PTF1 J191905.19+481506.2—A PARTIALLY ECLIPSING AM CVn SYSTEM DISCOVERED IN THE PALOMAR TRANSIENT FACTORY

AM CVn systems are rare, ultra-compact, semi-detached, white dwarf binaries with periods ranging from 5 to 65 minutes. First identified over 40 yr ago by Smak (1967), it has only been in the last decade that the number of known systems has risen above 10. Yet despite the recent discovery of almost 25 additional systems, their rich and complex phenomenological behavior has limited our understanding of this class of post-period minimum binaries. They are considered to be the helium analog of cataclysmic variables (CVs) and are important as strong, low-frequency Galactic gravitational wave sources (Nelemans et al. 2004; Roelofs et al. 2007b; Nissanke et al. 2012), the source of the proposed "Ia" supernovae (Bildsten et al. 2007), and one of the believed end-points of binary white dwarf evolution (Nelemans et al. 2001). However, many questions about these unique systems remain, including their population density, their evolutionary pathways, and the interactions between the two components, including the He-rich accretion disk. We refer the reader to Nelemans (2005) and Solheim (2010) for general reviews.

Follow-up study of a newly identify AM CVn system typically begins with a spectroscopic determination of its orbital period using phase-resolved spectroscopy of the "hot spot" (Nather et al. 1981), which is the location where matter from the donor hits the accretion disk. However, these measurements are incomplete, as the inclination of the systems cannot be determined. Only those systems showing eclipses and radial velocity variations of two components (e.g., the donor and accretor or the accretor and hot spot) allow for a full determination of the system's parameters (e.g., component masses, inclination, period, etc.). Eclipses also allow for the extremely precise measurement of the orbital period change. For example, the orbital period of the double detached white-dwarf binary SDSSJ0651+2844 was measured to be decreasing by Hermes et al. (2012) while the only known eclipsing AM CVn system, SDSSJ0926+3624 (Anderson et al. 2005; Copperwheat et al. 2011, hereafter C11), was recently shown to have an increasing orbital period (Szypryt et al. 2014).

AM CVn systems have been observed to have three relatively distinct phenomenological states. Systems with orbital periods (hereafter P_orb) less than 20 minutes have been observed in a "high" state, characterized by optically thick accretion disks and absorption-line spectra. In contrast, systems with P_orb > 40 minutes are observed to have emission-line spectra from what are believed to be optically thin accretion disks and are said to be in quiescence (though see Woudt et al. 2013 for an example of a longer period system that has been observed to outburst). Systems in these two classes do not show photometric variability over 0.5 mag.

Between these orbital period limits are outbursting systems, which are observed to have dwarf nova-type outbursts of 3–6 mag as well as variability at the 10% level in both quiescence and outburst (e.g., Patterson et al. 1997). Recent studies have shown that the frequency of these outbursts decreases as the orbital period increases (Levitan et al. 2011; Ramsay et al. 2012).

The population density of AM CVn systems has not been conclusively determined: population synthesis estimates of the space density (Nelemans et al. 2001) have not been observationally confirmed by color-selected samples (Roelofs et al. 2007b; Carter et al. 2013). The latest results of Carter et al. (2013) suggest a space density of (5 ± 3) × 10⁻⁷ pc⁻³, a factor of 50 lower than the population synthesis estimates by Nelemans et al. (2001). The reason for this discrepancy is currently unknown. We note that these color-selected samples are mostly sensitive to longer period systems.

Over the last two years, we have conducted a search for outbursting AM CVn systems using the Palomar Transient Factory⁸ (PTF; Law et al. 2009; Rau et al. 2009) large-area synoptic survey, in part, to use a different approach to the discovery of AM CVn systems that does not rely on their colors. The PTF uses the Palomar 48'' Samuel Oschin Schmidt telescope to image up to ${\sim }2000\deg ^{2}$ of the sky per night to a depth of R ∼ 20.6 or g' ∼ 21.3. After identifying outbursting systems, we obtain classification spectra. To date, this survey has identified $\mathord {>}340$ cataclysmic variables (CVs), six new AM CVn systems, and one extremely faint AM CVn candidate (Levitan et al. 2011, 2013).

In this paper, we present PTF1 J191905.19+481506.2—a new AM CVn system discovered by the PTF. This system is particularly interesting for the following reasons.

• 1.  
  Shallow eclipses are present, making it only the second known eclipsing AM CVn system.
• 2.  
  Its orbital period is the shortest known of outbursting AM CVn systems.

We note that while this system is in the Kepler field, it, unfortunately, fell into a gap between the detectors of the Kepler satellite.

This paper is organized as follows. In Section 2 we discuss our data reduction and analysis methods. In Section 3 we present both short-term and long-term photometric and spectroscopic data and perform a period analysis. We consider the system's geometric structure and the rate of eclipsing AM CVn systems in Section 4. We summarize in Section 5.

2.1. Photometric Data

2.2. Spectroscopic Data

2.3. Period Estimation

PTF1J1919+4815 was detected as a possible transient on 2011 July 11. A spectrum obtained at the Palomar 200'' (P200) on 2011 July 23, likely while the system was still in outburst, showed a strong He ii 4686 emission line, but no other significant spectral lines (Figure 1). A second spectrum taken on 2011 August 03 detected the system while it was in a much fainter state, resulting in much lower signal-to-noise with few discernible lines. The combination of the crowded field (relative to the PTF's pixel scale of 101 pix⁻¹) and lack of a clear spectral signature led to the object being left for future study.

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In 2011 December, PTF1J1919+4815 was identified as a candidate outbursting source as part of the PTF search for outbursting sources (Levitan et al. 2013). The combination of no Balmer lines and the presence of He ii in the initial spectrum, as well as the outbursting behavior, led to its initial classification as an AM CVn system candidate. The likelihood that PTF1J1919+4815 was a short-period system made phase-resolved observations necessary to understand its nature.

Between 2012 May and 2013 April, we obtained both high-speed photometry (Section 3.1) and phase-resolved spectroscopy (Section 3.2). Simultaneously, we began long-term photometric monitoring of PTF1J1919+4815 (Section 3.3). A summary of all non-PTF observations is presented in Table 1. A finding chart, useful given the crowded nature of the field, is shown in Figure 2. A long-term light curve of PTF1J1919+4815, indicating the times of the higher-cadence observations described below, is presented in Figure 3. Most data presented here are publicly available on the PTF Web site.

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3.1. High-cadence Photometric Observations

The high-cadence photometric observations provide the most unambiguous measurements of PTF1J1919+4815's geometric configuration. In 2012 May, June, and July we obtained several series of exposures with the Lick 3 m Shane telescope using a g' filter and 15–30 s exposures (see Table 1). The three-month baseline of these observations provides the best estimate for a period. However, the relatively long exposure times result in a poor time resolution relative to the short orbital period of AM CVn systems. Our more recent data set, from the P200 telescope with the CHIMERA instrument, has similar signal-to-noise per exposure but with only 5 s integrations and effectively no dead time between exposures. This light curve is ideal for studying the intra-orbital photometric variability of PTF1J1919+4815, but its short length of only $\mathord {\sim }1$ hr precludes its use for a period determination. We use the CHIMERA data set to show the presence of an eclipse (Section 3.1.1) and the Lick data set to identify the precise orbital period (Section 3.1.2).

3.1.1. CHIMERA Light Curve

We present the CHIMERA light curve in Figure 4. At the time of observation, PTF1J1919+4815 was in outburst, and shows the characteristic superhumps seen in other outbursting AM CVn systems (e.g., Figure 4 of Wood et al. 2002). Superhumps are believed to be caused by deformation of the disk while the system is in outburst, and typically have periods a few percent longer than the orbital period (Warner 1995). The most prominent features noticeable besides the sawtooth shape of the superhumps are the three "dips" in luminosity for each orbit of PTF1J1919+4815. We now consider the nature of these dips.

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We begin by considering the dip which occurs shortly before peak luminosity for each superhump cycle. If one were to remove the other two dips and draw a straight line from peak luminosity to minimum luminosity, then this last dip constitutes the minimum of the superhump sawtooth shape. Hence, we do not believe it to have a geometrical cause beyond the superhump phenomenon itself.

We now turn our attention to the remaining two dips. The presence of one dip with a $\mathord {\le }5\%$ decrease (labeled as "superhump dip" in Figure 4) has been observed in other AM CVn systems (e.g., CR Boo and V803 Cen both show such features; Patterson et al. 1997, 2000), but its cause is unknown. However, a second decrease of any kind has not been observed in other systems except SDSSJ0926+3624 (C11), the first-discovered eclipsing AM CVn system. Hence, we tentatively conclude that the dip near the peak of the superhump is a feature intrinsic to the superhump itself, while the second dip is an eclipse of some localized emission on the disk, possibly the hot spot. This conclusion is supported by additional data we acquired from the Shane telescope, which we discuss in Section 3.1.2.

The CHIMERA light curve indicates only a single eclipse, with a relatively symmetric ingress and egress. We did consider the possibility that the third dip, which we earlier identified as part of the sawtooth pattern, is an eclipse of the hot spot, making the earlier dip an eclipse of the white dwarf. However, the separation of $\mathord {\sim }4$ minutes between these two dips means that the donor would have to eclipse the hot spot at an orbital phase offset of $\mathord {\sim }0.2$, something that is highly unlikely.

In addition to the prominent dips, we also highlight the presence of significant variability on the one- to two-minute timescale, particularly immediately after the superhump dip. The source of this variability is unknown and further observations are necessary to ensure that it is, in fact, real. If so, it may be related to the flickering phenomenon seen in CVs (Bruch 1992).

3.1.2. Lick Light Curves

The data we obtained from the Lick Shane telescope, while of coarser time resolution, are particularly useful for period analysis due to its long baseline. We present light curves, Lomb–Scargle periodograms, and folded light curves in Figures 5 and 6 for the nights of 2012 May 16 and 2012 July 20, respectively. For the data from each night, we subtracted a first order linear fit to remove any linear trend. The remaining nights show similar characteristics to these, and are both included in the overall period analysis in Section 3.1.3. All light curves are available on the PTF Web site.

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Both the 2012 May and 2012 July data show variability with an amplitude of $\mathord {\sim }0.15$ mag, including both a small increase in luminosity as well as the eclipse feature. The eclipse is of similar depth to that seen in the CHIMERA data. We believe that the luminosity increase in the 2012 July data is a result of the hotspot rotating into view, as has been shown for CVs (e.g., OY Car; Schoembs & Hartmann 1983) and AM CVn systems. Levitan et al. (2011) explicitly showed the link between the hot spot and the variability for PTF1J0719+4858, but the variability in quiescence has been observed in several AM CVn systems (e.g., Patterson et al. 1997; Wood et al. 2002).

The origin of the 2012 May variability is more difficult to determine. PTF1J1919+4815 was in outburst at the time the data were obtained (see Figure 3) and this would typically indicate the presence of superhumps (as seen in the CHIMERA data; see Section 3.1.1). Superhumps, however, typically have a period slightly longer than the orbital period (Patterson et al. 2005) and hence should be out of phase with the eclipse. We note that WZ Sge, a CV which also shows only a disk eclipse, has shown significant differences in eclipse shape between outburst and quiescence (Patterson et al. 2002). Given the faintness of PTF1J1919+4815 in quiescence, we do not have sufficient signal-to-noise to clearly see any such differences in this system.

In this case, the eclipse in the 2012 May data is at the same phase with respect to the luminosity increase as in the 2012 July data. With fewer than two hours of data, it is impossible to measure the period of the variability with sufficient precision to distinguish between the orbital period and any slightly longer superhump period. Hence, we cannot say whether this variability is from the hot spot (thus indicating the lack of superhumps for part of the outburst; see, e.g., C11) or a chance superposition of the orbital period and the superhump period (as can be seen, for example, in Figure 3 of C11 for SDSSJ0926+3624).

We now consider the origin of the observed eclipse. The accretor is believed to be significantly more compact than the donor and contributes the majority of the luminosity of at least some AM CVn systems (e.g., Bildsten et al. 2006; C11). Any eclipse of the accretor would likely have a sharp ingress and egress and a depth of greater than 10% (as observed for PTF1J1919+4815). This leaves only a localized area of emission on the disk as a possible explanation. It is likely that this is the hot spot typically observed for semi-detached binaries, particularly for the data taken in quiescence. However, we are cautious about concluding that all data, including that taken in outburst, shows an eclipse of the hot spot as opposed to some other disk feature, as it is believed that the superhump phenomena results in a somewhat more extended area of disk emission (Simpson & Wood 1998).

3.1.3. Orbital Period Analysis

We stress that while the eclipse feature is certainly not strong, it is stable and repeating over data sets obtained with two different telescopes and over a time span of over 10 months, as well as in both quiescent and outburst states of PTF1J1919+4815. We thus conclude that this is an eclipse of the hot spot and proceed with a period analysis.

Outbursting AM CVn systems have been observed to show photometric variability at two periods, one in outburst (from the superhumps) and one in quiescence (likely from the hot spot; Levitan et al. 2011). We begin our analysis with the 2012 July data, which, being in quiescence, should show variability at the orbital period of the system. The best periods obtained from the individual light curves were 22.41 ± 0.09 minutes and 22.44 ± 0.11 minutes for the nights of 2012 July 18 and 2012 July 20, respectively. If we analyze all the 2012 July data simultaneously, we find a best period of 22.46 ± 0.09 minutes. Finally, if we analyze all data from 2012 May, June, and July simultaneously, we find a best period of 22.4559(3) minutes. All errors are calculated from bootstrap simulations, as described in Section 2.3. The times used in all calculations are based on the local time stamps provided by the Kast instrument, offset to UTC time, and subsequently corrected for light-travel time by conversion to HJD. We note that the error estimates of the period do not change linearly with respect to the baseline, as might be expected. This may be due to slight changes in the light curve between nights, perhaps similar to that seen for WZ Sge (Patterson et al. 2002). Lastly, the closest alias is at 22.4671 minutes, which is 37σ from the best period and shows no eclipse in a phase-binned light curve.

The long exposure times used, the weakness of the eclipse, and the faint nature of the system make ephemeris determination non-trivial. To obtain an estimate with an accurate error estimate, we used a matched filter technique. For our template, we selected a triangular function of width 3.7 minutes and depth 0.1 mag (roughly consistent with the shape of the eclipse). We calculated the cross-correlation of the template with the light curve using a Δt = 1 s and identified the maxima of this function for each orbital period. We limited the data used to that from 2012 May 16, the first half of 2013 July 18, and 2013 July 20. If the remainder of the data is included, the resulting ephemeris does not change significantly, but the error, calculated based on the scatter of the observed minima with the predicted minima, is significantly higher due to poorer conditions in that data.

Given the agreement in the periods measured from these light curves, we propose 22.4559(3) minutes as the orbital period of PTF1J1919+4815 and an ephemeris of

$$\begin{equation*} {\rm HJD}=2456063.8650(5)+0.0155944(2)E \end{equation*}$$

for the mid-point of the eclipse. We acknowledge that the error measurement of the orbital period may be overly optimistic, as it is from a bootstrap simulation and does not include systematic errors. However, it is consistent with the long baseline of observations and when all data are folded at this period, the eclipse is clearly visible (Figure 7).

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3.2. Phase-resolved Spectroscopy

We obtained phase-resolved spectroscopy on six separate nights, although the outburst state was different for almost every one of these nights (see Figure 3 and Table 1). For each night, we co-added the flux around several He emission lines in velocity space and folded the spectra on the orbital period identified in Section 3.1.3. However, no clear "S-wave" was observed, as would be expected for an AM CVn system (Nather et al. 1981). Likewise, a search for an orbital period in this data did not yield any discernible orbital period. We present one of these trailed, phase-folded spectra in Figure 8.

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Next, we generated Doppler tomograms (Marsh & Horne 1988) using multiple emission lines (which provides for better signal to noise in faint, shorter period systems—see Kupfer et al. 2013 for an example). The tomogram should boost the signal of the hot spot since all light is concentrated at one point in velocity space, yet we did not see any strong hot spot in any of the Doppler tomograms (Figure 8).

The lack of a clear signal in the trailed, phase-folded spectra and the Doppler tomograms is mysterious. We do not attribute this to a lack of data, as Levitan et al. (2013) found a faint S-wave in a g' > 22 AM CVn system with a similar amount of data as what we obtained for PTF1J1919+4815. Time variability of the visibility of the hot spot has been observed for other AM CVn systems (Roelofs et al. 2009) and may be the reason we could not identify an S-wave in the data. We thus conclude that the hot spot in PTF1J1919+4815 had lower contrast with the accretion disk than in other systems during the times it was observed, but note that this does not change our conclusions about the orbital period, as no clear signal was found in the spectroscopic data and the photometric signal is much less ambiguous.

3.3. Long-term Photometric Variability

We explore the long-term photometric variability of PTF1J1919+4815 using the 355 photometric measurements over $\mathord {\sim }200$ days. We identify the presence of super-outbursts and the much shorter normal outbursts. The latter, in particular, appear to last $\mathord {\sim }1$ day and were observed to occur 2–3 times between super-outbursts, consistent with CR Boo (Kato et al. 2000) and PTF1J0719+4858 (Levitan et al. 2011), both of which have similar orbital periods as PTF1J1919+4815.

Our data covers approximately 5 super-outburst cycles and these super outbursts are seen to recur every $\mathord {\sim }35$ days. We use a periodogram of the P60 data to calculate the strongest period of 36.8 ± 0.4 days. Folding the light curve (see Figure 9) at this period results in a remarkably well-defined super-outburst light curve, something that is not seen in most AM CVn systems. We note that we use the periodogram itself only to refine the already evident recurrence time.

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This super-outburst recurrence time is slightly shorter than that observed for CR Boo, which has an orbital period of 24.5 minutes. The duration of the super-outburst is $\mathord {\sim }13$ days and its amplitude (peak luminosity minus quiescent luminosity) is $\mathord {\sim }3$ mag. Although the data shown here indicates regular super-outbursts, long-term study of CR Boo (Honeycutt et al. 2013) suggests that this is likely to only be a temporary state. Moreover, with only five observed super-outbursts, the scatter in the super-outburst recurrence time measured here is likely to be underestimated as most AM CVn systems show changes of up to 15% (see, e.g., Ramsay et al. 2012).

3.4. Median Spectra and Long-term Variability

AM CVn systems are known to have significant changes in their spectra corresponding to their different states. PTF1J1919+4815, with its frequent changes between outburst and quiescence, provides a good opportunity to consider the changes in spectra. In Figure 10 we present median spectra from three nights. In particular, these spectra show the evolution of the spectrum from the emission-line spectrum that AM CVn systems exhibit in quiescence (the 2012 June 13 spectrum) to the outburst spectrum showing absorption lines (the 2012 June 20 spectrum). The figure also shows the spectrum taken on 2012 August 28, at the start of a super-outburst (see Figure 3), which shows a distinct lack of features. We present equivalent widths of the significant lines identified in the outburst and quiescent spectra in Table 2.

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Table 2. Equivalent Widths of Prominent Lines

Line	Quiescence	Outburst
	(Å)	(Å)
He i 3705	⋅⋅⋅ a	4.4 ± 0.2
He i 3819	−2.7 ± 0.5	3.9 ± 0.2b
He i 3888	−4.9 ± 0.4
He i 3964	−1.1 ± 0.4	6.8 ± 0.2c
He i 4026	⋅⋅⋅  a
He i 4120/4143	X	4.4 ± 0.2
He i 4388	−0.5 ± 0.4	1.5 ± 0.1
He i 4471	−2.2 ± 0.4	3.3 ± 0.2
He ii 4685	−3.4 ± 0.4d	−0.5 ± 0.1
He i 4713		⋅⋅⋅  a
He i 4921	⋅⋅⋅  a	2.8 ± 0.2
He i 5015	−3.6 ± 0.4	1.8 ± 0.2
He i 5875	−2.7 ± 0.6	1.9 ± 0.3
He i 6678	2.6 ± 0.5	⋅⋅⋅  e
He i 7065	−2.1 ± 0.5	X

Notes. Quiescent equivalent widths are measured from the medium spectrum taken on 2012 August 16. Outburst equivalent widths are measured from the median spectrum taken on 2012 August 17. An X indicates that the line is not detectable above the noise level of the spectrum obtained. ^aLine present but insufficient S/N to measure. ^bCombined equivalent width of He i 3819 and He i 3888. ^cCombined equivalent width of He i 3964 and He i 4026. ^dCombined equivalent width of He ii 4685 and He i 4713. ^eLine present but contaminated with atmosphere.

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Although the presence of the disk eclipse allows for a precise period estimate for such a faint system, the lack of an accretor eclipse precludes a full solution of the system structure. It also prevents us from measuring any change in the orbital period for this system, as the uncertainty on the disk radius and its likely change due to the presence of outbursts will overwhelm any expected change due to orbital evolution. However, the presence of the disk eclipse does allow us to estimate the inclination of the system if we make certain assumptions about the masses and radii of the two components.

Only one AM CVn system, SDSSJ0926+3624 (P_orb = 28.558 minutes; C11), has been observed to have eclipses of the accretor: and as a result, has had a full system solution derived. Here, we assume that the evolutionary origin and, hence, system properties of PTF1J1919+4815 are similar to that of SDSSJ0926+3624. We caution that AM CVn systems are believed to have several possible evolutionary pathways (Nelemans et al. 2001) and that no conclusive evidence has been found that can differentiate a system's origin (Roelofs et al. 2007a; Nelemans et al. 2010).

From C11, the parameters of SDSSJ0926+3624 are: M₁ = 0.85 M_☉, M₂ = 0.035 M_☉, R₁ = 9.7 × 10⁻³ R_☉, and q = M₂/M₁ = 0.041. We use these values for PTF1J1919+4815, despite the difference in orbital period. Assuming standard mass transfer, only $\mathord {\sim }0.01\,M_\odot$ would have been transferred during the evolution between the orbital periods of PTF1J1919+4815 and SDSSJ0926+3624, which is certainly below the uncertainties of the systems' respective evolutionary histories. Similarly, we adopt the radius of the accretor from SDSSJ0926+3624 to be that of the accretor in PTF1J1919+4815. The donor's radius will be that of the Roche Lobe as a first estimate, at the given period for the given mass ratio. Using the assumed masses for PTF1J1919+4815 and its orbital period with the approximation in Eggleton (1983), we find that the orbital separation, a, and the Roche-lobe radius, r_L, are a = 0.25 R_☉, and r_L = 0.039 R_☉. The last measurement needed is the radius of the disk. Based on the measurements of SDSSJ0926+2624, we assume a disk radius of R_disk ≈ 0.35a. The derived values in C11 vary between observing runs, and since the size of the disk is expected to change between outburst and quiescence, we believe this is a reasonable, first assumption. We can constrain the inclination by determining for which inclinations the shadow of the donor star (as seen from Earth) passes over the hot spot, assumed to be a point source, but not the accretor. Combining the radii of disk, donor, and accretor, we can constrain the inclination to 76° < i < 79°. We present a summary of system properties in Table 3.

Table 3. System Properties

System Property	Value
Ua	19.814 ± 0.048
ga	20.155 ± 0.016
ra	20.522 ± 0.065
ia	20.409 ± 0.121
Orbital period	22.4559(3) minutes
Super-outburst recurrence time	36.8 ± 0.4 days
Super-outburst strength	$\mathord {\sim }3$ mag
Super-outburst duration	$\mathord {\sim }13$ days
Inclinationb	76° < i < 79°

Notes. ^aPhotometric measurements are from the Kepler INT Survey DR2 (Greiss et al. 2012a, 2012b) and are typical of the quiescent values for the system. ^bThese values assume that the radius of the accretion disk, R_disk = 0.35a, where a = 0.25 R_☉. See Section 4 for further details.

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Given the uncertainty in the system properties used, we consider the impact of several parameters on the results. The upper limit on i is primarily related to the accretor and donor masses. However, their masses are constrained by the initial conditions of AM CVn system formation and as such are not likely to change significantly even with moderate changes of the initial masses. The lower bound, however, is significantly affected by the radius of the disk. If, for example, the radius of the disk was 0.5a, the lower bound would drop to 72°. Had the inclination been greater than the upper limit here, an eclipse of the white dwarf should have been observable. Conversely, an inclination of less than the lower limit would result in no eclipse of either the hot spot or the accretor.

4.1. Number of Eclipsing AM CVn Systems

We have established PTF1J1919+4815 as an eclipsing AM CVn system from its He-rich and H-deficient spectrum, its phenomenological behavior, and its photometrically determined period. Its orbital period was measured to be 22.4559(3) minutes, but we did not find a significant hot spot in the S-wave and Doppler map. PTF1J1919+4815 shows a well-defined super-outburst that is 36.8 ± 0.4 days. Such a non-variable recurrence time is atypical for AM CVn systems. We used these measurements to constrain the inclination of the systems given assumptions about its evolution and to consider the geometric structure of the system.

T.K. acknowledges support by the Netherlands Research School of Astronomy (NOVA). We thank Dong Xu for reducing the initial classification spectra.

Observations obtained with the Samuel Oschin Telescope at the Palomar Observatory as part of the Palomar Transient Factory project, a scientific collaboration between the California Institute of Technology, Columbia University, Las Cumbres Observatory, the Lawrence Berkeley National Laboratory, the National Energy Research Scientific Computing Center, the University of Oxford, and the Weizmann Institute of Science. Some of the data presented herein were obtained at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W. M. Keck Foundation. The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. Based in part on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), Ministério da Ciência, Tecnologia e Inovação (Brazil) and Ministerio de Ciencia, Tecnología e Innovación Productiva (Argentina). The Gemini data were obtained under Program ID GN-2012B-Q-110. Based in part on observations made with the Gran Telescopio Canarias (GTC) installed in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofísica de Canarias, on the island of La Palma.

Facilities: PO:1.2m - Palomar Observatory's 1.2 meter Samuel Oschin Telescope, PO:1.5m - Palomar Observatory's 1.5 meter Telescope, Gemini:Gillett (GMOS-N) - Gillett Gemini North Telescope, GTC (OSIRIS) - Gran Telescopio CANARIAS 10.4m telescope (GRANTECAN), Hale (CHIMERA, DBSP) - Palomar Observatory's 5.1m Hale Telescope, Keck:I (LRIS) - KECK I Telescope, Shane (Kast Double spectrograph) - Lick Observatory's 3m Shane Telescope