Photometric Study of KR Aurigae during the High State in 2001 (original) (raw)

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Accepted:

17 September 2002

Published:

25 December 2002

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Taichi Kato, Ryoko Ishioka, Makoto Uemura, Photometric Study of KR Aurigae during the High State in 2001, Publications of the Astronomical Society of Japan, Volume 54, Issue 6, 25 December 2002, Pages 1033–1039, https://doi.org/10.1093/pasj/54.6.1033
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Abstract

We photometrically observed the VY Scl-type cataclysmic variable KR Aurigae after its final rise from a fading episode in 2000–2001. A time-resolved observation revealed that the light curve is dominated by a persistent short-term variation with time-scales of minutes to tens of minutes. On some nights, quasi-periodic variations with periods of 10–15 min were observed. No coherent variation was detected. The power spectral density of the variation has a power-law component (⁠|$f^{-1.63}$|⁠). The temporal properties of the short-term variations in KR Aur present additional support for the possibility that flickering in CVs may be better understood as the result of a self-organized critical state, as in black-hole candidates. The light curve lacks “superhump”-type signals, which are relatively frequently seen in VY Scl-type systems and which are suggested to arise from a tidal instability of the accretion disk induced by changing mass-transfer rates. The present observation suggests a borderline of superhump excitation in VY Scl-type stars between mass ratios of |$q = 0.43$| (MV Lyr) and |$q = 0.60$| (KR Aur).

1. Introduction

Cataclysmic variables (CVs) are close binary systems consisting of a white dwarf and a red-dwarf secondary transferring matter via the Roche lobe overflow. The resultant accretion disk becomes thermally stable in systems with high mass-transfer rates (⁠|$\dot{M}$|⁠). Such systems are called nova-like (NL) stars, because they lack outbursts characteristic of dwarf novae [see Osaki (1996) for a review]. Among NL stars, there exists a small group which shows a temporary reduction or cessation of |$\dot{M}$| from the secondary. These systems are called VY Scl-type stars or anti-dwarf novae (Warner 1995).

Although accretion disks in NL stars are thermally stable, the disk can be tidally unstable (Osaki 1995, 1996). Numerical simulations have shown that this instability (tidal instability: Whitehurst 1988) only appears below a certain mass ratio |$(q = M_2/M_1): q < 0.25 \hbox{--} 0.33$| depending on the calculations (Hirose, Osaki 1990; Wood et al. 2000; Whitehurst, King 1991; Molnar, Kobulnicky 1992; Murray 1998). In recent years, several systems above this stability limit are known to show superhumps (Patterson 1999). Since many of them are VY Scl-type stars, there has been a theoretical interpretation that accretion disks can be tidally unstable upon the variation of |$\dot{M}$| (Murray et al. 2000), even in intermediate |$q$| systems. The temporary appearance of superhump signals in a recent low state of a VY Scl-type star, BZ Cam (Kato, Uemura 2001) may support this interpretation. However, observations have not yet fully illustrated the upper |$q$| limit for superhumpers in VY Scl-type stars. KR Aur (⁠|$q = 0.60$|⁠) is an ideal system to examine such a condition, since it has the smallest |$q$| among VY Scl-type stars which have not been reported to show (or studied for) superhumps.

2. Past Studies of KR Aurigae

KR Aur is a variable star originally discovered by M. Popova in 1960. Popova originally classified the object as a ‘unusual nova’ (cf. Duerbeck 1987). The object has received much attention since Popova and Vitrichenko (1977) claimed, from the detection of a large infall velocity of |$3200 \,\mathrm{km} \,\mathrm{s}^{-1}$| in its |$\mathrm{H}\beta$| line, that the object may be a solitary black hole accreting interstellar matter. Popova and Vitrichenko (1978) further reported a long-term light curve, which could not fit any existing classifications of variable stars, a large radial velocity (⁠|$+130 \,\mathrm{km} \,\mathrm{s}^{-1}$|⁠), a steep Balmer decrement, and a spectral energy distribution approximated by |$F_{\nu} =$| const. Popova and Vitrichenko (1978) considered that these features are also consistent with an accreting black hole. However, Doroshenko et al. (1977) reported that the time-scales of a continuum variation (a few hours to a few tens of seconds) and the |$UBV$| colors resemble those of CVs, known as “nova-like” stars at the time of the reporting.

Doroshenko and Terebizh (1978) studied short-term light variation, and detected |$\sim 0.4 \,\mathrm{mag}$| variations with time-scales of several minutes. Liller (1980) studied archival plates, and obtained the first long-term light curve of this object. Liller (1980) showed that KR Aur sometimes shows excursions to faint states [which was also noted by Popova and Vitrichenko (1978)], but no clear periodicity. Popov (1982) reported an ongoing new fading episode.

The nature of the object was revealed with spectroscopic observations by Shafter (1983) and Hutchings et al. (1983), who detected a binary motion from a radial-velocity study. The object was thus confirmed to be a CV. Williams (1983) showed that the object shows a spectrum typical of a CV with relatively weak emission lines. However, since the characteristics of light variations did not fit any known subclass of CVs, long-term light variations were subsequently studied in detail, mainly from photographic survey materials (Gotz 1983, 1984, 1985, 1986, 1987, 1988, 1989, 1990; Fujino et al. 1987; Liller, Popova 1984; Popov et al. 1994; Antov et al. 1996). All of these observations showed rather erratic light curves, and occasional fadings with time-scales of months to a few years. From these characteristics, KR Aur is now considered to belong to VY Scl-type nova-like CVs, which occasionally show fading episodes (cf. Warner, van Citters 1974; Garnavich, Szkody 1988; Greiner 1998; Leach et al. 1999). KR Aur was also independently selected as an ultraviolet-excess object (KUV 06126|$+$|2836: Kondo et al. 1984), for which Wegner and Swanson (1990) gave a spectroscopic classification as a CV.

Popov and Antov (1989) reported rapid variations in KR Aur. Biryukov and Borisov (1990) further studied this star and reported the presence of a 25-min periodicity. The light curve presented by Biryukov and Borisov (1990) suggests a moderate degree of coherence.

3. Observations

The observations were mainly performed using an unfiltered ST-7E camera attached to 25-cm/30-cm Schmidt–Cassegrain telescopes at Kyoto University. A single snapshot observation during a low state was taken on 1998 December 21, using an unfiltered ST-7 camera attached to a 25-cm Schmidt–Cassegrain telescope. All systems gave magnitudes close to |$R_\mathrm{c}$|⁠. The images were dark-subtracted, flat-fielded, and analyzed using a Java|$^\mathrm{TM}$|-based aperture photometry package developed by one of the authors (TK). The differential magnitudes of the variable were measured against GSC 1889.700 (Tycho-2 magnitude |$V = 10.27, B-V = +0.88$|⁠), whose constancy during the run was confirmed by a comparison with fainter check stars in the same field. The log of the observations is summarized in table 1. The total number of useful frames was 3139. Barycentric corrections were applied before the following analysis.

Table 1

Log of observations.

| Date | BJD * (start–end) | |$N$| † | Magnitude ‡ | Error § | Instrument || | ------------------------------------- | ------------------------------------------------ | --------------------------------------------------- | ----------------------------------------------- | ------------------------------------------------------ | | 3 | 6.408 | 0.036 | 1 | | 428 | 3.466 | 0.006 | 2 | | 69 | 3.464 | 0.010 | 2 | | 270 | 3.526 | 0.006 | 2 | | 252 | 3.376 | 0.008 | 2 | | 142 | 3.369 | 0.008 | 2 | | 259 | 3.508 | 0.008 | 2 | | 355 | 3.427 | 0.007 | 2 | | 102 | 3.320 | 0.018 | 2 | | 281 | 3.595 | 0.010 | 3 | | 359 | 3.465 | 0.006 | 2 | | 61 | 3.285 | 0.011 | 3 | | 200 | 3.455 | 0.010 | 2 | | 51 | 3.460 | 0.017 | 3 | | 53 | 3.614 | 0.015 | 3 | | 94 | 3.308 | 0.022 | 3 | | 91 | 3.229 | 0.018 | 3 | | 69 | 3.272 | 0.009 | 3 |

| Date | BJD * (start–end) | |$N$| † | Magnitude ‡ | Error § | Instrument || | | ---------------- | ---------------------------------------------------------- | ------------------------------------------------ | --------------------------------------------------- | ----------------------------------------------- | ------------------------------------------------------ | | 1998 December 21 | 51169.248–51169.248 | 3 | 6.408 | 0.036 | 1 | | 2001 November 23 | 52237.059–52237.244 | 428 | 3.466 | 0.006 | 2 | | 2001 November 24 | 52238.082–52238.111 | 69 | 3.464 | 0.010 | 2 | | 2001 November 25 | 52239.055–52239.172 | 270 | 3.526 | 0.006 | 2 | | 2001 November 26 | 52240.042–52240.177 | 252 | 3.376 | 0.008 | 2 | | 2001 November 27 | 52241.041–52241.105 | 142 | 3.369 | 0.008 | 2 | | 2001 November 28 | 52242.029–52242.152 | 259 | 3.508 | 0.008 | 2 | | 2001 November 30 | 52244.149–52244.352 | 355 | 3.427 | 0.007 | 2 | | 2001 December 1 | 52245.050–52245.126 | 102 | 3.320 | 0.018 | 2 | | 2001 December 4 | 52248.211–52248.346 | 281 | 3.595 | 0.010 | 3 | | 2001 December 6 | 52250.007–52250.168 | 359 | 3.465 | 0.006 | 2 | | 2001 December 7 | 52251.116–52251.152 | 61 | 3.285 | 0.011 | 3 | | 2001 December 11 | 52255.003–52255.166 | 200 | 3.455 | 0.010 | 2 | | 2001 December 13 | 52257.104–52257.137 | 51 | 3.460 | 0.017 | 3 | | 2001 December 14 | 52258.109–52258.134 | 53 | 3.614 | 0.015 | 3 | | 2001 December 15 | 52259.261–52259.073 | 94 | 3.308 | 0.022 | 3 | | 2001 December 16 | 52260.073–52260.128 | 91 | 3.229 | 0.018 | 3 | | 2001 December 17 | 52261.092–52261.126 | 69 | 3.272 | 0.009 | 3 |

BJD|$-$|2400000.

†

Number of frames.

‡

Averaged magnitude relative to GSC1889.700.

Standard error of the averaged magnitude.

1: Kyoto (25-cm|$+$|ST-7), 2: Kyoto (25-cm|$+$|ST-7E), 3: Kyoto (30-cm|$+$|ST-7E).

Table 1

Log of observations.

BJD|$-$|2400000.

†

Number of frames.

‡

Averaged magnitude relative to GSC1889.700.

Standard error of the averaged magnitude.

1: Kyoto (25-cm|$+$|ST-7), 2: Kyoto (25-cm|$+$|ST-7E), 3: Kyoto (30-cm|$+$|ST-7E).

4. Results and Discussions

4.1. Long-Term Variation

Figure 1 shows the long-term light variation of KR Aur between 1994 and 2002 drawn from visual observations reported to the VSNET Collaboration.1 Typical errors of visual observations were 0.2–0.3 mag, which would not affect the following analysis. The light curve shows predominant low states (durations: several months to two years) when the object is fainter than 14 mag, and occasional high states when the object is typically between 13 mag and 14 mag. Such a high occurrence of low states is relatively rare among VY Scl-type stars (Greiner 1998). The long-term behavior of the system most resembles “superminimum” of MV Lyr (Robinson et al. 1981; Wenzel, Fuhrmann 1983; Fuhrmann 1985). At a closer look, the system underwent a small dwarf-nova-like outburst around JD 2450745–2450755. The object then became inactive. Before a full recovery, the object underwent a rather complex brightening around JD 2451580–2451630. This behavior is quite reminiscent of the fading and recovery processes observed in MV Lyr (Pavlenko 1998; Shugarov, Pavlenko 1998; Pavlenko, Shugarov 1999). Although such behavior is unexpected for variable mass-transfer rates on a usual CV (Honeycutt et al. 1994), Leach et al. (1999) showed that, in the presence of heating from a hot white dwarf, the irradiation on the accretion disk suppresses the thermal instability, which can reproduce the observed light curve of VY Scl-type systems. From these observations, we propose that MV Lyr and KR Aur comprise the most “active” subgroup of VY Scl-type stars.

Fig. 1

Long-term light variation of KR Aur between 1994 and 2002 drawn from visual observations reported to VSNET. The large dots and small dots represent positive and negative (upper limit: object undetected) observations, respectively. Typical errors of visual observations are 0.2–0.3 mag. The open circle represents our snapshot CCD observations. The mean epoch of our time-series CCD photometry is shown by the arrow.

Figure 2 presents the overall light variation of KR Aur in 2001 November–December, drawn from observations in table 1. Although some irregular variation is superimposed, no major “flare”-like brightening (cf. Popova, Vitrichenko 1978) was observed.

Overall light variation of KR Aur in 2001 November–December.

Fig. 2

Overall light variation of KR Aur in 2001 November–December.

4.2. Superhumps

Figure 3 presents a period analysis of the November 23–December 17 data, after subtracting a linear fit to the overall light curve. The upper and lower panels represent the results of a period analysis with Phase Dispersion Minimization (Stellingwerf 1978) and with the CLEAN algorithm (Roberts et al. 1987), respectively. The orbital period is marked with a tick on the upper panel. No significant coherent periodicity was detected at or near the orbital period, indicating that no detectable periodic or quasi-periodic variations related to orbital modulations or superhumps were present (see also).

Period analysis of the November 23–December 17 data. The upper and lower panels represent the results of a period analysis with Phase Dispersion Minimization (Stellingwerf 1978) and with the CLEAN algorithm (Roberts et al. 1987), respectively. The orbital period is marked with a tick on the upper panel. No significant coherent periodicity was detected at or near the orbital period.

Fig. 3

The absence of superhumps makes a clear contrast to other VY Scl-type stars with superhumps [TT Ari: Udalski (1988); Skillman et al. (1998); Andronov et al. (1999); Kraicheva et al. (1999); Stanishev et al. (2001), MV Lyr: Skillman et al. (1995), V751 Cyg: Patterson et al. (2001)]. This finding suggests that the mechanism proposed by Murray et al. (2000) is ineffective in the mass ratio of KR Aur. We propose that there is a borderline of superhump excitation between mass ratios |$q = 0.43$| (MV Lyr) and |$q = 0.60$| (KR Aur) in VY Scl-type stars. A future determination of the mass ratios in longer |$P_\mathrm{orb}$| systems (i.e. candidates for systems with higher |$q$| than MV Lyr; such as BZ Cam, V751 Cyg, and TT Ari) is expected to provide a more stringent constraint to this limit (see table 2 for a summary of the superhumps and binary parameters of VY Scl-type stars).

Table 2

Parameters of VY Scl-type stars. *

| Object | |$P_\mathrm{orb}$| (d) | |$P_\mathrm{SH}$| (d) | |$M_1$| | |$q$| | | -------------------------------------------------- | ------------------------- | ------------------------ | ------------ | ------------ | | VY Scl | 0.232 | |$\cdots$| | 1.22 | 0.32 | | KR Aur | 0.16280 | |$\cdots$| | 0.59 | 0.60 | | LX Ser † | 0.158432 | |$\cdots$| | 0.41 | 0.88 | | BZ Cam | 0.153693 | 0.15634 | |$\cdots$| | |$\cdots$| | | V794 Aql | 0.1533 | |$\cdots$| | |$\cdots$| | |$\cdots$| | | V425 Cas | 0.1496 | |$\cdots$| | |$\cdots$| | |$\cdots$| | | VZ Scl | 0.144622 | |$\cdots$| | 1:9 | 0.7: | | V751 Cyg | 0.144464 | 0.1394 | |$\cdots$| | |$\cdots$| | | PG 1000|$+$|667 | 0.144384 | |$\cdots$| | |$\cdots$| | |$\cdots$| | | TT Ari | 0.137551 | 0.133160 | |$\cdots$| | |$\cdots$| | | 0.148815 | | | | | | DW UMa † | 0.136607 | 0.1330 | 0.9 | 0.32 | | MV Lyr | 0.1329 | 0.138 | |$\cdots$| | 0.43 | | V442 Oph † | 0.124330 | 0.12090 | |$\cdots$| | |$\cdots$| |

References: VY Scl: Martínez-Pais et al. (2000); KR Aur: Shafter (1983); LX Ser: Young et al. (1981); BZ Cam: Patterson et al. (1996), Kato and Uemura (2001); V794 Aql: Honeycutt and Robertson (1998); V425 Cas: see Kato et al. (2001) and the references therein; VZ Scl: Warner and Thackeray (1975), Robinson (1976), O’Donoghue et al. (1987), Sherrington et al. (1984); V751 Cyg: Patterson et al. (2001); PG 1000|$+$|667: Hillwig et al. (1998); TT Ari: Cowley et al. (1975), Thorstensen et al. (1985), Udalski (1988), Skillman et al. (1998), Andronov et al. (1999), Kraicheva et al. (1999), Stanishev et al. (2001); DW UMa: Shafter et al. (1988), Patterson (1999), Bíró (2000); MV Lyr: Skillman et al. (1995); V442 Oph: Patterson (1999), Hoard et al. (2000), Diaz (2001).

LQ Peg (Kato, Uemura 1999) is also known as a VY Scl-type star. The orbital period has not been reported.

†

SW Sex star (see Hellier 2000 for a recent review; see alsoThorstensen et al. 1991). A few other SW Sex stars [PX And (Still et al. 1995); BH Lyn (Hoard, Szkody 1997)] are claimed to show some degree of low/high state transitions (mainly from spectroscopic observations). Such a variation may more represent a change in the disk state (cf.Groot et al. 2001) rather than VY Scl-type temporary reduction or cessation of |$\dot{M}$| from the secondary. Some DQ Her stars (intermediate polars) also show temporary fadings (Garnavich, Szkody 1988; Hessman 2000), but they are not usually considered as VY Scl-type stars.

Table 2

Parameters of VY Scl-type stars. *

LQ Peg (Kato, Uemura 1999) is also known as a VY Scl-type star. The orbital period has not been reported.

†

4.3. Short-Term Variations

Figure 4 shows typical examples of nightly light curves. Three longest runs (November 23 and 30, December 6) were selected as representatives of different epochs of the present observation. All light curves show distinct short-term (minutes to tens of minutes) quasi-periodic variations. No systematic variation close to the orbital period (0.16280 d) was observed (cf. subsection 4.2). Aside from the lack of superhumps, these variations look remarkably similar to those of quasi-periodic oscillations (QPOs) observed in another VY Scl-type star, TT Ari (e.g. Mardirossian et al. 1980; Jensen et al. 1983; Hollander, van Paradijs 1992; Tremko et al. 1996; Andronov et al. 1999).

Representative nightly light curves at different epochs. All light curves show distinct short-term (minutes to tens of minutes) quasi-periodic variations. No systematic variation close to the orbital period (0.16280d) was observed (cf.

Fig. 4

Figure 5 shows power spectra of nightly observations. On some nights (November 23, 30, December 11), an increased power around frequencies of |$100 \hbox{--} 150 \,\mathrm{d}^{-1}$| (corresponding to the periods of 10–15 min) was observed. There was no common period between the nightly observations. This finding confirms the quasi-periodic nature of the short-term variations. We have not been able to confirm the presence of a 25-min periodicity claimed by Biryukov and Borisov (1990). The periods of the presently observed QPOs are close to the periodicities (480–780 s) recorded by Singh et al. (1993). We also have confirmed a night-to-night variation of dominant periods, as was originally claimed by Singh et al. (1993).

$Power spectra of nightly observations. On some nights (November 23, 30, December 11), increased power around the frequencies <math xmlns="http://www.w3.org/1998/Math/MathML"><semantics><mrow><mn>100</mn><mrow><mtext>–</mtext></mrow><mn>150</mn><mtext> </mtext><msup><mi mathvariant="normal">d</mi><mrow><mo>−</mo><mn>1</mn></mrow></msup></mrow><annotation encoding="application/x-tex">100 \hbox{--} 150 \,\mathrm{d}^{-1}</annotation></semantics></math>100–150d−1 (corresponding to the periods of 10–15 min) was observed. There was no common period between the nightly observations.$

Fig. 5

Power spectra of nightly observations. On some nights (November 23, 30, December 11), increased power around the frequencies |$100 \hbox{--} 150 \,\mathrm{d}^{-1}$| (corresponding to the periods of 10–15 min) was observed. There was no common period between the nightly observations.

We also note that the overall profiles of short-term variations (figure 4) and the nightly variation of the power spectra are also similar to those observed in the peculiar symbiotic binary V694 Mon (Michalitsianos et al. 1993; Dobrzycka et al. 1996; Ishioka et al. 2001).

There are two major types of “quasi-periodic” oscillations in CVs: dwarf nova oscillations (DNOs), which are oscillations observed only in dwarf nova outbursts, having periods of 19–29 s, and have long (several tens to |$\sim 100$| wave numbers) coherence times (Robinson 1973; Szkody 1976; Hildebrand et al. 1980), and (in a narrower sense) QPOs (for a review, see Warner 1995). The present QPOs in KR Aur correspond to the latter classification. Several models have been proposed to account for QPOs, including vertical or radial oscillations of the accretion disk (Kato 1978), reprocessing of the light by the orbiting blobs (Patterson 1979), non-radial pulsations of the accretion disk (Papaloizou, Pringle 1978; van Horn et al. 1980), radial oscillation of the accretion disk (Cox 1981; Blumenthal et al. 1984; Okuda, Mineshige 1991; Okuda et al. 1992), excitation of trapped oscillations around the discontinuity of physical parameters (Yamasaki et al. 1995). There has also been a suggestion that the magnetism of the white dwarf can be responsible for some kinds of QPOs (Mikolajewski et al. 1990). Although the present observations are not able to constrain the origin of QPOs, a future search for coherent X-ray, ultraviolet or optical pulsations, which are a well-known signature of a magnetic white dwarf, would be helpful in discriminating the possibilities.

4.4. Power Spectrum of Flickering

Figure 6 shows the power spectral density (PSD) of the entire data set. High-frequency white noise has been subtracted from the PSD. Above the frequency |$\log f = 1.5 \,\mathrm{d}^{-1}$| (corresponding to time-scales shorter than |$\sim 45 \,\mathrm{min}$|⁠), the PSD is proportional to |$f^{-1.63}$|⁠.

$Power spectral density (PSD) of the entire data set. High-frequency white noise has been subtracted from the PSD. Above the frequency <math xmlns="http://www.w3.org/1998/Math/MathML"><semantics><mrow><mi>log</mi><mo>⁡</mo><mi>f</mi><mo stretchy="false">(</mo><msup><mi mathvariant="normal">d</mi><mrow><mo>−</mo><mn>1</mn></mrow></msup><mo stretchy="false">)</mo><mo>=</mo><mn>1.5</mn></mrow><annotation encoding="application/x-tex">\log f (\mathrm{d}^{-1})=1.5</annotation></semantics></math>logf(d−1)=1.5, the PSD is proportional to <math xmlns="http://www.w3.org/1998/Math/MathML"><semantics><mrow><msup><mi>f</mi><mrow><mo>−</mo><mn>1.63</mn></mrow></msup></mrow><annotation encoding="application/x-tex">f^{-1.63}</annotation></semantics></math>f−1.63.$

Fig. 6

Power spectral density (PSD) of the entire data set. High-frequency white noise has been subtracted from the PSD. Above the frequency |$\log f (\mathrm{d}^{-1})=1.5$|⁠, the PSD is proportional to |$f^{-1.63}$|⁠.

These short-term variations (flickering) are one of the most characteristic features in CVs. Although the existence of flickering in CVs has been well-documented since the 1940’s (see Bruch 1992 for an extensive historical review), their physical origin has not been well understood. Historically, Warner and Nather (1971) demonstrated that flickering almost disappeared during eclipses of the eclipsing dwarf nova U Gem. This finding indicated that the origin of flickering is strongly associated with the hot spot (the stream-impact point on the accretion disk). James (1987) noted the presence of a power-law-type frequency-dependence, and proposed that a multiple scattering from the flickering source (hot spot) is responsible for this frequency-dependence. More recently, Bruch (1991) more extensively studied the properties of flickering, and summarized an observational review (Bruch 1992). From an analysis of the frequency and color dependencies, Bruch (1992) and Bruch and Duschl (1993) suggested that the inner part of the accretion disk (rather than the hot spot) is more responsible for flickering. More direct observational evidence for a major contribution from the inner accretion disk to flickering has been demonstrated through eclipse observations of CVs (HT Cas: Welsh, Wood 1995; Bruch 2000, Z Cha: Bruch 1996). It is now widely believed that the originally proposed idea of stream-impact-type flickering (Warner, Nather 1971) applies to a only limited sample of CVs, or contributes to a small extent to overall flickering activity, and that most of CVs have a strong concentration of flickering activity toward the inner accretion disk (Bruch 1996).

In order to reproduce this power-law spectrum, Yonehara et al. (1997) proposed a superposition of “shots” in a self-organized critical state (SOC), which was originally introduced to explain time variations in black-hole candidates (BHCs) (e.g. Mineshige et al. 1994a,b; Takeuchi et al. 1995; Kawaguchi et al. 2000). Taking this analogy into account, the power-law-type temporal properties of the short-term variations in KR Aur present additional support for the possibility that flickering in CVs may be better understood as a result of SOC, as in BHCs. Although the detailed mechanism of energy release was not identified at the time of Yonehara et al. (1997), Williams and Maletesta (2002) recently tried to explain flickering by a superposition of flares, resulting from an energy release in the photosphere of the accretion disk of injected high-energy electrons originating from reconnections of magnetic field lines. This type of energy release in visual wavelengths would be a promising candidate for explaining flickering in CVs. Further quantitative comparisons with numerical simulations and observed properties in CVs will be a next step toward understanding flickering in CVs.

5. Summary

We photometrically observed the VY Scl-type cataclysmic variable KR Aurigae after its final rise from a fading episode in 2000–2001. We show that the long-term light curve of KR Aur is densely populated with low states, making KR Aur an exceptionally active VY Scl-type star. The object showed a complex recovery process from a faint state, which may be understood as being the result of heating on the accretion disk from a hot white dwarf. A time-resolved observation during the high state revealed that the light curve is dominated by a persistent short-term variation with time scales of minutes to tens of minutes. On some nights, quasi-periodic variations with periods of 10–15 min were observed. No coherent variation was detected. The power spectral density of the variation has a power-law component (⁠|$f^{-1.63}$|⁠) at high frequencies. The temporal properties of the short-term variations in KR Aur present additional support for the possibility that flickering in CVs may be better understood as being the result of a self-organized critical state, as in black-hole candidates. Contrary to large-amplitude short-term variations, the light curve lacks “superhump”-type signals, which are seen relatively frequently in VY Scl-type systems. The present observation suggests a borderline of superhump excitation in VY Scl-type stars between mass ratios |$q = 0.43$| (MV Lyr) and |$q = 0.60$| (KR Aur).

The authors are grateful to observers who reported vital observations to VSNET. This work is partly supported by a grant-in aid (13640239) from the Japanese Ministry of Education, Culture, Sports, Science and Technology. Part of this work is supported by a Research Fellowship of the Japan Society for the Promotion of Science for Young Scientists (MU).

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February 2018	14
March 2018	14
April 2018	9
May 2018	12
June 2018	13
July 2018	19
August 2018	10
September 2018	10
October 2018	7
November 2018	13
December 2018	25
January 2019	8
February 2019	11
March 2019	21
April 2019	20
May 2019	10
June 2019	11
July 2019	12
August 2019	11
September 2019	5
October 2019	8
November 2019	7
December 2019	8
January 2020	9
February 2020	5
March 2020	4
April 2020	3
May 2020	1
June 2020	10
July 2020	3
August 2020	6
September 2020	13
October 2020	9
November 2020	8
December 2020	6
January 2021	6
February 2021	4
March 2021	5
April 2021	1
May 2021	2
July 2021	5
August 2021	2
September 2021	5
October 2021	9
November 2021	9
December 2021	1
February 2022	6
March 2022	2
April 2022	2
May 2022	2
June 2022	2
July 2022	7
August 2022	7
September 2022	9
October 2022	6
November 2022	5
December 2022	11
January 2023	19
March 2023	1
April 2023	6
May 2023	2
June 2023	3
August 2023	3
September 2023	3
October 2023	2
November 2023	3
December 2023	11
January 2024	6
February 2024	5
March 2024	3
April 2024	4
May 2024	6
June 2024	7
July 2024	15
August 2024	2
September 2024	7

Citations

20 Web of Science

Photometric Study of KR Aurigae during the High State in 2001 (original) (raw)

Cite

Abstract

1. Introduction

2. Past Studies of KR Aurigae

3. Observations