits are identified.

we could adjust the model parameters, such as the inclina-tion, so that it could fully explain the I-band maximum, wefelt that the deviation from the irradiated model on eitherside of minimum in this bandpass indicated that there wasexcess emission present at all phases. There is a large ex-cess in the near-infrared light curves from φ = 0 to φ = 0.5.There is a significant excess above the irradiated secondarystar model at nearly all phases in the UBV R data. We also

at New

Mexico State U




ownloaded from



0 0.5 1 1.5 2

18.5

18

17.5

17

16.5

Figure 8. The JHK light curves of V1500 Cyg for 30 October

as in Fig. 7.

0 0.5 1 1.5 219.5

19

18.5

18

17.5

0 0.5 1 1.5 219.5

19

18.5

18

17.5

Figure 9. The V -band light curve of 2007 October 21 (top), andfor October 20 (bottom). The solid green line in both panels is the

irradiated secondary star model described in the text, offset by

∆ V = −0.2 mag. The light curve for 20 October is incomplete,with a gap between phase 0.2 and 0.4.

plot this model, offset by −0.2 mag, on to the V -band lightcurves of 2007 October shown in Fig. 9. The higher tempo-ral cadence of the 2007 October 21 V -band data reveals anexcess above the light curve model that is almost identicalin amplitude and morphology to the JHK excesses seen in2014 October. The data on 20 October finds larger pre- andpost-maxima excesses, though the light curve model passesthrough the photometry near phases 0.0 and 0.5.

Since the light curve models are normalized to fit thedata, we need to examine the spectral energy distribution(SED) to insure that our model is realistic. In Fig. 11 we plotthe faintest U -band data point from the XMM light curve,the 31 AugustBV RI data, and the JHK fluxes at phase 0.0.The solid circles in this plot are the minimum light data, and

0 0.5 1 1.5 2

18.5

18

17.5

Figure 10. The V -band light curves of V1500 Cyg in 1995, ex-

tracted from Semeniuk et al. (1995), and phased to their minimum

light ephemeris. Data for 27 July are plotted in blue, 28 July ingreen, and 29 July in red. The data for 27 July have been offset

by ∆V = −0.18 mag, and those for 28 July have been offset by

∆V = −0.05 mag, to achieve maxima that agree with that for 29July. The irradiated secondary star model discussed in the text is

plotted as the solid line, after being offset by ∆V = −0.87 mag.

have been fitted with the sum of two blackbodies: Tbb1 =52,000 K + Tbb2 = 3,100 K. The resulting SED is consistentwith the observed data at light curve minimum with onlythe stellar components as sources of luminosity. Note thatfor an M4V (T2 = 3100 K) to explain the spectrum requiresa distance of 1 kpc to V1500 Cyg. The Wilson-Divinneycode also predicts the luminosity of the white dwarf in eachof the modeled bandpasses. For the hot blackbody to re-produce the observed V -band flux also requires d ∼ 1 kpc.Lance et al. (1988) have summarized the distance estimatesfor V1500 Cyg and conclude that 1.2 ± 0.2 kpc is a rea-sonable mean. Note that more recently, Slavin, O’Brien &Dunlop (1995) found a distance of 1.5 kpc for V1500 Cygfrom a nebular expansion parallax measurement. As shownin Harrison et al. (2013b), when compared to actual paral-laxes, the secondary distance estimation techniques for CNeare not reliable. In the four cases examined, only the nebularexpansion distance for GK Per agrees with its astrometricdistance. For DQ Her, the nebular expansion parallax is 40%larger than its true distance. Thus, if we assume d = 1 kpc,our light curve model reproduces the observed spectrum atphase 0. If, of course, the distance to V1500 Cyg is 1.5 kpc,the light curve model would need to be scaled downward by0.9 mag, resulting in very large excesses in every bandpassand at all phases of the orbit.

It is important to mention that Barman, Hauschildt,& Allard (2004) used PHOENIX atmosphere models to ex-plore the process of heavily irradiating a CV secondary starjust like that found in V1500 Cyg. Instead of the typicalbolometric albedo of 0.5 used for convective stars in lightcurve modeling (as used here), they found a much highervalue of ≈ 0.8 for scenarios like V1500 Cyg. Essentially, theirradiation of the convective star makes it appear to be morelike a non-convective star, and thus the bolometric albedo

at New

Mexico State U




ownloaded from



Figure 11. The spectrum of V1500 Cyg at φ = 0 (solid circlesand lines), and φ = 0.5 (open circles, dashed lines). The data

for minimum light have been fitted with a two blackbody model

comprised of a hot white dwarf (Teff = 52000 K, blue), and a coolsecondary star (Teff = 3100 K, red). The sum of both components

is plotted as a black solid line. The data for φ = 0.5 have also

been fit with a two blackbody, but now the cooler componenthas Teff = 5300 K (red dashed line). The sum of the hot white

dwarf and the hotter hemisphere of the irradiated secondary star

is plotted as a black dashed line. The BV RI excesses at this phaseare apparent.

is closer to the value (1) found for such stars. If we set thebolometric albedo to 0.8, then we have to lower the valuefor the inclination to i = 35◦ to reproduce the IJHK lightcurves. Horne & Schneider (1989) argue that the strong or-bital modulation in both the continuum and the emissionlines, and the lack of an eclipse, constrain the inclination tobe 40◦ ≤ i ≤ 70◦. If we use our lower limit on the white dwarfmass derived from the shock temperature, the estimated or-bital inclination from their Fig. 14 is i > 50◦. Similar limitswere derived by Schmidt et al. (1995).

Returning to the SED of V1500 Cyg at φ = 0.5, theUBV RI data have excesses above the irradiated light curvemodel, but J through K do not. We can estimate the ap-parent temperature of the heated face of the secondary starusing the procedure outlined in Brett & Smith (1993), andfind T2 = 5300 K. When we add a blackbody with thistemperature to the same hot blackbody used at phase 0, seeFig. 11, we can reproduce the JHK observations. The modelSED at I falls 0.2 mag below the observed flux, which is theobserved difference between the light curve model and thephotometry for 31 August. Thus, the irradiated light curvemodel remains dominant at both inferior and superior con-junctions.

3.4 Possible Origins for the Light Curve Excesses

If we accept the presence of an irradiated secondary star,and assume the white dwarf spin is now synchronized withthe orbital period, we can attempt to explore the origin of

Figure 12. The difference plot, where we have subtracted theirradiated secondary star model from the V -band data for 2007

October 21, and from the J- and K-band data for 2014 October

30. The locations in phase of the transits of the two accretionregions are identified.

the strong excesses seen in the light curves. Using the timeof X-ray maximum listed in Table 2 (this is the time of thecentroid of the third maximum), we can predict its phase po-sition in the other light curves. In Figs. 5 through 8 we delin-eate the locations of the primary (solid lines) and secondary(dashed lines) X-ray maxima for the BV RI and JHK lightcurves. The primary X-ray maximum, presumably when wehave the most direct view to the accretion region, occursat binary orbital phase φ = 0.822, leading the secondarystar around the orbit. This is the configuration observedfor most polars (e.g., Cropper 1988). Examination of Fig. 8shows that the large near-IR excesses (above the light curvemodel) are not centred on the transit of either accretion re-gion. There appears to be a small peak, or inflection point,in all three infrared bandpasses at the phase when we haveour best view of the secondary pole. However, the bulk ofthe excess occurs before this pole reaches the centre of thedisc, and quickly weakens as it rotates away. In Fig. 12, wesubtract6 the light curve model from the V -band data for2007 October 21, and from the J and K-band light curveson 2014 October 30. The peaks in the excesses in the V -band are close to, but not quite centred on, the transit ofthe two accretion regions. The excesses in the near-IR areskewed relative to the secondary pole, and reach their great-est deviation well before the accretion spots are centred onthe facing hemisphere of the white dwarf. It is importantto note that the same excess in the V -band data for 2007October 20 began earlier, and was substantially larger, thanthat for 21 October, and more closely resembled those seenin the near-IR.

6 We have used IRAF to perform this subtraction, and thus manyof smaller fluctuations seen in the actual light curve data areremoved due to the interpolation process used inside IRAF.

at New

Mexico State U




ownloaded from



Figure 13. A difference plot for the UBV RI data for 2014 Au-

gust 31 as in Fig. 12.

In the middle panel in Fig. 6, we plot the light curve oftwo round spots of identical temperature, at the locationsderived from the X-ray observations (note that this modelis only for the visualization of the accretion pole positions).The secondary pole is in the hemisphere above the view-ing plane, and thus is visible for longer than the primarypole; it remains visible for the entire orbit, though grazingthe horizon at times. The rise in the excess shortly afterthe minimum near φ = 0 suggests that the excesses seenin BV RIJHK over the interval 0.0 ≤ φ ≤ 0.5, originatefrom the secondary pole. If we trust the normalization ofthe irradiated secondary star model, the primary pole doesnot produce a significant excess in the near-IR. This is truefor all three epochs of JHK light curve data. This is notthe case in the visual: there is a strong excess in BV R inthe phase interval 0.5 ≤ φ ≤ 1.0. This excess is highly vari-able compared to that associated with the secondary pole.In Fig. 13, we have subtracted the light curve model fromthe UBV RI photometry. It is clear that we now have an ex-cess in all the visual wavebands that can be associated withthe primary pole. This excess is weakest in I, and strongestin U and V . The U -band data reveals a minimum near thetime when the secondary pole transits our line-of-sight.

As discussed in Schmidt et al. (1995), the most likelysource for the optical/IR excesses is cyclotron emissionand/or emission from the accretion columns. The fact thatthe U -band reaches a maximum near φ = 0, when bothaccretion regions are near the limb, shows that two hotblackbody sources on the white dwarf photosphere are notresponsible for the excesses seen in this bandpass. If such

sources are not apparent in the U -band, they will be evenless significant at longer wavelengths. As shown in Harrison& Campbell (2015, their Fig. 6), the majority of the fluxfrom optically thin cyclotron emission from the fundamen-tal harmonic is beamed along the axis of the magnetic field,while the higher harmonics emit the majority of their fluxperpendicular to field. As the temperature and/or opticaldepth increase, the emission from the fundamental becomesmuch more aligned with the field axis, with much less emis-sion in the perpendicular direction (and, to a lesser extent,this trend is also seen for the higher harmonics). Since theoptical/IR excesses are near maximum when the viewingangles to the two poles are relatively low, optically thick cy-clotron emission is the leading candidate for the observedexcesses.

We do not believe that emission from the accretionstream/funnel is as viable an explanation for the light curveexcesses. The main reason is that at no time have we ob-served an excess in the near-IR in the phase interval 0.5≤ φ ≤ 1.0, while the visual data shows a strong, and vari-able excess during this time. A free-free emission source thatis strong in the visual, would be strong in the near-infrared(see Harrison 2014). The opposite is true for the phase inter-val 0.0 ≤ φ ≤ 0.5: the strong excesses seen in the visual areseen in the near-IR. If we consider the same mechanism tobe responsible for both excesses, as suggested by the similarmorphologies of the differenced light curves, this then rulesout dilution by the stellar components as the cause for di-minished free-free emission from the primary pole region atlonger wavelengths. Another inference for ruling out the ac-cretion stream/funnel is provided by the constraint derivedby Schmidt et al. that the source of additional flux over thewavelength interval 1600 A ≤ λ ≤ 7000 A, had a luminosityof ∼ 0.25 L�, or about 5% of the luminosity of the whitedwarf in 1992. Schmidt et al. found that they could probablyexplain the origin of the excess they observed with clumpy,large M , optically thick accretion columns that interceptand re-radiate 5% of the system’s luminosity. However, thecurrent peak B-band excesses have relative luminosities thatare 50% that of the system in this bandpass.

3.4.1 A Cyclotron Interpretation for the Excesses

In the possible locations for the accretion regions discussedabove, the two scenarios were that either both spots remainvisible at all times, or the primary pole is located at a co-latitude that is below our sight line, while the secondarypole is visible for nearly the entire orbit. We now attemptto reconcile the light curves with a cyclotron interpretation,starting with the excesses observed in the phase interval 0.0≤ φ ≤ 0.5, which we associate with the secondary X-raypole. A cyclotron model for these excesses needs to explainthe following light curve features: 1) the near-IR excess turnson near φ = 0.0, 2) the excess is not at its strongest whenthe secondary accreting pole is near the limb, 3) the ex-cess is largest as the secondary pole rotates toward transit,but drops as the spot transits, 4) the excess rapidly de-clines after spot transit, disappearing before the X-ray spotis self-eclipsed, and 5) the excesses have slightly differentmorphologies in different bandpasses.

A small accretion region located at β < 60◦ would bevisible for most of the orbit, and thus the light curve of the

at New

Mexico State U




ownloaded from



Figure 14. The multi-wavelength light curves generated by a

constant-Λ cyclotron model with B = 80 MG, T = 10 keV, Λ =5, β = 60◦, and i = 60◦ over two rotational cycles.

cyclotron emission would be broad, and slowly changing. Asthe pole rotated beyond the limb, possibly suffering a briefself-eclipse near anti-transit, the orientation of the field lineswould result in the majority of the cyclotron emission beingradiated away from the observer. We have modeled this case,using the constant-Λ cyclotron code discussed in Harrison &Campbell (2015). The light curve for a cyclotron model withB = 80 MG, T = 10 keV, Λ = 5, β = 60◦, and i = 60◦ ispresented in Fig. 14. Note that the phasing in this diagramhas φ = 0 at spot transit. The cyclotron emission turns-on assoon as the pole is on the limb, and disappears as it rotatesbeyond the limb. There is a dip in all of the bandpasses atthe time of transit. Obviously, this cyclotron model has littlein common with the observed light curves, except that theexcesses behave in a similar fashion in all of the bandpasses.Every cyclotron model, regardless of the co-latitude of thespot, or value for the orbital inclination, will have symmetricemission around the time of pole transit. Thus, such modelscannot explain the skewed excesses seen in the light curvedata.

As discussed in Beuermann, Stella, & Patterson (1987),if a field line is offset from the magnetic axis by an angleθ, the field line is tilted from the radial direction by ∼ θ/2.Here we investigate what the cyclotron model would looklike if we tilt the field lines (in longitude only). A modelwith a tilt of 20◦, corresponding to a phase offset of 40◦

(∆φ = 0.11), is plotted in Fig. 15. The main result is thatthe duration of the cyclotron excess is shortened by about0.1 in phase. The minima near transit seen for the un-bent

Figure 15. A cyclotron model as in Fig. 13, but with field lines

tilted by 20◦.

field line model are no longer significant, as the viewing an-gle at this time is non-zero (20◦). This leads us to proposethat both accretion regions have a core-tail type distribution(c.f., Cropper 1986), with the tail aligned in the direction ofrotation. Regions in the tail, being located well away fromthe main accretion region, could have sufficiently bent fieldlines so as to reduce the duration of the time of maximumemission, while also peaking before transit. This would pro-duce a heavily skewed excess, that begins to rapidly declineafter transit as the tail rotates toward/beyond the limb. Toexplain the light curves requires that the cyclotron emissionfrom the actual X-ray pole has to be much lower than thatfrom the tail.

In producing the cyclotron model for the secondarypole, we specifically chose parameters so that the emissionwould be similar in each of the bandpasses. To explain simi-lar visual and near-IR excesses requires a field strength near80 MG, and a high optical depth. For B = 80 MG, and logΛ= 5, the n= 1 cyclotron fundamental harmonic supplies sim-ilar fluxes in all three near-IR bands, while the n = 2 and 3harmonics produce the visual excesses. Field strengths muchdifferent from 80 MG, cannot easily reproduce the observa-tions for the secondary pole as the cyclotron fundamentalmoves out of the J (B lower) or K (B higher) bandpasses.A much lower field strength near 27 MG can result in some-what similar JHK light curves, but with differing, deeperminima at transit due to this emission coming from differentharmonics in J vs. K. The excesses in the visual from sucha field are substantially weaker, as they result from higherharmonics (e.g., 7 and 8 for the V -band), and would require

at New

Mexico State U




ownloaded from



much hotter temperatures to pump them up to a level suffi-cient to explain the bluer excesses. Also note that the fieldstrength at the X-ray poles has to be slightly higher thanthose in the tail. For example, a region with B = 80 MGlocated 20◦ from the pole implies that the field strength atthe pole would be B ∼ 85 MG (see Achilleos & Wickra-masinghe 1989). Obviously, the expectation is that the tailregion would consist of a gradient of field strengths and fieldangles, allowing for a complex light curve morphology.

An alternative scenario is to suppose that there is an-other accreting pole located close to the secondary pole, butone in which the accretion rate is lower, and the cyclotronoptical depth is lower. Like the previous model, we proposethat the secondary X-ray accreting pole does not stronglyemit cyclotron emission. If the cyclotron-emitting pole islocated ∆φ = 0.18 ahead of the X-ray pole, it easily ex-plains the observational constraints. In addition, the sym-metry of the excess is better explained by such a model.There remains a slower rate of decline of the excess afterX-ray pole transit, which could be attributed to a smallamount of cyclotron emission coming from the X-ray pole.The different light curve morphologies in the various band-passes can also be more easily explained by cyclotron emis-sion from two poles with different conditions, visibilities, andfield strengths.

We do not have as tight of constraints on the phasing ofthe excesses associated with the primary pole since there isno emission above the light curve model in the near-IR, andour temporal resolution and coverage in the optical data arenot as complete. In fact, the excesses shown in Fig. 13 forBV RI are reasonably symmetric about the transit of theprimary pole. While a skew in the V -band excess in the in-terval 0.5 ≤ φ ≤ 1.0 is apparent in Fig. 12, it is smaller thenthe excess that occurs 0.5 phase later. Clearly, a core-tail ortwo accretion pole model like that evolved for secondary poleprovides a better match to the data than a single isolatedcyclotron source. The V -band excess for 2007 October 20turns on near φ = 0.55, just after the primary pole appearsfrom behind the limb. The excess disappears when the pri-mary pole becomes self-eclipsed. If the cyclotron emittingpole was offset by an identical amount from the primaryX-ray pole as that for the secondary X-ray pole, it wouldtransit at φ = 0.64, in agreement with the observed offset.

The uniformity of the excesses across the optical/IRbands suggest optically thick cyclotron emission. For thesecondary X-ray pole, using a constant−Λ cyclotron model(see Harrison & Campbell 2015), emission from a 80 MGfield could explain the lack of strong wavelength dependence,whether in an extended tail, or from an additional, discretepole. The field strength for the primary pole accretion regionmust be substantially higher than that for the secondarypole. To explain the lack of cyclotron emission in the near-IR requires field strengths of B ≥ 120 MG. The excesses inthe visual bands for the primary pole are weaker than thosefor the secondary pole, suggesting a source with higher op-tical depth, and weaker cyclotron emission. This result is inagreement with one of the early predictions from modelingthe accretion regions of polars: if either the mass accretionrate or magnetic field strength are increased, cyclotron emis-sion is suppressed vis-a-vis bremsstrahlung emission (Lamb& Masters 1979).

Schmidt et al. (1995) present results from optical spec-

tropolarimetry that show that the circular polarization re-mains positive over the phase interval 0.99 . φ . 0.56, butis negative only over the range 0.67 . φ . 0.90. At theepoch of the spectropolarimetry, the spin and orbital phasesdiffered by ∆φ = 0.06. In the geometry and synchronizedphasing that we have derived, the centres of these two po-larization ranges are very close to being co-located with thetwo X-ray poles (note that the temporal resolution of thespectropolarimetry was quite low). This than associates thepositively polarized emission with the secondary X-ray pole,and the negatively polarized emission with the primary pole.

3.4.2 The V -band Light Curves of 1995

It is worthwhile to more closely examine the older V -bandlight curves obtained by Semeniuk et al. (1995). Those data,presented in Fig. 10, were obtained on three different nights,and we have offset the individual light curves so that all oftheir maxima have similar values. We also plot the irradiatedsecondary star model that we derived above, offset by ∆V= −0.87 mag, so it is consistent with the maxima of thesedata. The result is somewhat surprising: both the pre- andpost-minima excesses seen in the recent data are present inthese earlier epoch observations. The data for 27 July closelyfollows the light curve model from phase 0.5 to phase 1.0.While on the following two days there was a large excessduring this phase interval. The excess from 0.0 ≤ φ ≤ 0.5 ismuch more stable. This is exactly the same trends observedin our more recent visual photometry. Note that the excessesin 1995 are much smaller relative to the model when com-pared to the data obtained in 2007 and beyond. This can beeasily explained by dilution since the source was a factor oftwo more luminous in 1995.

If these V -band excesses are truly tied to the accretionregions on the white dwarf as we argue above for the currenttime, the similar appearance of the older V -band data wouldsuggest that V1500 Cyg was synchronously rotating just afew years after it was shown to be asynchronous! Of course,the excesses observed in 1995 do not have to arise from thesame sources as those in the current system, since the accre-tion rate was much higher at that time. If we do, however,propagate the spin ephemeris out to the epoch of the Seme-niuk et al. observations, we find that the spin and orbitalphases differed by 0.01 on 27 July, 0.2 on 28 July, and 0.36on 29 July. Currently, we believe the main cyclotron emit-ting regions are advanced by 0.18 on the orbital period. Ifthe same mechanisms and phasing existed at the epoch ofthese earlier data, the light curves for 27 July should havecyclotron excesses centred near φOrb = 0.08, and 0.58. Thisis not in agreement with the light curve for this date. For28 July, the excesses would have been centred near φOrb =0.28, and 0.78, which is consistent with their V -band data.On 29 July the excesses would have been centred near φOrb

= 0.43, and 0.93. While there is a significant excess near φOrb

= 0.93, any excess at 0.43 is quite small. It appears unlikelythat the V -band light curves of Semeniuk et al. are consis-tent with the recent data if we assume the predicted spinperiod at the time of their observations. Either the sourceof those excesses was not located at the same sites as thecurrent emission, or the spin rate at that time was differ-ent from that predicted by the ephemeris of Schmidt et al.(1995).

at New

Mexico State U




ownloaded from



3.4.3 The U-band Light Curve

The morphology of the U -band light curve remains puz-zling. As noted earlier, each of the strongest dips in theX-ray light curve are reflected in the U -band data. This con-firms that the excess in this bandpass is due to the accretionprocess, since the response appears to be immediate. Sincewe do not have another U -band light curve, or co-temporaloptical/near-IR data, we cannot know if these data are typ-ical, or represent a peculiar state. The only features of thislight curve is the maximum at φ ∼ 0, and the minimumat φ ∼ 0.5. Neither of these phases are special in the senseof the cyclotron models evolved above, but suggest a phe-nomenon tied to the orbital period. But as the irradiatedsecondary star model plotted in Fig. 1 demonstrates, varia-tions due to the stellar components are small at this wave-length. It is interesting that the HST “FOS Optical” lightcurve, which has an effective wavelength similar to the U -band, and is plotted as Fig. 3 in Schmidt et al., is essentiallyidentical to our U -band data set, except that the phasing ofthe two light curves differ by ∆φ = 0.5 (our data is phased tothe minimum light ephemeris, theirs to the maximum lightephemeris). At the time of the HST observations, the spinand orbital phases were offset, and the light curve maximumin the FOS data corresponds to a spin phase of φspin = 0.35.The definition of the spin phase is the zero crossing of thepolarization curve, a longitude that is located halfway be-tween the two poles. For the XMM data, this correspondsto an orbital phase of φ = 0.57. A spin phase of φspin =0.35 is equivalent to a current orbital phase of φorb = 0.92.It appears the U -band maxima in both data sets are consis-tent with falling at the same “spin phases”. This remarkableresult, barring an incredible coincidence, only arises if thewhite dwarf in V1500 Cyg was asynchronous in 1992, andsynchronized in 2014.

What is the source of the U -band excess? Obviously, thelight curve could bear a considerable imprint from cyclotronemission. The relatively flat region of the differential U -bandlight curve from φ = 0, to φ = 0.3, could result from a risein the cyclotron component responsible for the optical/IRexcesses at these phases, combined with a fall in the sourcethat has its peak near φ = 0.95. Thus, it is possible that nomore than one half of the U -band excess at φ = 0 is due tothis unidentified source. This source could be the accretionstream/funnel, as it follows the main accretion pole aroundthe orbit, but precedes the donor star. This source has tobe prominent in U , but not be seen in B. Perhaps the addi-tional component of the excess is due to very strong Balmercontinuum emission.

4 DISCUSSION AND CONCLUSIONS

We have obtained a multi-wavelength, multi-epoch data setfor V1500 Cyg. The X-ray light curve is dominated by ac-cretion on to a single pole, being visible for one half of theorbit. Analysis of this light curve for periodicities indicatesthat the white dwarf now spins at a rate close to the orbitalperiod. Even with the uncertainties in the period searchingprocess, none of the periods identified in our search wereclose to the predicted spin period. Folding the X-ray datasuggests the presence of a second pole located opposite to the

primary pole that has a much lower accretion rate and/ortemperature. This result is consistent with the polarimetricobservations of Schmidt & Stockman (1991) where a sym-metric light curve in the circularly polarized flux was found,suggesting diametrically opposed accretion regions. Analysisof the X-ray data show that the spectrum and luminosity ofV1500 Cyg are consistent with a high state polar.

The multi-epoch optical/near-IR data we have compiledshows large excesses that appear to have identical phasingover a eight year span. By using the time of X-ray maxi-mum, and assuming the white dwarf rotates at the orbitalperiod, we find that these excesses are consistent with cy-clotron emission from a distributed accretion region, or frommultiple poles. While the secondary pole is quite weak atX-ray energies, the optical/IR data require an excess asso-ciated with its location. This excess is generally larger inamplitude, and less variable than the excess associated withthe primary pole. Attributing the observed excesses to cy-clotron emission leads to large field strengths for both poles:B1 = 120 MG, and B2 = 80 MG. Broadband photometry canonly provide limited insight into the nature of a cyclotronsource, but it is obvious here that the strong photometricexcesses are offset in phase relative to the location of the X-ray emitting poles. The simplest explanation is to proposecyclotron emitting regions located ∆φ = 0.18 ahead of bothpoles. These could be either a tails of emission, or nearby ac-creting poles. Distributed emission from tails, or arcs, havebeen proposed for polar accretion regions on a number ofoccasions (e.g., Buckley et al. 2000, Ramsay et al. 1996,Wickramasinghe et al. 1991), and appear to be the mostlikely explanation for the behavior of V1500 Cyg. Harrison& Campbell (2015) found several examples of polars (e.g.,VV Pup) where cyclotron emission from two poles locatedclose to each other on the white dwarf photosphere, butwith dramatically different magnetic field strengths, wererequired to explain WISE observations.

As noted earlier, if the observed spin-down rate had re-mained unchanged, the white dwarf would have had a rota-tion period at the time of the XMM observations of Pspin

∼ 11,880 s. Thus, either the original spin-down rate wasunderestimated, or there was a sudden change in the spin-down rate sometime after 1992. While the V -band data from1995 superficially resemble the more recent light curves, thespin phases of the excesses seen at that time are not con-sistent with those observed currently. Thus, either the spinephemeris was incorrect, or the emission process was differ-ent at the earlier epoch. The only other data set availableto us to help establish limits on when synchronization mighthave occurred is the first, brief XMM data set discussed inHarrison et al. (2013a). We have extracted a light curve fromthe PN data set for 2002 November 2 and overplot it (red)in Fig. 16 on top of the recent XMM PN that was data pre-sented in Fig. 1. The x-axis is now in orbital phase. To getthe two data sets to overlap requires us to shift the old databy 0.22 in phase. Obviously, we cannot know if the apparentrise to maximum in the older data set has any correlationto the behavior observed recently, but the suggestion is thatV1500 Cyg was still an asynchronous system at that time.

The rapid synchronization of V1500 Cyg appears to beat odds with existing models on how rapidly the white dwarfin such an object could be spun-down. As noted above, ifthe mass of the white dwarf was low, ≤ 0.5 M�, Camp-

at New

Mexico State U




ownloaded from



0 0.5 1 1.5 2

0

1

2

3

4

Figure 16. The XMM PN light curve for 2002 November 2 (red)overlaid on the recent PN data from Fig. 1 (black), plotted vs.

orbital phase. To get the older data set to align with the more

recent data required an offset of ∆φ = −0.22. The 2002 data werebinned by 60 s, whereas the 2014 data are binned by 120 s. The

older PN data were normalized so as to best-match the recent

light curve.

bell & Schwope (1999) predicted synchronization could oc-cur within 50 years by employing a torque that resultedfrom the currents and electric fields induced by the asyn-chronous rotation. That prediction is sufficiently close tobe in agreement with our results. But given the violenceof the CNe eruption, the probability of an ONeMg whitedwarf primary, and the high X-ray shock temperatures, theprimary in V1500 Cyg almost certainly has a mass that ismuch greater than 0.5 M�. Thus, their mechanism mightseem insufficient to explain the current observations. How-ever, the time scale they derive for synchronization onlyincreases linearly with white dwarf mass, while being in-versely quadratic in magnetic field strength. If the magneticfield strength of the primary pole is truly near 120 MG, thecanonical synchronization time scale is shortened by a factorof nine. Thus, the torque mechanism detailed by Campbell& Schwope would be sufficient to explain the current stateof V1500 Cyg.

V1500 Cyg remains a difficult source to fully under-stand. It is obvious that there is a hot white dwarf thatirradiates a low mass star. The X-ray data clearly indicatea system whose spin rate is much reduced from expecta-tions, and suggest it is close to synchronization. The recentground-based data are also consistent with that interpreta-tion, and suggest it has been synchronized since 2006/2007.However, it remains difficult to fully characterize the sourceof the optical/IR excesses. If the excesses are due to cy-clotron emission, then they require either a core-tail, or amultipole geometry. A cyclotron interpretation would alsoargue for high magnetic field strengths for both poles. Per-haps it is these larger field strengths that aided in the morerapid spin-down observed for V1500 Cyg, when compared

to other asynchronous polars such as BY Cam. Harrison &Campbell (2015) found that BY Cam had very large am-plitude variations in the WISE bandpasses, and suggestedthat this was due to cyclotron emission from a field witha strength of B ≤ 28 MG. Additional X-ray observationson several closely spaced epochs should provide for a morerobust periodicity determination for V1500 Cyg. It also re-mains possible that discrete cyclotron harmonics might beobserved from the secondary pole region, since the excessesassociated with that pole are large, and do not appear toexhibit the same variations seen in the excesses that corre-spond to the visibility of primary pole. Such observationswill be required to better constrain the field strength, andgeometries of the cyclotron emitting regions.

ACKNOWLEDGEMENTS

This was unfunded research. Based on observations obtainedwith XMM-Newton, an ESA science mission with instru-ments and contributions directly funded by ESA MemberStates and NASA. We would like to thank our anonymousreferee for useful comments.

References

Achilleos N., Wickramasinghe D.T., 1989, ApJ, 346, 444Barman T.S., Hauschildt P.H., Allard F., 2004, ApJ, 614,338Beuermann K., El Kholy E., Reinsch K., 2008, A&A, 481,771Beuermann K., Schwope A. D., 1989, A&A, 223, 179Beuermann K., Stella L., Patterson J., 1987, ApJ, 316, 360Brett J.M. Smith R.C., 1993, MNRAS, 264, 641Buckley D.A.H., Cropper M., van der Hayden K., PotterS.B., Wickramasinghe D.T., 2000, MNRAS, 318, 187Campbell C.G. Schwope A.D., 1999, A&A, 343, 132Carter J.A., Read A.M., 2007, A&A, 464, 1155Christian D.J., 2000, AJ, 119, 1930Collazzi A.C., Schaefer B.E., Xiao L., Pagnota A., Kroll P.,Lochel K., Henden, A.A., 2009, AJ, 138, 1846Cropper M., 1988, MNRAS, 231, 597Cropper M., 1986, MNRAS, 222, 853Harrison T.E., Campbell R.D., 2015, ApJS, 219, 32Harrison T.E. 2014, ApJ, 791, L18Harrison T.E., Campbell, R.D., Lyke J.E., 2013a, AJ, 146,37Harrison T.E., Bornak J., McArthur B.E., Benedict G.F.2013b, ApJ, 767, 7Harrison T.E., Hamilton R.T., Tappert C., Hoffman D.I.,Campbell R.K., 2013c, AJ, 145, 19Horne K., Schneider D.P., 1989, ApJ, 343, 888Ishida M., Matsuzaki K., Fujimoto R., Mukai K., Osborne,J.P., 1997, MNRAS, 287, 651Kaluzny J., Semeniuk I., 1987, ActaA, 37, 349Lance C.M., McCall M.L., Uomoto A.K., 1988, ApJS, 66,151Lamb D.Q., Masters A.R., 1979, ApJ, 234, L117Landi R., Bassani L., Dean A.J., Bird A.J., Fiocchi M.,Bazzano A., Nousek J.A., Osborne, J.P., 2009, MNRAS,392, 630Lenz P., Breger M., 2005, Commun. Asteroseismology, 146,53

at New

Mexico State U




ownloaded from



Lumb D.H., Schartel N., Jansen F.A., 2012, OptEn, 51,1009Patterson J., 1979, ApJ, 231, 789Politano M., Starrfield S., Truran J. W., Weiss A., SparksW.M., 1995, ApJ, 448, 807Ramsay G., Cropper M., Wu K., Potter S., 1996, MNRAS,282, 726Ramsay G., Cropper M., 2003, MNRAS, 338, 219Schmidt G.D., Liebert J., Stockman H.S., 1995, ApJ, 441,414Schmidt G.D., Stockman H.S., 1991, ApJ, 371, 749Schwope A.D., et al., 1998, in Wild Stars in the Old West,ed. S. Howell, E. Kuulkers, & C. Woodward, ASP Conf.Ser., 137, 44Semeniuk I., Olech A., Nalezyty M., 1995, ActaAst, 45, 747Semeniuk I., Kruszewski A., Schwarzenberg-Czerny A.,Chlebowski T., Mikolajewski M., Wolcyzk J., 1977, Ac-taAst, 27, 301Slavin A.J., O’Brien T.J., Dunlop J.S., 1995, MNRAS, 276,353Stockman H.S., Schmidt G.D., Lamb D.Q., 1988, ApJ, 332,282Warner B., 1995, Cataclysmic Variable Stars (Cambridge:Cambridge Univ. Press)Wickramasinghe D.T., Bailey J.A., Meggitt S.M.A., Fer-rario L., Hough J., Touhy I.R., 1991, MNRAS, 251, 28Wolff M.T., Wood K.S., Imamura J.N. 1999, ApJ, 526, 435

at New

Mexico State U




ownloaded from


the apparent synchronization of v1500...

Documents