the 2011 eruption of the recurrent nova t pyxidis: the

The Astrophysical Journal, 773:55 (23pp), 2013 August 10 doi:10.1088/0004-637X/773/1/55C© 2013. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

THE 2011 ERUPTION OF THE RECURRENT NOVA T PYXIDIS: THE DISCOVERY, THE PRE-ERUPTION RISE,THE PRE-ERUPTION ORBITAL PERIOD, AND THE REASON FOR THE LONG DELAY

Bradley E. Schaefer1, Arlo U. Landolt1, Michael Linnolt2, Rod Stubbings3, Grzegorz Pojmanski4, Alan Plummer5,Stephen Kerr6, Peter Nelson7, Rolf Carstens8, Margaret Streamer9, Tom Richards10,

Gordon Myers11, and William G. Dillon121 Physics and Astronomy, Louisiana State University, Baton Rouge, LA 70803, USA

2 American Association of Variable Star Observers, 49 Bay State Road, Cambridge, MA 02138, USA3 Tetoora Observatory, Tetoora Road, Victoria, Australia

4 Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland5 Variable Stars South, Linden Observatory, 105 Glossop Road, Linden, NSW, Australia

6 American Association of Variable Star Observers, Variable Stars South, Astronomical Association of Queensland, 22 Green Avenue, Glenlee, Queensland, Australia7 Ellinbank Observatory, 1105 Hazeldean Road, Ellinbank 3821, Victoria, Australia

8 American Association of Variable Star Observers, Variable Stars South, CBA, Geyserland Observatory, 120 Homedale Street, Rotorua 3015, New Zealand9 Lexy’s Palace Observatory, 3 Lupin Place, Murrumbateman, NSW, Australia

10 Variable Stars South, Pretty Hill Observatory, P.O. Box 323, Kangaroo Ground 3097, Victoria, Australia11 Center for Backyard Astrophysics, Columbia University, 538 West 120th Street, New York, NY 10027, USA12 American Association of Variable Star Observers, 4703 Birkenhead Circle, Missouri City, TX 77459, USA

Received 2011 August 30; accepted 2013 May 29; published 2013 July 25

ABSTRACT

We report the discovery by M. Linnolt on JD 2,455,665.7931 (UT 2011 April 14.29) of the sixth eruption of therecurrent nova T Pyxidis. This discovery was made just as the initial fast rise was starting, so with fast notificationand response by observers worldwide, the entire initial rise was covered (the first for any nova), and with high timeresolution in three filters. The speed of the rise peaked at 9 mag day−1, while the light curve is well fit over onlythe first two days by a model with a uniformly expanding sphere. We also report the discovery by R. Stubbings ofa pre-eruption rise starting 18 days before the eruption, peaking 1.1 mag brighter than its long-time average, andthen fading back toward quiescence 4 days before the eruption. This unique and mysterious behavior is only thefourth known (with V1500 Cyg, V533 Her, and T CrB) anticipatory rise closely spaced before a nova eruption. Wepresent 19 timings of photometric minima from 1986 to 2011 February, where the orbital period is fast increasingwith P/P = +313,000 yr. From 2008 to 2011, T Pyx had a small change in this rate of increase, so that the orbitalperiod at the time of eruption was 0.07622950 ± 0.00000008 days. This strong and steady increase of the orbitalperiod can only come from mass transfer, for which we calculate a rate of (1.7–3.5) × 10−7 M� yr−1. We report6116 magnitudes between 1890 and 2011, for an average B = 15.59 ± 0.01 from 1967 to 2011, which allows foran eruption in 2011 if the blue flux is nearly proportional to the accretion rate. The ultraviolet–optical–infraredspectral energy distribution is well fit by a power law with fν ∝ ν1.0, although the narrow ultraviolet region has atilt with a fit of fν ∝ ν1/3. We prove that most of the T Pyx light is not coming from a disk, or any superpositionof blackbodies, but rather is coming from some nonthermal source. We confirm the extinction measure from IUEwith E(B − V ) = 0.25 ± 0.02 mag.

Key words: novae, cataclysmic variables – stars: individual (T Pyx)

Online-only material: machine-readable table

1. INTRODUCTION

T Pyxidis (T Pyx) was the first recognized recurrent nova(RN) with three known eruptions, and this over a dozenyears before any other RN was recognized with their seconddiscovered eruption (Schaefer 2010). T Pyx is one of thequintessential RNe, with six eruptions in 1890, 1902, 1920,1944, 1967, and now 2011. Over the past few decades, thequiescent magnitude of T Pyx has been near V = 15.5, whilepeaking at V = 6.4 (Schaefer 2010). T Pyx is slow among RNe,with its overall rise taking 40 days and the decline by 3 mag fromthe peak (t3) taking 62 days (Schaefer 2010). T Pyx is one of thetwo RNe with an orbital period (P = 0.076 days; Schaefer et al.1992; Patterson et al. 1998; Uthas et al. 2010) that is inside orbelow the period gap. T Pyx is the only RN that has a nova shell,and this shell has a surprisingly slow expansion velocity plusa structure of thousands of discrete knots (Duerbeck & Seitter1979; Williams 1982; Shara et al. 1989, 1997). The homologous

expansion of the individual shell fragments proves that theysuffer no significant deceleration, and the expansion velocities,viewed over 13 yr by Hubble Space Telescope, prove that thenova shell was ejected in the year 1866 ± 5 at a velocity of500–715 km s−1 (Schaefer et al. 2010a). This velocity is greatlytoo low to be from an RN event, and with the measured massin the shell (∼10−4.5 M�) being much greater than is possiblefor an RN event, we know that the 1866 event must have beena normal classical nova eruption with a very long time of priorquiescence (Schaefer et al. 2010a).

A key and unique fact about T Pyx is that its quiescent bright-ness has been secularly fading from 1890 to 2011 (Schaefer2005). Before the 1890 eruption, archival plates show T Pyxto be constant at B = 13.8, the next four inter-eruption inter-vals have B-band magnitudes of 14.38, 14.74, 14.88, and 14.72,while after the 1967 eruption the quiescent B magnitude fadesfrom 15.3 (around 1968) to 15.5 and fainter (from the 1980s un-til 2011). This secular decline is fundamental for understanding

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http://dx.doi.org/10.1088/0004-637X/773/1/55

The Astrophysical Journal, 773:55 (23pp), 2013 August 10 Schaefer et al.

the evolution and fate of T Pyx, as well as the timing of theeruptions.

T Pyx has long been a perplexing system because there wasno known way to combine a short orbital period and a highaccretion rate into one system. (The accretion rate must behigh, ∼10−7 M� yr−1, so as to provide enough mass to triggereruptions with a typical recurrence time interval of two decades.)That is, a binary with a period below the gap should have angularmomentum loss only from gravitational radiation, and this willlead to an accretion rate many orders of magnitude smallerthan required (Patterson 1984). This dilemma was solved byKnigge et al. (2000), who realized that the high accretion ratecould be powered by a hot-luminous source on the white dwarfthat is irradiating the companion star. For T Pyx, this idea wasextended by Schaefer et al. (2010a), where the classic novaeruption of 1866 started a high accretion rate which drives alargely self-sustained supersoft source (with nuclear burning onthe surface of the white dwarf). This mechanism can only workfor binaries with short orbital periods and where a substantialmagnetic field channels the material into a small area. Thesecular fading demonstrates that the supersoft source is notfully self-sustaining. Schaefer & Collazzi (2010) generalize thissituation to a whole class of stars, called V1500 Cyg stars, whichsuffer this same secular fading after their eruptions, and whichalso have short orbital periods and highly magnetized whitedwarf stars. So the picture we have for T Pyx is that it startedout as an ordinary cataclysmic variable with a long durationin quiescence, then an ordinary nova eruption in 1866, whichignited a supersoft X-ray source whose irradiation drives a highaccretion rate and hence RN eruptions. Over time, the accretionrate declines and the quiescent brightness fades.

By far the most important question relating to RNe is whetherthey are the progenitors of Type Ia supernovae. RNe are one ofthe best progenitor candidate classes because they certainly havewhite dwarfs near the Chandrasekhar mass being fed with a veryhigh accretion rate. Nevertheless, T Pyx (and any short-periodRNe) cannot be a Type Ia progenitor (Schaefer et al. 2010a;Schaefer 2010). The reason is that the ordinary nova eruption in1866 easily dominates the dynamics of the system, and ordinarynova eruptions eject much more mass than has been accretedby the white dwarf, so the white dwarf is actually losing massand will not become a supernova. Another way of saying thisis that the secular fading shows that the interval during whichT Pyx is an RN is short, perhaps under two centuries, and thuswill have little effect even if the RN event causes a net gain ofmass on the white dwarf. The long-period RNe remain as strongprogenitor candidates (due to their high accretion rate beingdriven by the evolution of the donor star so as to be sustainedand long in duration), but the short-period RNe cannot sustainthe high accretion and so are not supernova progenitors.

The secular fading is also critical for understanding the timingand triggering of the RN events. Schaefer (2005) demonstratedthat the inter-eruption intervals during which the quiescentmagnitude was faint on average lead to a long time betweeneruptions, while a bright interval leads to a short time betweeneruptions. The physics is simple, with the amount of accretedmass required to trigger an eruption being constant for anyone RN, with the accretion rate being given by the blue fluxdominated by the accretion disk, so that a high quiescentbrightness implies a high accretion rate and a short intervalneeded to accrete the trigger mass. With calibration fromprior inter-eruption intervals, it becomes possible to predict theapproximate date of upcoming eruptions. As such, Schaefer

(2005) predicted that U Sco would erupt in 2009.3 ± 1.0,and this advance notice allowed for an incredible coverage ofthe eruption in 2010.1, so that this now is the all-time best-observed nova event (Schaefer et al. 2010b). Schaefer (2005)also predicted that the next eruption of T Pyx would be in 2052 ±3, while Schaefer et al. (2010a) extended this to predict that thenext eruption will be after the year 2225. This prediction wasshown to be horribly wrong by the T Pyx eruption in 2011, sothe issue now is to understand where the prediction went wrong.

2. DISCOVERY OF THE 2011 ERUPTION

Starting with the 1890 eruption, the inter-eruption intervalswere 11.9, 17.9, 24.6, and 22.1 yr (Schaefer 2010). In themiddle-1980s many people around the world (including manyof the authors of this paper) realized that T Pyx “should” go offsoon. The simple logic was that the prior inter-eruption intervalwas 22 yr, and 22 yr after the 1967 eruption placed the expectederuption around 1989. Due to this, many of the world’s bestobservers kept T Pyx under nightly watch.

As the 1990s stretched on, the expected eruption did not occur.Vigilance was maintained and occasional public exhortationswere published in many journals. Any eruption after 1967 wouldcertainly have been discovered because the T Pyx eruptionduration is two-thirds of a year and so no eruption would havebeen missed due to observational gaps. By the year 2005, theinterval since the prior eruption extended to 38 yr, with noexplanation for this long delay.

Schaefer (2005) provided the explanation. The key pointwas that T Pyx was declining overall in brightness ever sincebefore the 1890 eruption, and this means that the accretion ratein the system was declining, so that it would take longer toaccumulate the trigger mass on the surface of the white dwarf.In particular, from the 1944 to the 1967 eruption the averageB-band magnitude was 14.72 ± 0.03, while from 1967 to 2005the average B-band magnitude was 15.49 ± 0.01 mag. Theblue flux from T Pyx is entirely dominated by the accretionlight, so the B-band flux is determined by the accretion rate.The dimming by 0.77 mag (a factor of two drop in accretionluminosity) requires a fall in the accretion rate by at least a factorof two, and this forces an increase in the inter-eruption timeinterval. The long delay since 1967 was simply because the rateof accumulation on the white dwarf surface had at least halved.

With the prediction from Schaefer (2005) that the next erup-tion was a long way in the future, professional preparations forthe next eruption (like the long-running Target of Opportunityprogram with the Hubble Space Telescope) were stopped. Butthis did not deter many amateur observers (who were generallyfully aware of the prediction) from continuing to regularly checkfor eruptions, report visual magnitudes, and make long photo-metric time series with CCDs. For example, Stubbings startedin early 1998 in making 2004 visual estimates of the bright-ness of T Pyx with a 16 inch reflecting telescope in Victoria,Australia, while Richards made CCD time-series observationson five nights in early 2008 with 661 brightness measures.

From 2008 until 2011 March 31 (JD 2,455,652), T Pyx heldat a steady magnitude of V = 15.5, with the usual periodicvariations of �0.2 mag. For the 610 observations in the last4000 days before 2011 March, Stubbings never saw T Pyxbrighter than V = 15.0. Thus, it was a distinct surprise whenStubbings saw T Pyx at V = 14.5 on JD 2,455,657.0104(UT 2011 April 5.51), with the realization that this could bean initial rise of the next T Pyx eruption. With this exciting

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Table 1Pre-eruption Rises

Nova Nova Classa,b Vpeaka VQ

a ΔVrisec T − Tstart

c,d Rise Morphologyc

(days)

T Pyx P(62), RN, short-Porb, shell 6.4 15.5 1.1 −18 to −4 “Bump”V533 Her S(43), CN, IP, shell 3.0 15.0 1.7 −550 to 0 Exponential riseV1500 Cyg S(4), neon CN, polar, shell 1.9 20.5c 7.0 < − 23 to 0 Fast and bright riseT CrB S(6), RN, red giant donor 2.5 10.5 1.4 −3200 to −220 “Bump”

Notes.a Data from Strope et al. (2010).b Class P has a plateau in the light curve. Light curve class S has a smooth monotonic decline. The parentheses contain the t3 decline rate.CN indicates a classical nova with no known recurrences. IP indicates an intermediate polar system with a highly magnetized white dwarf.c Data from Collazzi et al. (2009) and this paper.d T − Tstart is the time from the beginning and end of the pre-eruption rise to the start of the eruption.

possibility, Stubbings checked T Pyx five more times beforedawn, but saw only variations from 14.4 to 14.7. On the nextnight T Pyx was slightly fainter, and it continued a slow fade untilJD 2,455,662.0417 (UT 2011 April 10.54). The light curve justbefore the sharp initial rise to the eruption shows a significantshort-duration pre-eruption rise, which started roughly 18 daysbefore the eruption, peaked at around V = 14.4 mag (close to1.1 mag above the prior quiescent level), and declined towardquiescence before the eruption.

The existence of this pre-eruption rise has been con-firmed with the archival data from the All Sky AutomatedSurvey (ASAS; Pojmanski 2002). For the year prior toJD 2,455,654.5972, T Pyx was seen to vary about its aver-age with an rms scatter of 0.14 mag. On JD 2,455,656.62097(0.4 days before Stubbings’s discovery of the rise), T Pyx was0.43 mag brighter than the average, with this being the brightestmagnitude by far in the prior year. This is the confirmation ofthe pre-eruption rise. The apparent duration of the pre-eruptionrise as seen with the ASAS data alone is roughly one week, orperhaps a bit longer.

This pre-eruption event is one of the unusual phenomenathat is still completely mysterious. Collazzi et al. (2009) reportan extensive survey of archival data for 22 novae with well-sampled pre-eruption light curves, which showed that three ofthese (V1500 Cyg, T CrB, and V533 Her) have pre-eruptionrises with greatly different timescales and morphologies. Table 1summarizes the properties of these events, and both the risesand the nova systems are a rather wide range of properties.In all cases, the pre-eruption rise cannot be some early partof the usual thermonuclear runaway because the anticipationtimes and the durations (weeks and years) are greatly longerthan both the theoretical runaway timescale (hours) and theobserved rise times (hours). In all cases, the pre-eruption event ishighly significant and unprecedented for many decade intervalsbefore the eruptions and after the eruptions. For the example ofV1500 Cyg, the fast apparently exponential rise was detected athigh significance on five plates, getting as bright as B = 13.5 mag(i.e., 7 mag brighter than quiescence), with no other suchbrightenings from 1890 to 2013. The uniqueness and the timecoincidence immediately before the eruption provide proof ofa causal connection. It is difficult to understand this connectionbecause the pre-eruption brightness level is dominated by theaccretion which is governed by the mass falling off the donorstar, while the timing of the eruption is governed by conditionsat the bottom of the accretion layer deep under the surface ofthe white dwarf. How could the donor star near the L1 pointpossibly anticipate that conditions deep in the white dwarf will

soon heat to the trigger threshold? We regard the explanation ofthese pre-eruption phenomena as a premier challenge for novatheory. Presumably, the durations, start times, and luminositieswill provide sensitive indicators of conditions at the time of thenova trigger.

Linnolt had also been monitoring T Pyx since 2009 March,waiting for an eruption, as part of a larger program of mon-itoring RNe for eruptions (e.g., Schaefer et al. 2010b). OnJD 2,455,664.8590 (UT 2011 April 13.36), Linnolt observedT Pyx at V = 14.5 and realized that this was unusually bright,and hence a possible indication of an eruption. With anticipation,Linnolt checked T Pyx early the next night (JD 2,455,665.7931,UT 2011 April 14.29), saw that the star was bright (V =13.0), and immediately realized that an eruption had started.Within 15 minutes, the discovery was announced on the AAVSOdiscussion group, which is closely monitored by interestedamateurs and professionals worldwide. The discovery wasquickly confirmed by Australian visual observers Plummer(JD 2,455,665.8847 V = 12.2) and Kerr (JD 2,455,665.941V = 11.3). CCD observers Nelson, Carstens, and Streamerstarted taking time-series photometry in multiple bands byJD 2,455,665.9973.

Serendipitously, Schaefer actually took the first images ofT Pyx certainly in eruption, 7.6 hr before Linnolt’s discovery.Fortuitously, on that night, Schaefer had half a night of serviceobserving on the SMARTS 0.9 m telescope on Cerro Tololoin Chile, with the bulk of the time being used to get a longtime series covering an eclipse of the recurrent nova U Sco (seeSchaefer 2011). This was scheduled many months in advance,and the night was filled out with images of other RNe, includingtwo B-band images of T Pyx. In the excitement and activity of theT Pyx eruption, these two images were overlooked for severalmonths, until the U Sco data were analyzed in due course. Thetwo 100 s exposures (JD 2,455,665.4752 and 2,455,665.4986)show T Pyx at B = 12.51 ± 0.01 and B = 12.50 ± 0.01,respectively, which is over 3 mag brighter than in quiescence,so the fast rise was certainly well underway. Intriguingly, thismagnitude (B = 12.50) is brighter than the discovery magnitudeof Linnolt (V = 13.0) made 7.6 hr later. One possibility is thatT Pyx had a stutter in its rise, with some local maximum in itslight curve hours before discovery, so that theorists would haveto explain why there should be a significant drop in brightnessnear the start of the fast initial rise. Another possibility is that theinitial light from the rise was extremely blue so that the V-bandinitial rise could have been monotonic and steady. Interpolationbetween Linnolt’s pre-discovery and discovery magnitudes tothe time of the SMARTS image pair gives V = 13.51 mag,

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and hence a required B − V = −1.0 mag for the extremecolor hypothesis. The two possibilities (a non-monotonic rise orB − V = −1.0) can be used to test the theoretical model.

To the best of our knowledge (e.g., Payne-Gaposchkin 1964;Robinson 1975; Collazzi et al. 2009; Strope et al. 2010; Schaefer2010), the 2011 eruption of T Pyx has by far the best-observedinitial rise from quiescence. Only a handful of old novae havehad just a few points taken during the later half of their initialfast rise, whereas T Pyx has high time resolution time series withnear full-time coverage and in multiple bands. Indeed, T Pyx isthe only nova with any observations in the middle or the firstpart of the initial fast rise. The very fast alert to the world and thevery fast reaction by observers worldwide made this possible.

3. OUR NEW OBSERVATIONS

We have collected a wide variety of new observations pre-sented here for the first time.

3.1. Pre-eruption Photometry of Quiescence

Schaefer has returned to the Harvard College Observatoryand remeasured the many archival photographic plates, withthis providing the photometric history of T Pyx from 1890to 1953 and the basis for knowing that T Pyx has suffered asevere decline in quiescent brightness from 1890 to present. Thebrightness of T Pyx was made by direct comparison with nearbycomparison stars with magnitudes in the B magnitude system(Schaefer 2010). The spectral sensitivity of the Harvard plates isidentical to that of the B magnitude system (indeed, these plateswere involved in the original definition of the system), andexplicit measures of the color terms (corrections proportional toB − V) always show that these are closely equal to zero. Withthis, the returned magnitudes are exactly in the B-magnitudesystem. Repeated experiments with the Harvard plates showthat an average plate (with a good nearby comparison sequence)will produce a magnitude with a 1σ uncertainty of 0.15 mag.Of particular interest are the first two magnitudes from Harvard,which are before the 1890 eruption. The 1890 eruption wasrecorded on a Harvard plate at B = 8.05 at a time 25 days afterthe two earliest plates, so we have to wonder whether we have arecord of some pre-eruption event or the initial rise (rather thana record of T Pyx in quiescence). We have two plates separatedby 2 days both showing the same magnitude (B = 13.80), so weknow that these plates could not show the initial rise. The 2011eruption light curve has B = 8.05 around 1 day and 70 days afterthe start of the initial fast rise, so the two plates at B = 13.80taken 25 days earlier cannot be part of either the initial fast riseor a pre-eruption bump. Therefore, these two plates show thepre-eruption normal quiescent phase of T Pyx.

Starting in 1996, Stubbings began a monitoring program onT Pyx, accumulating 2004 visual estimates of T Pyx. In his15 yr of seeking the next eruption, T Pyx was checked morefrequently than once every two nights, all visually with theAAVSO sequence and a 16 inch Newtonian reflector telescopefrom Tetoora Road Observatory.

Starting in early 2003, Dillon began a long-term monitoringcampaign on T Pyx, accumulating 125 magnitude estimatesbefore the start of the eruption. These magnitudes were takenwith a CCD camera through a V filter with the AAVSO sequenceof comparison stars.

Starting in 2009 March, Linnolt included T Pyx in hisprogram of regularly checking for eruptions of RNe. Hisobservations were from several sites on different islands in

the Hawaiian Islands. He used a 20 inch aperture Newtonianreflector telescope operating visually at 500× magnification,so that observations down to V = 17 were possible. Before heobtained his 20 inch telescope, his use of a smaller telescopeallowed him to see T Pyx crowded by the nearby star, so hereported only an upper limit because he could not confidentlyresolve T Pyx as a separate star.

From 2005 to 20011, Schaefer has used the SMARTS tele-scopes on Cerro Tololo to monitor the slowly fading quiescentmagnitude of T Pyx. From 2005 to 2009, BVRI magnitudes weremeasured with CCD images taken with the SMARTS 1.3 m tele-scope. There was a hiatus in 2010 due to all the observationalresources being focused on the eruption of U Sco. The recom-mencement of monitoring in 2011, made with the SMARTS0.9 m telescope, fortuitously caught the first rise of the T Pyxeruption.

From 2010 July until the eruption, ASAS covered T Pyxfrequently as a regular part of its sky survey. ASAS uses atelescope with a 200 mm diameter f/2.8 lens located at the LasCampanas Observatory in northern Chile (Pojmanski 2002). Theimages were taken through a Bessel V filter, and magnitudes aretied to a large number of comparison stars in the field withmagnitudes in the Tycho catalog. This procedure can lead tomoderate offsets for an object with unusual color (like T Pyx).Nevertheless, the differential photometry is excellent to withinthe statistical uncertainties (typically 0.09 mag). From this, wehave 40 V-band magnitudes, critically including a confirmationof the pre-eruption rise.

All our pre-eruption photometries are presented in Table 2.The first column gives the Julian Date of the observation (orthe middle of the CCD exposure), the second column gives theband, the third column gives the magnitude and its 1σ errorbar, while the last column gives the source of the magnitude.We have also included some scattered observations from theAAVSO database, with the observer being indicated by theirthree-letter observer code in parentheses. We have also includeda variety of published T Pyx magnitudes in quiescence, primarilyfor the convenience of readers so that they do not have toscramble around the literature. The result is that Table 2 contains3850 magnitudes for T Pyx in quiescence, and this representsessentially all the useable quiescent light curves from 1890 to2011. The total number of magnitudes in Table 2 (includingquiescence, the pre-eruption rise, and the initial rise untilJD 2,455,668.0) is 6116. The entire table is available only inthe online version of the journal, while the printed paper hasonly the first 10 lines to show format, plus the middle lines thatcover the critical pre-eruption rise and the initial fast rise. Table5 of Schaefer (2010) has 603 magnitudes of T Pyx in eruptionfrom 1967 and earlier.

The entire B-band quiescent light curve from 1890 to 2011is plotted in Figure 1, with the vertical lines indicating thetimes of eruptions. Importantly, we see a steady decline inquiescent brightness from before the 1890 eruption (with B =13.8) up until the 2011 eruption (with B = 15.7). With theB-band brightness depending entirely on the accretion light inthe system, we see that the accretion rate has been falling offdramatically over the 122 yr interval. This falloff is critical forunderstanding the changing time between eruptions (Schaefer2005; Section 5.4), for deriving the evolutionary status of T Pyx(Schaefer et al. 2010a), and for knowing that T Pyx will notbecome a Type Ia supernova (Schaefer et al. 2010a).

The B-band light curve of T Pyx in quiescence from 1967to 2011 is plotted in Figure 2. The earliest points show T Pyx

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Table 2T Pyx in Quiescence (1890–2011) and the Initial Rise of the 2011 Eruption

Julian Date Band Magnitude Source

2,411,491 B 13.80 ± 0.15 HCO plates2,411,493.6 B 13.80 ± 0.30 HCO plates2,412,211 B 14.10 ± 0.15 HCO plates2,412,236 B 14.30 ± 0.15 HCO plates2,412,936 B 14.30 ± 0.15 HCO plates2,413,570 B 14.30 ± 0.15 HCO plates2,413,891 B 14.30 ± 0.15 HCO plates2,414,424 B 14.50 ± 0.15 HCO plates2,414,454 B 14.70 ± 0.15 HCO plates2,414,766 B 14.70 ± 0.15 HCO plates. . .

2,455,635.1653 Vis. 15.40 ± 0.15 Stubbings2,455,641.7291 ASAS-V 15.16 ± 0.09 ASAS2,455,643.7340 ASAS-V 15.59 ± 0.10 ASAS2,455,645.7339 ASAS-V 15.03 ± 0.09 ASAS2,455,647.7350 ASAS-V 15.12 ± 0.08 ASAS2,455,649.0903 Vis. 15.30 ± 0.15 Stubbings2,455,650.0708 Vis. 15.40 ± 0.15 Stubbings2,455,650.8194 Vis. 15.20 ± 0.15 Linnolt2,455,652.1556 Vis. 15.40 ± 0.15 Stubbings2,455,652.5797 ASAS-V 15.18 ± 0.10 ASAS2,455,654.5972 ASAS-V 15.00 ± 0.09 ASAS2,455,656.6210 ASAS-V 14.69 ± 0.08 ASAS2,455,657.0104 Vis. 14.50 ± 0.15 Stubbings2,455,657.0153 Vis. 14.50 ± 0.15 Stubbings2,455,657.0222 Vis. 14.50 ± 0.15 Stubbings2,455,657.0285 Vis. 14.40 ± 0.15 Stubbings2,455,657.0361 Vis. 14.50 ± 0.15 Stubbings2,455,657.0424 Vis. 14.60 ± 0.15 Stubbings2,455,657.0993 Vis. 14.70 ± 0.15 Stubbings2,455,657.1354 Vis. 14.70 ± 0.15 Stubbings2,455,657.9667 Vis. 14.70 ± 0.15 Stubbings2,455,658.6256 ASAS-V 15.07 ± 0.08 ASAS2,455,659.0007 Vis. 14.80 ± 0.15 Stubbings2,455,660.6229 ASAS-V 15.26 ± 0.08 ASAS2,455,662.0417 Vis. 15.00 ± 0.15 Stubbings2,455,664.8590 Vis. 14.50 ± 0.15 Linnolt2,455,665.4752 B 12.51 ± 0.01 Schaefer2,455,665.4986 B 12.50 ± 0.01 Schaefer2,455,665.7931 Vis. 13.00 ± 0.15 Linnolt2,455,665.8847 Vis. 12.20 ± 0.15 Plummer2,455,665.9313 Vis. 12.00 ± 0.15 Plummer2,455,665.9410 Vis. 11.30 ± 0.15 Kerr2,455,665.9667 Vis. 11.50 ± 0.15 Plummer2,455,665.9681 Vis. 11.10 ± 0.15 Kerr2,455,665.9938 Vis. 11.10 ± 0.15 Kerr2,455,665.9944 Vis. 11.40 ± 0.15 Stubbings2,455,665.9973 V 11.54 ± 0.01 Nelson2,455,665.9978 V 11.55 ± 0.01 Nelson2,455,665.9985 V 11.56 ± 0.01 Nelson2,455,665.9990 V 11.56 ± 0.01 Nelson2,455,665.9997 V 11.54 ± 0.01 Nelson· · ·

(This table is available in its entirety in a machine-readable form in the onlinejournal. A portion is shown here for guidance regarding its form and content.)

somewhat above quiescence in the tail of the 1967 eruption.When photometry picks up in 1976, we see a flat light curve,with the usual flickering superposed. From 1976 to 2004.5, theaverage B magnitude is 15.50, with a formal 1σ error bar (theobserved rms scatter divided by the square root of the numberof observations) of 0.01 mag. After 2004.5, the quiescentbrightness appears to have taken a step fainter, with this same

Figure 1. T Pyx in quiescence since 1890. This light curve has all the B-bandmagnitudes from Table 2, with the vertical lines indicating the times of the sixhistorical eruptions. The key point is that T Pyx is suffering from a large andhighly significant secular fade from B = 13.8 in the quiescence before the 1890eruption to B = 15.7 just before the 2011 eruption. This fairly steady declineis by 1.9 mag, a factor of 5.7× in flux, which shows that the accretion ratehas declined by at least a factor of 5.7 over this 122 yr interval. The cause forthis decline is interpreted as due to a fading of the white dwarf after its 1866classical nova eruption, with the white dwarf light irradiating and puffing upthe companion star, leading to weirdly high accretion rate in the system. T Pyxappears to be going into hibernation now, although it is unclear how far downthe accretion rate will ultimately fall to. This falling accretion rate explains whythe inter-eruption intervals have been increasing over time, simply because ittakes longer and longer for the white dwarf to accumulate the necessary triggermass. This secular decline has the implication that T Pyx will not become aType Ia supernova.

Figure 2. T Pyx in quiescence from 1968 to 2011. The earliest points in this lightcurve, from 1969, show T Pyx in the late tail of its 1967 eruption. The nightswith time-series photometry have been displayed with only the nightly average.The B-band photometry from 1976 until 2004.5 shows a flat line, with the usualflickering, at B = 15.50 ± 0.01. Both the B-band and V-band photometry showthat T Pyx had an apparently sudden drop in brightness around the middle of2004, fading to B = 15.68±0.01. The light from T Pyx is essentially all from theaccretion process, so a sudden drop in brightness can only come from a suddenchange in the accretion rate. The steady orbital period change is proportionalto the accretion rate, so this light curve implies that P will have a small-stepchange starting in 2004.5.

effect visible in the V-band light curve. From the middle of 2004until early 2011, the average is B = 15.68 ± 0.01. We know ofno special event in 2004, but this appears to be just a normalpart of the star fading since 1890.

We have used our full table of quiescent magnitudes todetermine the average B-band magnitude between eruptions forsix inter-eruption intervals from before the 1890 eruption up to

5


Table 3T Pyx Inter-eruption Brightness and Flux

Interval ΔT Nights 〈B〉〈FB 〉〈FB 〉ΔT 〈F 1.13B 〉ΔT

(yr) (mag)

Pre-1890 · · · 2 13.80 ± 0.13 7.59 ± 0.99 · · · · · ·1890–1902 11.86 12 14.38 ± 0.05 4.52 ± 0.20 54 ± 2 65 ± 31902–1920 17.92 7 14.74 ± 0.05 3.20 ± 0.15 57 ± 3 67 ± 41920–1944 24.62 68 14.88 ± 0.03 2.88 ± 0.08 71 ± 2 82 ± 21944–1967 22.14 28 14.70 ± 0.04 3.35 ± 0.11 74 ± 2 87 ± 31967–2011 44.33 336 15.59 ± 0.01 1.47 ± 0.011 65 ± 1 69 ± 1

the 1967–2011 interval. The V magnitudes are converted to Bmagnitudes by using the long-term average of B − V = +0.12(Schaefer 2010). (The R and I magnitudes are few in numberand usually nearly simultaneous with B and V magnitudes, sothe R and I magnitudes were not used.) Nightly averages areformed, so that a few nights with many magnitudes as part ofa time series (which might or might not be representative ofthe whole interval) do not dominate the average. Later in thispaper, we will be needing an average of the B-band flux (notmagnitude), so we have defined a unit of flux equal to the fluxwhen T Pyx is at B = 16.00 mag. That is, FB = 10−0.4(B−16).These fluxes are then computed for each nightly average, andthen combined as an average over each inter-eruption interval togive 〈FB〉. The uncertainty for each average is taken to be the rmsscatter divided by the square root of the number of nights. (If thenightly averages are first formed into yearly averages and theseare averaged, then we get fluxes which are essentially identical.)Table 3 gives the pre-eruption intervals, their durations (ΔTfrom Schaefer 2010), the average B-band magnitudes over eachinterval (〈B〉), the average FB for those intervals, the values of〈FB〉ΔT (which should be nearly constant), and the values of〈F 1.13

B 〉ΔT . This last column will be discussed in Section 5.4.

3.2. Pre-eruption Minima Times

On five nights in 2008 January, Richards made long time-series photometry on T Pyx, all with 100 s integrations ona 0.41 m Ritchey-Chretein reflecting telescope at Pretty HillObservatory near Melbourne, Australia. The images were alltaken with an unfiltered CCD, with the T Pyx magnitudebeing taken by differential photometry with respect to theV-band magnitudes of the comparison stars (designated as “CV”band), so the effective magnitudes are similar to the V-bandbut nevertheless different. This non-standard band is perfectfor getting the shape of the light curve, as needed to time theminima.

On three nights in 2011 February, Myers made long time-series photometry of T Pyx. The reason for this photometry, andfor the nearly contemporaneous time series of Nelson, was asa support for observations of T Pyx with the Chandra X-raysatellite (Balman 2012). The CCD images were made withouta filter (in the “CV” band) with 120 s exposures, with readouttimes of 64 s, with a 0.25 m telescope. The observing site is inMayhill, New Mexico.

On five nights in 2011 February, Nelson made long time-series photometry on T Pyx using a 0.32 m corrected Dall-Kirkham telescope at Ellinbank Observatory, in Victoria, Aus-tralia. A total of 192 magnitudes were made with an unfilteredCCD and V magnitudes for the comparison stars (hence, the“CV” band). The integration times were always 120 s, resultingin an average time between exposures of 170 s.

The magnitudes from these pre-eruption time series areincluded in Table 2. The magnitudes were grouped togetherinto five intervals, each from a single observer, and thesegroups were then used to measure the times of minimum inT Pyx’s light curve. We have performed chi-square fits ofthe folded light curve to a template. The folded light curveis folded around the period from the parabolic ephemeris ofEquation (1) (see below). (We are dealing with time seriesof such short length that any approximate period would havebeen adequate.) The template we used was a flat light curvefrom phases 0.25–0.75 with a linear and symmetric decline toa minimum at zero phase of 0.15 mag fainter (see Figure 10 ofPatterson et al. 1998). The derived times of minima were thencorrected to Heliocentric Julian Dates (HJDs). This procedureyields a formal uncertainty on the time of minimum, but this onlyaccounts for the measurement errors. The uncertainty is actuallydominated by the orbit-to-orbit variations in the light curve (seeFigure 54 of Schaefer 2010 and Figure 8 of Patterson et al. 1998)caused by the usual flickering plus other unstable periodicities.When one orbit is considered, the light curve occasionally doesnot have an obvious minimum. However, when many orbitsare considered together, the combined light curve always lookssimilar to the template. So the accuracy of the times of minimashould scale as the inverse square root of the total duration ofdata in the time series in terms of the orbital period. For a groupof similar observations over a relatively small time range, thescale of this uncertainty can be set by looking at the scatterin the O − C curve about some smooth best-fit line. So, forexample, the uncertainty in the minimum times reported byPatterson et al. (1998) can be calculated as 0.0033(Δ/0.076)−0.5

days, with Δ being the total duration of the observation (notnecessarily contiguous) in days. With this, Table 4 presents theavailable times of minimum (in HJDs) with their uncertainties.Included are six times of minimum from Patterson et al. (1998).

3.3. Photometry of the Initial Rise

The pre-eruption rise was covered by the visual observationsof Stubbings and the CCD observations of ASAS, as describedin Section 3.1. The initial fast rise was recorded by the SMARTSCCD images of Schaefer and the visual discovery observationsof Linnolt, as described in Section 3.1. A blowup of the lightcurve around the time of the pre-eruption rise and the initial fastrise is displayed in Figure 3.

Immediately after the announcement of the discovery by Lin-nolt, we began an intensive set of time-series images in the B, V,R, and I bands. This intensive set of observations has continuedall the way through the peak, and indeed continues with dailylong time series until the time of this writing (2013 February).Plummer, Kerr, and Stubbings provided immediate visual con-firmation of Linnolt’s discovery. Nelson, Carstens, and Streamerimmediately began all-night fast time-series photometry in the

6


Table 4T Pyx Minimum Times

Year Observer Δ HJD Minimum N Linear O − C Quadratic O − C(days) (days) (days)

1986.022 Warnera 0.226 2,446,439.5121 ± 0.0035 −68378 0.1267 0.00731987.943 Szkodya 0.091 2,447,141.0770 ± 0.0056 −59174 0.0960 0.00651988.347 Schaefera 1.038 2,447,288.5651 ± 0.0017 −57239 0.0844 0.00061988.840 Schaefera 0.154 2,447,468.9843 ± 0.0043 −54872 0.0737 −0.00331989.053 Bonda 0.183 2,447,546.6591 ± 0.0039 −53853 0.0729 −0.00121989.537 Schaefera 0.196 2,447,723.4983 ± 0.0038 −51533 0.0649 −0.00301990.342 Buckleya 0.459 2,448,017.3506 ± 0.0025 −47678 0.0611 0.00301990.348 Walkera 0.193 2,448,019.6358 ± 0.0038 −47648 0.0595 0.00141996.112 Kempb 3.667 2,450,124.8300 ± 0.0005 −20030 0.0095 −0.00101996.220 Kempb 0.892 2,450,164.3930 ± 0.0009 −19511 0.0106 0.00061996.289 Kempb 0.328 2,450,189.5480 ± 0.0015 −19181 0.0106 0.00091996.352 Kempb 0.386 2,450,212.4910 ± 0.0014 −18880 0.0092 −0.00021997.031 Kempb 1.784 2,450,460.6090 ± 0.0007 −15625 0.0075 0.00101997.272 Kempb 0.686 2,450,548.4970 ± 0.0011 −14472 0.0055 −0.00012008.000 Richardsc 0.160 2,454,467.1391 ± 0.0015 36935 0.0334 −0.00162008.019 Richardsc 0.220 2,454,474.0788 ± 0.0013 37026 0.0364 0.00122008.058 Richardsc 0.292 2,454,488.1785 ± 0.0011 37211 0.0341 −0.00152011.099 Myersc 0.148 2,455,598.7619 ± 0.0016 51780 0.0627 −0.00592011.118 Nelsonc 0.297 2,455,606.0039 ± 0.0011 51875 0.0631 −0.0057

Notes.a Schaefer et al. (1992).b Patterson et al. (1998).c Table 2.

Figure 3. Pre-eruption rise. T Pyx never appeared brighter than V = 15.0 forthousands of days before the eruption, and then it suddenly brightened to V =14.4 mag just 9 days before the eruption. The close time coincidence and theuniqueness of the bump demonstrate some causal connection between the bumpand the eruption. The decline before the very fast rise demonstrates that thissignificant bump in the light curve is not some early part of the thermonuclearrunaway. There is no understanding of the cause of the bump or why it anticipatesthe eruption. All of the magnitudes on this plot are V-band magnitudes, exceptfor the pair of B-band magnitudes circled in the upper right corner of the plot.These two magnitudes are serendipitously taken 7.6 hr before the discoveryof the eruption by Linnolt (indicated by the large square), with the earlierobservations being 0.5 mag brighter than the later discovery. We are left witheither that the initial fast rise had a large amplitude dimming just hours beforediscovery or that T Pyx was extremely blue (B − V = −1.0) in its earliestmoments. Explaining the pre-eruption phenomena is now a premier challengefor nova theory.

V band. These time series provide a unique record of the fastinitial rise from the earliest time, for direct comparison withtheory. Streamer also took parallel time series in B, V, and I,so we have a totally unique complete record of the color devel-opment of the fast initial rise from the earliest times. Further

Figure 4. Initial rise. T Pyx now has the only nova light curve where thebrightness was followed over the entire initial rise. For the first time, we see thatthe initial rise starts suddenly and is very fast. The rise is as fast as 9 mag day−1.A uniformly expanding shell is expected to rise fast in brightness followed bya substantial slowing (in magnitude units) after the first day, with the best-fitmodel displayed by the curve in the figure. We see a sudden small fading andbreak from the model about two days after the start, with this perhaps indicatingthe recession of the photosphere.

observations have been taken from the archives of the AAVSO.In this paper, we will only cover the pre-eruption quiescence,the pre-eruption rise, and the first 2.2 days of the eruption. Allthese magnitudes are recorded in Table 2. Figure 4 shows thelight curve of the initial fast rise.

We have large numbers of observations of time series inmultiple colors, where pairs of images through different filterswere taken within seconds of each other, and so we can easilycalculate the T Pyx colors at many times throughout the fastinitial rise. The first measures of the color (when T Pyx wasaround eleventh magnitude) have median B − V = −0.02,

7


Figure 5. B − V color over initial rise. We have the first and only measure ofthe colors of the initial rise of a nova. The color does change somewhat over thefirst two days of the eruption (covered by the model fit in Figure 3), suggestingthat the photospheric temperature does change over this interval, with this beingcontrary to the simple model of a uniform expanding shell with a constanttemperature. The initial colors (B − V = −0.02 and V − I = +0.39) arenot those of a blackbody. In addition, the SMARTS B-band magnitudes fromJD 2,455,665.4752 imply B − V = −1.0 if the earliest moments of the lightcurve are smooth and monotonic in the rise.

but this slowly changes to B − V = +0.23 by two days afterthe start. At this time, just as the light curve is nearly flat,the B − V color curve also goes nearly flat. Just after discovery(around JD 2,455,666.07), the median colors are B−V = −0.02and V − I = +0.39. These colors are inconsistent with ablackbody, as for example an A0 star with a temperature near10,000 K (B − V = 0.00) has the V − I color equal to 0.00,which is substantially different from that seen in T Pyx. Halfa day after discovery (JD 2,455,666.65), the median colors areB − V = +0.06, V − R = +0.27, and V − I = +0.37. Atthe time when the fast initial rise breaks (JD 2,455,668.0), themedian colors are B −V = +0.23 and V − I = +0.58. Figure 5shows our derived B − V color curve through the initial fast rise.

3.4. GALEX Spectrum and the Spectral Energy Distribution

We have obtained an ultraviolet spectrum of T Pyx inquiescence with the GALEX satellite. This is from GALEXCycle 2, with Schaefer as principal investigator. The obser-vation time was 2005 December 20 starting at UT 03:30:05(JD 2,453,724.646), with an effective total exposure of 2503 s.GALEX gives a grism spectrum with a far-UV (FUV) bandcovering 1344–1786 Å, and a near-UV (NUV) band covering1771–2831 Å. We have used the standard data analysis pipeline,which returns a fully calibrated combined spectrum with thespectral flux, fλ, in units of erg s−1 cm−2 Å−1. The FUV mag-nitude is 15.81 ± 0.01 and the NUV magnitude is 16.00 ± 0.01.The resultant spectrum is of good quality, although the noise atthe edges of the bands gets fairly large.

The GALEX spectrum shows a very blue continuum with noblatant emission features, plus a prominent broad dip centeredon 2175 Å. We have dereddened the spectrum by fitting forthe extinction that best removes the 2175 Å bump, for whichwe find E(B − V ) = 0.239 ± 0.004 mag. This spectrum, withmoderate binning plus clipping of several end points with verylarge error bars, is displayed in Figure 6. We see a smoothcontinuum, plus only one significant spectral line (the C iv lineat 1550 Å). The He ii λ1640 Lyα line is visible as a 2σ bump inthe GALEX spectrum, although this line is highly significant and

Figure 6. GALEX spectrum. This spectrum of T Pyx in quiescence was taken in2005 December with the GALEX satellite. The error bars are moderately higharound 1800–1900 Å due to that being the region where the NUV and FUVcameras join together. The original spectrum showed a prominent broad dipcentered on 2175 Å, due to the usual absorption by the interstellar medium,with our dereddening corrections designed to eliminate this dip. The resultantspectrum is a good power law with an index of −2.25 ± 0.03. The C iv lineat 1550 Å is significant, and the He ii line at 1640 Å is likely visible but notsignificant. This spectrum can be compared with the ultraviolet spectrum derivedfrom greatly longer exposures with the IUE satellite that covers a somewhatbroader wavelength range (Selvelli et al. 2010).

variable in the IUE spectrum (Gilmozzi & Selvelli 2007). Thisspectrum is well fit by a power law, for which we find an indexof −2.25 ± 0.03, while the normalization has a formal error of1.1%. Our best fit is fλ = 2.78 × 10−6λ−2.25 erg s−1 cm−2 Å−1,for wavelength λ measured in angstroms. For a broad spectralenergy distribution (SED), it is usual to represent the spectralflux as fν (in units of Jansky, 1 Jy = 10−23 erg s−1 cm−2 Hz−1),where we have the relation νfν = λfλ. In a plot of fν versusfrequency ν, our best-fit spectral slope is fν ∝ ν0.25. This isessentially identical to the best-fit power law from the averageIUE data (Gilmozzi & Selvelli 2007).

We have also constructed an SED for T Pyx from 1456 Å to22200 Å by using the GALEX spectrum along with optical andinfrared photometry. Schaefer (2010) measured the UBVRIJHKcolors of T Pyx on several occasions contemporaneous withthe GALEX observations and has converted these magnitudesinto fν with extinction corrections. (The original paper hadalready pointed out the U-band fν as being a likely error, andwe have now recognized that the U − V color for T Pyx hadbeen incorrectly transferred from Tables 25–30 of Schaefer(2010) without the negative sign, and so the SEDs in Figures70 and 71 show a false turn down in the ultraviolet. This erroris now corrected in this paper.) To display the SED, we plotlogarithmic scales with frequency, ν, on the horizontal axis,and νfν on the vertical axis, in part because this presentationshows the peaks where the most energy is coming out. ForT Pyx, our SED is plotted in Figure 7. We see a rising spectrumtoward the ultraviolet, so most of the T Pyx accretion energyis coming out for wavelengths shorter than 1456 Å. This willnot surprise anyone, but it does illustrate that we are missingmost of the accretion light. The overall SED rises remarkablyclose to a simple power law, even though we can see deviations,for example where the slope is different in the ultraviolet. Ourbest-fit power law for fν has an index of essentially unity, sowe have fν = 0.009Jy(ν/1015 Hz). Gilmozzi & Selvelli (2007)had a similar result with fλ ∝ λ−2.9 or fν ∝ ν0.9. This is

8


Figure 7. T Pyx spectral energy distribution (SED). This SED plot is constructedfrom the GALEX spectrum plus UBVRIJHK photometry in Schaefer (2010). Toa close approximation, the SED is a simple power law with νfν ∝ ν2 orfν ∝ ν. Within this broad overall shape, the ultraviolet portion of the spectrumis somewhat flatter, with a slope that goes as fν ∝ ν0.25. This is completelydifferent from a blackbody spectrum, as the Rayleigh–Jeans tail has fν ∝ ν2.This is also completely different from a disk spectrum, which should go likefν ∝ ν1/3 only for frequencies above the observed range (due to any diskrequiring a temperature of the outer edge of 20,000 K or hotter), while anydisk SED must have turned over to a Rayleigh–Jeans tail in most of the plottedspectral range. There is one additional point that has not been plotted, and thatis that the X-ray brightness of the central source can be represented by a pointgreatly below the plot at a frequency of ∼2 × 1017 Hz. This point forces thepower-law rise to high frequency to undergo some sort of a cutoff between2 × 1015 and 2 × 1017 Hz. Unfortunately, there is no way to better constrain thecutoff, and this leads to a factor of 10,000× uncertainty in the accretion energy.

completely different from a hot blackbody spectrum, where theRayleigh–Jeans region has fν ∝ ν2.

We can also place an X-ray point onto the SED. Balman(2012) reports that the central source has a 0.2–9.0 keV flux of(1.0–5.8) × 10−14 erg s−1 cm−2. This translates into somethinglike 10−9 Jy for a frequency around 2×1017 Hz, so νfν is around2 × 108 Hz Jy. This point is half-again off the right side of thebox in Figure 7, and far below the bottom. So we know thatthe power-law rise so prominent in Figure 7 must break sharplydownward somewhere between 2 × 1015 Hz and 2 × 1017 Hz.

This SED will give the accretion luminosity of T Pyxas L = 4πD2

∫fνdν. For a fiducial distance of 3000 pc,

the spectral range (1–20) × 1014 Hz gives a luminosity of2 × 1035 erg s−1. If the observed power-law rise to high photonenergies continues until a sharp break (as constrained by theX-ray observation), for a total range in frequency of (1–2000)×1014 Hz, we get a maximal accretion luminosity of 2 ×1039 erg s−1. Even for a set fiducial distance, this allows arather wide range of possible accretion luminosities.

4. MEASURING PROPERTIES OF THE T Pyx SYSTEM

We have used our observations, plus results from the litera-ture, to derive the properties of the T Pyx system.

4.1. The Steady Period Change P

In the classic review of RNe, Webbink et al. (1987) concludedthat T Pyx must have a relatively short orbital period (like thoseof the typical classical novae) due to the lack of any red excess.Early photometric periods suggested were 0.069 days (Szkody& Feinswog 1988) and 0.1433 days (Barrera & Vogt 1989). Witha much larger set of time-series photometry, Schaefer (1990)found a stable and highly significant modulation but had severe

daily alias problems, for which the highest peak in the Fouriertransform was at 0.099 days, while another lower peak was at0.076 days. Schaefer et al. (1992) collected an even larger set oftime series from 1966 to 1990 and determined the period to be0.07616 ± 0.00017 days. Patterson et al. (1998) made long runsfrom 1996 to 1997, confirmed this orbital period, discovered asteady period change, and discovered a complicated structure ofweaker periodicities that vary in amplitude. Uthas et al. (2010)demonstrated that the stable 0.076 day photometric periodequals the orbital period as measured with a radial velocitycurve. They also used many minimum timings from 1996 to2009 (from the records of the Center for Backyard Astronomy)to give a high-precision ephemeris that includes a parabolicterm. This gives the minimum times (expressed in Julian Dates)as

TUthas = 2,451,651.65255 + 0.076227249N

+ 2.546 × 10−11N2, (1)

where N is the number of orbital cycles from the epoch. Thequadratic coefficient has an uncertainty of 2.1%.

The observed minimum times (Tmin) in Table 4 (seeSection 3.2) can be plotted by means of an O − C diagram.For this plot, we compare the observed minus the predictedminimum times (O − C, measured in days), where the predictedtimes are taken from some linear ephemeris. We have chosen alinear ephemeris of

Tlinear = 2,451,651.65255 + 0.076227249N, (2)

which is just the Uthas ephemeris without the quadratic term.The observed-minus-calculated difference is

O − C = Tmin − Tlinear. (3)

The O − C curve is a plot of O − C versus time T. If the starobeys the linear ephemeris, then the O − C curve will show allthe minimum times to have the usual scatter along the horizontalaxis. If the period is different than given in the linear ephemeris,then the O − C curve will have the points follow along a straightline with some non-zero slope. If the period changes linearlywith time, then the slope of the plot will change linearly withtime, and this will appear as a parabola in the O − C diagram.The N and O − C for the linear ephemeris are tabulated foreach of our Tmin in Table 4. The O − C plot for our 19 Tminmeasures and the linear ephemeris from Equation (2) is givenin the upper panel of Figure 8. We see a good parabola, whichpoints to a steady period change in the T Pyx orbital period.This is not surprising, because Uthas et al. used our data forall the pre-1992 times. Our best-fit parabola to Figure 8 (upperpanel) returns an ephemeris very close to that of Uthas et al., sowe will simply adopt her ephemeris as the best measure of thesteady period change in T Pyx.

For later use in equations involving mass transfer, we needthe steady period change, P , cast into a dimensionless form.This period derivative can be calculated from Equation (1), withthe period over each orbit equaling the time difference betweensuccessive minima, so P equals twice the quadratic coefficientin Equation (1) divided by the orbital period in days. Thus,P = 6.68 × 10−10, with an uncertainty of 2.1%. A useful wayto express this is P/P = 313,000 yr, about a third of a millionyears.

4.2. Orbital Period Just Before the Eruption Ppre

A highly accurate pre-eruption orbital period (Ppre), togetherwith a post-eruption period, will give the change in orbital period

9


Figure 8. T Pyx O − C curve. T Pyx has a stable periodic modulation in itslight curve, and we have minimum times from 1986 until 2011 February (just48 days before the eruption). The upper panel shows the O − C curve for alinear ephemeris (Equation (2)) in which we see a closely parabolic shape. Thisindicates that T Pyx has a steady change in orbital period, just as expectedfor steady conservative mass transfer at a rate of (1.7–3.5) × 10−7 M� yr−1.The lower panel shows the O − C curve for the predicted times coming fromthe quadratic ephemeris (see Equation (1) of Uthas et al. 2010). This steadyperiod change well describes the minimum times from 1988 to 2008. But thetwo minimum times from 2011 February are consistent and highly significantlybelow the predictions from a single parabola, so apparently T Pyx underwentsome sort of a small P change sometime around 2008, such that its orbitalperiod just before the 2011 eruption is slightly different (7.5 parts per million)from the period given by Equation (1). Such a change in P is expected due tothe observed 18% decrease in system brightness (and hence an 18% decrease inaccretion rate) from 2004.5 to 2011.

across the 2011 eruption. This period change will then givethe ejected mass with high accuracy and no need for uncertainquantities (like the distance) or dubious assumptions (like theshell filling factors). For the case of T Pyx (but not for theother RNe), the dynamical friction of the companion star withinthe nova envelope will also cause a small period change (Livio1991; Schaefer et al. 2010a), and this effect must also be takeninto account. The ejected mass is critical for knowing whetherthe RN white dwarf is gaining or losing mass over the eruptioncycle, and hence whether the RN will soon become a Type Iasupernova. Previously, only two other RNe (U Sco and CI Aql)had a measure of the period change across the eruption (Schaefer2011). The T Pyx eruption of 2011 provides an opportunity tomeasure the period change for a third RN.

As the period of T Pyx changes, it is vital to determine thepre-eruption orbital period for the time immediately precedingthe eruption. To good accuracy, Equation (1) should work well

to give the pre-eruption orbital period. With this, for an epochof N = 52,662, the pre-eruption orbital period is 0.07622993 ±0.00000017 days. However, this period is based on minima timesonly until 2009. We now have the time series in 2011 February,up to 48 days before the eruption. The possible existence ofsmall period changes just before the eruption is critical wherewe need part-per-million accuracy.

To search for deviations from the parabolic ephemeris, weconstructed another O − C curve where the ephemeris wasEquation (1). That is, O − C = Tmin − TUthas. The resultantquadratic O − C values are presented in the last column ofTable 4, and the O − C diagram is shown in the lower panelof Figure 8. If the Uthas et al. (2010) model is perfect, thenall the observed minimum times should lie along the horizontalaxis with O − C = 0 to within the plotted error bars. Indeed,we see that the early time series from Schaefer et al. (1992),Patterson et al. (1998), and from early 2008 by Richards all fallclose to the axis. This is to say that the parabolic ephemeris isa good match to the observed times from 1988 to 2008. Thetwo times in 1986 and 1987 are roughly 2σ and 1σ above theline, with a weak suggestion of deviation from Equation (1)before 1988, although we judge this suggestion to be too weakto accept. However, the two times from 2011 February are bothclosely consistent and with 3.7σ and 5.1σ deviations from theparabolic ephemeris. These two measures indicate that T Pyxunderwent a very small but significant change in period from2008 to 2011.

If the O − C curve is linear from 2008 to 2011, then theorbital period (gotten from just our five minimum times in2008 and 2011) is Ppre = 0.07622916 ± 0.00000008 days.But the steady period change is caused by the mass transfer inthe system, and we know that this has not stopped (T Pyx wasstill flickering in early 2011), so the shape of the O − C curvemust still be a parabola from 2008 to 2011. In all scenarios, Pis proportional to M , and from Section 5.3 we have that 〈FB〉is proportional to M . For a magnitude difference of 0.18 mag(see Section 3.1 and Figure 2), the blue-band flux ratio is 0.85,so both M and P should also have that same ratio for the twotime intervals. With this, for the last few years before the 2011eruption, P = 5.66 × 10−10, and P/P = 366,000 yr. We onlyhave minima times in 2008 and 2011, so with the derived P ,we have enough to state a quadratic ephemeris for T Pyx in theyears just before eruption. We then fit the times to get

Tpre-eruption = 2,455,665.9962 + 0.07622950N2011

+ 2.16 × 10−11N22011. (4)

Again, the coefficient of the quadratic term is just P P /2.Here, the epoch was chosen to be the time of minimumclosest to the start of the eruption. That is, the start of the2011 eruption corresponds to N2011 = 0. With this, we havePpre = 0.07622950 ± 0.00000008 days.

4.3. Extinction E(B − V )

The extinction of T Pyx is required for various attempts toget the distance, luminosity, and accretion rate. The extinctioncan be quantified by either the color excess, E(B − V ), orthe V-band dimming in magnitudes, AV . These are related asAV = RV E(B − V ). For the nearby and normal sightline toT Pyx, we can be very confident in using the canonical valueof RV = 3.1. For a given RV , the dimming in magnitudes overultraviolet/optical/infrared wavelengths is accurately given by

10


the analytic equation of Cardelli et al. (1989), as a function onlyof the parameter E(B − V ).

Various values for the extinction to T Pyx have been reportedin the literature. Bruch et al. (1981) used an IUE spectrumof T Pyx to derive E(B − V ) = 0.35 ± 0.05 mag as basedon the broad 2175 Å dip. A later and better analysis of manymore IUE spectra with the same method yielded E(B − V ) =0.25±0.02 mag (Selvelli et al. 2003; Gilmozzi & Selvelli 2007).Weight et al. (1994) quote E(B − V ) = 0.08 mag, but do notmention the method or source of this result.

Shore et al. (2011) find E(B − V ) ≈ 0.5 ± 0.1 as based onthe strength of diffuse interstellar bands and some correlationsbetween these strengths with E(B − V ). However, their textgives a different value from their abstract, reporting insteadE(B − V ) = 0.49 ± 0.19, with the difference of a factor of twobetween the error bars apparently being due to the averaging ofE(B−V ) values from four bands. (Our detailed repetition of thisprocedure yields an average extinction of 0.47 mag for T Pyxwith their reported equivalent widths and their choice of fourbands.) But this improvement by a factor of two by averagingover four bands is not valid, because the equivalent widths of thebands are very tightly correlated (Friedman et al. 2011) so theyare not independent. Indeed, the uncertainty should not be givenby the fractional rms scatter of the not-independent values fromthe four bands. A good way to calculate the real uncertaintyis to repeat the procedure of Shore et al. for the many stars inFriedman et al., and calculate the rms scatter of the averagederived E(B − V ) values for those stars with similar E(B − V )as claimed for T Pyx. With this, we find that the 1σ error baris 49%. Thus, the data and method of Shore et al. are actuallyproducing E(B − V ) = 0.47 ± 0.23.

Here, we report on a new measure of the extinction to T Pyx.This is based on the GALEX spectrum, where the 2175 Ådip is prominent. We have dereddened the fluxed and binnedspectrum using the extinction model in Cardelli et al. (1989).With a chi-square fit, we find the E(B − V ) value for whichthe resultant spectrum from 1440 to 2844 Å is best fit by somepower law. For this fit, we have not included the bins withthe C iv emission line at 1550 Å or the He ii emission lineat 1640 Å. Our best-fit dereddened spectrum and the best-fitpower-law model are shown in Figure 6. With this, we getE(B − V ) = 0.239 ± 0.004 mag. This is essentially the sameresult from the same method as Selvelli et al. (2003), except thatwe are independent of them and have used a different satellitefor the ultraviolet spectrum.

The standard way to determine the total extinction alonga line of sight is to use the far-infrared maps of IRAS andCOBE/DIRBE as given by Schlegel et al. (1998). Their mapsshow T Pyx to be on the edge of the Milky Way. Their extinctionalong the entire line of sight (for an assumed value of RV = 3.1)is E(B − V ) = 0.275 mag (with a range of 0.253–0.290 mag).This sets a confident upper limit on the possible interstellarextinction to T Pyx. So we know that the value reported byShore et al. is roughly 1σ high, which is to say about a factorof two high. This extinction limit is equal to that measured forT Pyx. This means, not surprisingly, that T Pyx is outside mostor all of the interstellar medium around the plane of our galaxythat is causing the extinction.

So what is the extinction to T Pyx? The only confident andaccurate method for determining the extinction to T Pyx is basedon the 2175 Å absorption in the ultraviolet spectrum. For this,the best measure is that by Selvelli and coworkers. So we adopttheir E(B − V ) = 0.25 ± 0.02 mag.

4.4. Distance D

Distances are always critical in astronomy, and the distance toT Pyx, D, will be needed to get the luminosity and accretion rate.Prior work has not turned up any decisive measure of D, althoughneeding some answer, a variety of methods have been applied.All papers hearken back to the original distance estimate ofCatchpole (1969), who used three different means, involvingthe strength of the interstellar calcium absorption line, thevelocity of the interstellar calcium line, and the rate of decline inbrightness. Additional methods key off the ultraviolet brightnessin comparison with some model, as well as observations of thenova shell. All of these methods have variations, such as tothe lines and the data source used. With various combinationsof these methods and differing input data, many authors havecome up with distances from 1000 pc up to �4500 pc (Catchpole1969; Webbink et al. 1987; Patterson et al. 1998; Selvelli et al.2008; Sion et al. 2010; Schaefer 2010; Shore et al. 2011).Unfortunately, with recent results and new analysis, all thesedistance determinations are seen to be either wrong or withsuch large uncertainties as to make them useless.

Catchpole (1969) started with the method that gets thedistance from the equivalent widths of the interstellar calcium Kline and the interstellar sodium D lines, EW(Ca) and EW(Na),both measured in kilometers per second. He used a calibrationfrom Beals & Oke (1953) that D = 34.83 × EW(Ca) and D =30.75 × EW(Na), with the distance in parsecs. Unfortunately,this method, and all similar analyses giving distances based oninterstellar absorption, suffer from one large problem and oneerror, either of which robs the method of the ability to giveany useful distance. The large problem is that the calibration ofmany sight lines shows a wild variation in the ratio of distanceto absorption. The 1σ scatter in the distance for some observedequivalent width is half an order of magnitude (e.g., Hunteret al. 2006), so the 1σ range in derived distance will be one fullorder of magnitude in distance. That is, some heavily averagedrelation might be linear, but any one star can vary around thisaverage by large factors, with the implication that any absorptiondistance to T Pyx has an uncertainty of around an order ofmagnitude. The error arises because the linear relation betweenabsorption and distance requires that the entire line of sight lieclose to the galactic plane. The calibration of these relationsuses sources almost all of whose whole lines of sight lie within100 pc of the plane (Beals & Oke 1953; Munch 1968; Hunteret al. 2006; Friedman et al. 2011). The scale height of interstellarabsorbers is around 150 pc, so any sightline extending outsidethis height above the galactic plane will be in error. T Pyx is ata galactic latitude of +9.◦7, so its sightline will rise 150 pc abovethe plane for a distance of 900 pc. For all D > 2600 pc, thesightline will have passed outside three scale heights and havebroken into regions that are free of interstellar absorption. Thatis, the equivalent widths in interstellar lines will be identicalfor distances of 2600 pc, 3500 pc, 4500 pc, or 100,000 pc.Indeed, given the thin nature of the interstellar medium morethan ∼200 pc above the plane as well as the unknowable naturalvariations along any line of sight, the absorption at 1200 pc isindistinguishable with that at 100,000 pc. The procedural error(assuming a linear distance–absorption relation for a distantsource at a galactic latitude of +9.◦7) invalidates all publisheddistances involving the strength of interstellar absorption. It isdisconcerting that this mistake was made in the first place, thenrepeated many times over the past 44 yr, and was never caughtuntil now. There is precedent for this debacle in the case of the

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recurrent nova RS Oph, where Hjellming et al. (1986) made thesame mistake (with RS Oph at a galactic latitude of +10.◦4) soas to get an erroneous distance of 1600 pc, with this distancebeing uncritically adopted in paper after paper by the bandwagoneffect (cf. Schaefer 2008) so that most papers in the past decadegive 1600 pc as the canonical distance by dint of sheer repetition(Schaefer 2009). For T Pyx, the bottom line is that all distancesinvolving the strength of interstellar absorption are now knownto be invalid.

Catchpole (1969) also was the first to try to use the velocitystructure of the interstellar absorption lines to get the distanceto T Pyx. Alas, again, this distance determination and all similardeterminations are fatally flawed. Both Catchpole (1969) andSelvelli et al. (2008) noted that the galactic coordinates of T Pyx(longitude 257.◦2 and latitude +9.◦7) are near to a node of therotation curve, so both declined to use the method becausethe uncertainty will inevitably be large. Shore et al. (2011)note that the interstellar absorption line for CH has a radialvelocity of +49.5 km s−1 and derived a minimum distance of4000 pc. Unfortunately, they had made an error in velocityrequiring a correction of −19.6 km s−1, so in a publishedcorrigendum, they give the CH line with a real radial velocity of+29.9 km s−1 as the highest velocity absorption feature in theT Pyx spectrum. No analysis is given as to the updated distanceestimate in their published corrigendum, although Shore et al.(2013) have modified the claimed distance from their first paperfrom the original �4500 pc to �3500 pc. These limits wereobtained by taking this radial velocity and formally applyinga galactic rotation curve. Unfortunately, such a procedure doesnot account for normal variations in cloud velocities from thegalactic average. In particular, nearby stars in the same directionhave clouds absorbing up to +50 km s−1. To take three nearbyexamples (Hunter et al. 2006), HD 72067 (galactic longitude262.◦1 and latitude −3.◦1) is at a distance of 488 pc and has anintervening cloud at +30 km s−1, HD 74966 (galactic longitude258.◦1 and latitude +3.◦9) is at a distance of 658 pc and has anintervening cloud at +50 km s−1, while HD 68761 (galacticlongitude 254.◦4 and latitude −1.◦6) is at a distance of 1473 pcand has an intervening cloud at +50 km s−1. So, T Pyx, withan intervening cloud at +29.9 km s−1, needs only be furtheraway than roughly 500 pc. That is, this whole method can onlyproduce a meaningless lower limit.

Many papers have used the method called the “maximummagnitude–rate of decline” (MMRD) relation, to go from thespeed class of the light curve (quantified by t3 or t2) to get thepeak absolute magnitude, and then the distance. Unfortunately,the observed scatter in the MMRD for novae in our Milky Wayhas a total range of 3 mag for decline rates near that of T Pyx(Downes & Duerbeck 2000), and this translates into a distanceuncertainty of a factor of four. So even if we take the MMRDat face value, we are left with too large an uncertainty to beuseful. However, there is substantial evidence that the MMRDcannot successfully be applied to RNe, in particular, 3 of the 10galactic RNe (V2487 Oph, U Sco, and V394 CrA) have MMRDdistances far outside our Galaxy, while 2 RNe (U Sco and RSOph) have MMRD distances in large disagreement with thereliable blackbody distances to the companion stars (Schaefer2010). So the MMRD is not confidently applicable to RNe,and T Pyx is such a weird and unique system that the physicalconditions of the eruption are likely to be so unusual that theMMRD does not apply. To make matters much worse, Kasliwalet al. (2011) have found that the novae in M31 show only avery wide scatter with no hint of an MMRD in any form. These

spectroscopically confirmed M31 novae are generally faster thanT Pyx in decline rate, but this cannot hide the fact that most M31novae are far from the MMRD relation. Recently, this sameresult has been found for another galaxy. With this, the veryexistence of the MMRD is disproven. There is no MMRD fornova sets with known distances, while our Milky Way showedan MMRD, so we think that this implies that distances to novaein our Milky Way are subject to substantial bandwagon effects.In summary, all MMRD distances have a uselessly large realuncertainty and the best idea is that the MMRD neither appliesto T Pyx nor any other nova.

A similar distance method, with a physical basis instead of anempirical basis for the absolute magnitude at peak, is to presumethat the nova cannot get much above the Eddington luminosity.Selvelli et al. (2008) calculate that novae with high-mass whitedwarfs have an Eddington luminosity that corresponds to anabsolute magnitude of −6.78. With this assumption, we canindeed derive a distance. Unfortunately, we know that theassumption of the Eddington luminosity as a limit is greatlywrong, even for slow novae. Kasliwal et al. (2011) show that80% of novae in the Andromeda galaxy violate the Eddingtonlimit by factors from 2 to 16.

Both Selvelli et al. (2008) and Sion et al. (2010) have deriveddistances to T Pyx based on the de-reddened ultraviolet fluxas compared to the models of Wade & Hubeny (1998). Thereare many large technical problems with this program. First,the model of Wade & Hubeny only goes up to white dwarfmasses of 1.21 M� and up to 10−8.5 M� yr−1, thus requiring farextrapolations. Second, the best-fit model has a greatly differentpower-law slope than is observed, and this shows that the modelis not applicable to T Pyx. Third, the ultraviolet luminosityhas a strong dependence on the accretion rate, which mightonly be known to within a factor of two or so, so any deriveddistance is merely returning the worker’s assumption on theaccretion rate. Fourth, the distance derived by Sion et al. (2010)is 1000 pc, while the distance derived by Selvelli et al. (2008) is∼3150 pc. That both were made with essentially identical dataand technique demonstrates that the real accuracy of the methodis a factor of three or worse. Fifth, the model presumption isthat the ultraviolet luminosity is provided by a simple accretiondisk, but this presumption is violated by T Pyx with the casethat a magnetic field from the white dwarf is disrupting anydisk (Schaefer & Collazzi 2010; Oksanen & Schaefer 2011)and with the case that there is continuing nuclear burning on thewhite dwarf (Patterson et al. 1998). More generally, and morecritically, we prove that the T Pyx light is not coming from anaccretion disk (see Sections 5.2.2 and 5.2.3). All this is to saythat any model with a simple accretion disk is invalid and itsapplication will give a wrong distance. In all, the method ofcomparing the ultraviolet flux to models results in uncertaintiesthat are too large to produce any useful distance estimate.

T Pyx has a nice shell, so might it be possible to get adistance from this? In particular, the expansion parallax distancemethod is the traditional means to get nova distances thoughtto be good. But this method is beset by several factor-of-twouncertainties that combine to make for poor real error bars,with these problems always being ignored. (1) It is unclearwhat angular radius to assign to the shell in any image, partlybecause all nova shells are out of round, often significantly,partly because many novae (including T Pyx) have multipleshells, and partly because it is unknown as to which isophotallevel (e.g., the outermost, the half-peak, the peak) should beused. (2) It is unclear as to what velocity from a line profile to

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use, with possibilities including the half-width at half-maximum(HWHM) and the half-width at zero intensity (HWZI) anddiffering by a factor of two. (3) Nova shells are generally far fromspherical in shape, often being bipolar, such that the expansionvelocity along the line of sight can easily be up to a factor of twodifferent from a tangential expansion velocity. These problemsshow that expansion parallaxes have a real uncertainty of atleast a factor of two in distance. For the case of T Pyx, wecannot use the expansion parallax for two reasons. First, theexplosion is likely bipolar (Shore et al. 2013; Chesneau et al.2011), so the relation between the transverse expansion velocityand the radial velocity can only be speculative. Second, theobserved shell actually comes from a slow-velocity classicalnova eruption around the year 1866 (Schaefer et al. 2010a), sowe have no spectrum to get any radial velocity. (The later RNeruptions have completely different expansion velocities, butwe see no shell from them.) So expansion parallax cannot yieldany distance, no matter how crude, to T Pyx.

T Pyx has a pre-existing shell, so in principle, we can watchthe light echo of an eruption sweeping through the clumps inthe shell so as to get a distance. The idea would be to measurethe polarization of the echo light as a function of angular radius,where the peak in the polarization fraction corresponds to lightthat is single scattered at a right angle. With the geometryof the nova, we can watch the ring of maximal polarizationexpand at the speed of light, with the angular size equalingthe light travel distance since peak light. Unfortunately, the2011 eruption of T Pyx had the ring of maximum polarizationsweeping through the shell around the time of solar conjunction,and no polarization data were taken.

The preceding seven paragraphs present a horrifying casewhere all prior distance estimates to T Pyx are either erroneous,not applicable, and/or uselessly poor. (Indeed, similar analysiscan be made for most nova distances, with the only reliabledistances being for the few novae with parallaxes from theHubble Space Telescope, novae in external galaxies, novae ingalactic clusters, and novae with blackbody distances to thecompanion star.) But all is not lost, as there are three constraintsthat can be confidently placed. (1) The first constraint is thatthe interstellar absorption is significant and comparable withthe total galactic extinction along the line of sight. With thiswe can be confident that T Pyx is at a distance such thatits line of sight passes through most of the ISM in the disk,and it must rise more than something like 150 pc above thegalactic plane. With suitably round numbers, we know that Dis greater than ∼1000 pc. (2) There is some real upper limit onthe peak absolute magnitude of novae, and this is far above theEddington limit. The M31 novae have the most luminous eventat absolute magnitude −10.7 (Kasliwal et al. 2011), and T Pyxwill confidently not exceed this. The theoretical models of Yaronet al. (2005) come with the same peak. For AV = 0.75 mag anda V-band peak of 6.4 mag, the upper limit on D is 37,000 pc,which is far outside our Galaxy. (3) T Pyx is in the Milky Wayand is very unlikely to be in the sparse outer regions. A distanceof 10,000 pc would place T Pyx at a height of 1700 pc above theplane and at a galactocentric distance of 14,000 pc. Somewherearound this position will have a sufficiently low star density thatno one would care to say that T Pyx is farther. And this is all wehave that we can have any confidence in. Thus, we conclude thatT Pyx has a highly uncertain distance with only poor constraints,such that 1000 pc � D � 10,000 pc.

At the time this paper was accepted, Sokoloski et al. (2013)report a new measure of the distance based on a light echo

observed to sweep through the T Pyx shell from 2011 July to2012 May with the Hubble Space Telescope. No polarizationmeasures were taken, but a distance can still be gotten if therelative radial position of the reflectors is known. Sokoloskiet al. make a good case that the geometry of the reflecting dustis largely that of a clumpy inclined ring, with the axis of tilt beingsuch that the echoes to the north and south would have the sameradial distance as the central binary. This model distributionplus the time delay then allows for a good distance measureof 4800 ± 500 pc. So finally, we have a reliable and accuratedistance to T Pyx.

4.5. Stellar Masses MWD and Mcomp

Theory has a very strong result that RNe must have the massof the white dwarf (MWD) close to the Chandrasekhar mass,roughly 1.2 M� < MWD < 1.4 M� (e.g., Starrfield et al. 1988;Prialnik & Kovetz 1995; Hachisu & Kato 2001). The reasonis that only high-mass white dwarfs will have a high enoughsurface gravity to compress a thin layer of accreted materialto the kindling temperature for hydrogen for an amount ofmaterial that can be accreted in under one century. Observationaltests from radial velocity curves of this theoretical claim haveonly turned up patchy agreement. Published MWD values havebeen >2.1 M� for T CrB (Kraft 1958), 0.29 ± 0.14 M� forU Sco (Johnston & Kulkarni 1992), 1.55 ± 0.24 M� forU Sco (Thoroughgood et al. 2001), 1.37 ± 0.13 M� for T CrB(Stanishev et al. 2004), and 0.7 ± 0.2 M� for T Pyx (Uthaset al. 2010). This dismal record for radial velocity analyses ofRNe is matched by the dismal record for getting MWD in allcataclysmic variables, with the reason being that the emissionlines come from the inner disk and do not perfectly follow thedynamics of the white dwarf.

Both Selvelli et al. (2008) and Schaefer et al. (2010a) haveused information on the recurrence timescale, the accretionrate, and the ejection velocity (for both the RN events andthe 1866 classical nova eruption) so as to try to narrow therange of 1.2 M� < MWD < 1.4 M�. The typical analysis isto interpolate within the grid of theoretical models in Yaronet al. (2005) so as to find the white dwarf mass and accretionrate that provides the closest match in recurrence timescale andejection mass. Schaefer et al. (2010a) came to the conclusionthat 1.25 M� < MWD < 1.30 M�. Selvelli et al. (2008) find theclosest agreement with the models has MWD ≈ 1.33 M�, withan uncertainty ∼0.03 M�. With no better information, we willhere take the middle ground that is acceptable to both studiesand conclude that MWD = 1.30 ± 0.05 M�.

Into this situation, Uthas et al. (2010) report a radial velocitycurve for T Pyx, with the results that the companion starmass, Mcomp, is 0.14 ± 0.03 M�, MWD is 0.7 ± 0.2 M�, themass ratio, q = Mcomp/MWD, is 0.20 ± 0.03, and the orbitalinclination is 10◦ ± 2◦. Here, the given white dwarf mass isin large contradiction to the very strong theory. Fortunately, itturns out that their given MWD can easily be changed, as theyallow in their paper. Without accurate inclination measures (say,from an eclipse or ellipsoidal effects), we need to only slightlychange the inclination and MWD would become a lot larger.What Uthas et al. are actually measuring are the amplitude of theradial velocity curve for the emission lines plus the wavelengthseparation between peaks in the emission line profiles. Fromthese two measures, they use a non-standard (but reasonable)technique to derive the mass ratio. Their main line of argumentis then to claim that the companion should follow a slightlyinflated main-sequence mass-radius relation (with the radius

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known from the Roche lobe geometry) so as to come up witha companion mass of 0.14 ± 0.03 M�, and then use the massratio to get MWD. But it is a large leap of faith to expect thatsuch a weird system as T Pyx will have a normal main-sequencecompanion. So we are left with a dilemma, either that the whitedwarf is half of the Chandrasekhar mass in contradiction tostrong theory or that the companion star is not a normal main-sequence star. Faced with this dilemma, we have to choose thatthe companion star is not a normal main-sequence star. Withthis, we have a near-Chandrasekhar-mass white dwarf, and withtheir mass ratio, we get Mcomp = 0.26 ± 0.04 M�.

Another possibility is that the mass ratio is the problem,specifically with the emission line profile not coming from adisk. In Section 5.2.2, we prove that the continuum light canat most have only a small fraction coming from an accretiondisk, so it is reasonable to think that the emission lines do notcome from a disk. (There are other possible explanations forthe double-peaked profile, including from the pair of accretioncolumns.) In the absence of a disk, the argument to go fromthe separation of the peaks to the mass ratio is broken. Then,we still have MWD = 1.30 ± 0.05 M�, but now we can acceptthe argument based on the orbital period (plus some assumedsequence of stellar structures) that Mcomp = 0.14 ± 0.03 M�.With this, we would have q = 0.108 ± 0.023.

We do not have a clear answer as to whether the companionhas a mass of 0.14 or 0.26 M�. The problem is that there arethree claims (the white dwarf is near the Chandrasekhar mass,the mass ratio comes from the peak separation in the line profiles,and the companion satisfies some mass-radius relation) that arecontradictory. Various researchers will pick out their favoriteclaim and derive their conclusion. For example, we know thatRNe require near-Chandrasekhar-mass white dwarfs. With allthree claims denied in turn, the companion is still either 0.14 or0.26 M�.

4.6. Accretion Rate M

We can use our data and analysis to address the question of thequiescent accretion rate. The value of M is key for understandinga variety of peculiarities in the T Pyx system. The most criticalneed for M is so that we can calculate how much mass hasarrived on the white dwarf from 1967 to 2011, for a comparisonwith the mass ejected by the 2011 eruption. This comparisonwill then decide whether the white dwarf has gained or lostmass over the entire eruption cycle, and hence whether it canbe a Type Ia supernova progenitor. Here are five methods fordetermining M , some new, some old, and some failing.

Method 1. The orbital period of T Pyx has been changingclosely linearly with time (see the good parabolic shape of theO − C curve in Figure 8). This large and steady period changecan only arise from the mass transfer in the system. The observedincrease in the orbital period is expected when a low-mass donorstar transfers matter to the high-mass white dwarf. The periodchange will be proportional to M , so we can use the observedperiod change to deduce the accretion rate. The default scenariois conservative mass transfer, in which all the mass leaving thedonor star arrives onto the white dwarf, and this material carrieswith it the specific angular momentum of the donor star. Forconservative mass transfer,

M = (Mcomp/[3 − 3q])/(P/P ). (5)

With Mcomp = 0.26 M�, we have M = 3.5 × 10−7 M� yr−1.With Mcomp = 0.14 M�, we have M = 1.7 × 10−7 M� yr−1.

The conservative case is where all matter leaving the companionfalls onto the white dwarf, with no angular momentum lostto the system. This case has to be fairly close (but certainlynot exact) for T Pyx. Depending on the details of how theaccretion process moves mass and angular momentum, theaccretion rate can vary somewhat from this derived M . Withinthis moderate uncertainty, this basic result is robust, dependingonly on timing measures and on simple dynamics of the binary.We will express the allowed range of the accretion rate asM = (1.7–3.5) × 10−7 M� yr−1.

There are alternative possibilities to produce a good parabolain the O − C curve, so we should consider these cases. Here arethe three other processes, all of which fail to explain T Pyx. (1)The O − C curves for cataclysmic variables occasionally showa sinusoidal oscillation with periods of years, likely associatedwith activity cycles on the companion star (Bianchini 1990). Ifonly a portion of the sine wave O − C curve is observed, thecurve might appear as a parabola, for example, when Robinsonet al. (1995) found that the Z Cam had a nearly parabolicO − C curve from 1972 to 1992, yet when the O − C curvewas extended to 2002 the shape looked more like one periodof a sine wave (Baptista et al. 2002). We have fit our O − Ccurve for T Pyx to a sine wave, and find acceptable fits onlyfor periods longer than 80 yr and half-amplitudes larger than0.2 days. This is greatly different than all solar cycle variationsin cataclysmic variables (Baptista et al. 2003) because theseall have periods less than 30 yr (with the median value near10 yr) and O − C half amplitudes of less than 3 minutes (withthe median value 1.2 minutes). The cataclysmic variables belowthe period gap have a maximum half-amplitude of less than oneminute (0.0007 day), and so this effect is a factor of >300×too small to account for T Pyx. (2) The O − C curve will alsochange as a sine wave for the case of T Pyx being in far orbitaround a third star, with the visible parabola merely being a partof the sine wave. For such a possibility, the orbital period of thethird body would have to be >80 yr and the orbital radius ofthe nova binary system would be >0.2 lt-day, or >35 AU. FromKepler’s Law, any third body must currently be >9 M� and itsprogenitor would have exploded as a supernova before the whitedwarf could have formed in the system. Even if such a triplecould have remained bound, we do not think that anyone will tryto claim that T Pyx has a massive black hole as a very wide thirdbody. (3) Cataclysmic variables above the period gap suffer asteady period change due to magnetic breaking (Patterson 1984),and this will lead to a parabolic O − C curve. Importantly, thismagnetic breaking turns off below the period gap, so this processwould not be applicable to T Pyx. Nevertheless, suspicionshave been raised of some small residual magnetic breakingbelow the gap, so we should still consider this possibility.This possibility is strongly refuted by either of two means.First, the magnetic breaking mechanism as applied to long-period cataclysmic variables always produces O − C changeswith P/P > 5,000,000 yr (Patterson 1984), so a small residualeffect for short-period systems must have P/P 108 yr or so,whereas T Pyx has P/P = 313,000 yr. Second, the magneticbreaking robs the system of angular momentum, so the orbitalperiod would be decreasing, whereas T Pyx has its orbital periodincreasing. In summary, the three alternative possibilities allcompletely fail.

The parabolic O − C curve is caused by the normal masstransfer. We know this because all alternatives fail completely.We also know this because the behavior seen (period increasing,P/P = 313,000 yr, and a good parabola shape) is just that

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expected from strong theory (see the discussion for the nexttwo methods). We also know that the parabolic O − C curvecomes from mass transfer because the small deviation from2008 to 2011 corresponds with the drop in the accretion rate.The measured period change comes from timing observations(always good for high reliability), and the derivation of M isfrom a purely dynamical method (where the physics is simpleand well known, being independent of distance, extinction,and anyone’s theory). The derived accretion rate has someuncertainty associated with the not perfectly known companionmass, and some additional comparable uncertainty will arisefrom the mass transfer not being perfectly conservative. Butwithin these error bars, our derived M is robust and by farthe best measure for T Pyx. In summary, we find M =(1.7–3.5) × 10−7 M� yr−1.

Method 2. We can set a very useful upper limit on the accretionrate onto the white dwarf, because if M is too high, thenthe hydrogen on the surface of the white dwarf will undergosteady nuclear burning. We see a normal nova eruption, so weknow that there is no steady hydrogen burning, so the accretionrate must be below some limit. The limit for steady hydrogenburning is (2–3)×10−7 M� yr−1 for a near-Chandrasekhar-masswhite dwarf (Nomoto et al. 2007; Shen & Bildsten 2007). Asubstantial problem for this method is that the quoted steadyhydrogen burning limit is for a white dwarf that is in thermalequilibrium with the given accretion rate, whereas T Pyx hasa complex history with rates apparently going from 10−11 to10−6 M� yr−1 in the recent past. With a cooler white dwarf thanexpected for the current accretion rate, the threshold for steadyhydrogen burning should be raised by some unknown amount.With this, the accretion rate limit from this method might haveto be somewhat raised.

Method 3. The 2011 eruption ties the inter-eruption intervalto 44 yr. With the approximate mass of the white dwarf, we canestimate the required trigger mass on the white dwarf, so theaccretion rate will be just the trigger mass divided by the inter-eruption interval. A substantial problem is that the accumulatedmass required for the trigger has a strong dependence on theadopted conditions. It varies by over an order of magnitudeover the acceptable range of white dwarf masses (Yaron et al.2005), it varies by a factor of four over a plausible range ofmetallicities, and it varies by a factor of five over a plausiblerange of 3He mass fraction (Shen & Bildsten 2007). In addition,the trigger mass varies by a factor of two with a change inM by a factor of 10 over the plausible range. Nevertheless,for the estimated conditions, the T Pyx white dwarf must haveaccumulated 10−6.0±0.3 M� (Yaron et al. 2005), 10−6.3±0.2 M�(Nomoto et al. 2007), or 10−5.5±0.3 M� (Shen & Bildsten 2007)to be able to erupt. This range of theoretical models can berepresented as 10−5.9±0.6 M�. With the 44 yr interval, we get thatthe average rate of deposition of mass onto the white dwarf is inthe range 10−7.5±0.6 M� yr−1. Importantly, this highly uncertainestimate is based on specific theoretical models, so we must becareful to not use this estimated M to evaluate future models.

Method 4. We can compare our GALEX spectrum with theaccretion disk models of Wade & Hubeny (1998) and pick outthe accretion rate that best reproduces the model. The availablemodels have the two highest mass white dwarfs of 1.21 and1.03 M�, while the highest mass model has the two highestaccretion rates of 10−8.5 and 10−9.0 M� yr−1, so we have toextrapolate far outside the range of available models to get tothe conditions likely to be relevant for T Pyx. The modelshave a wavelength range of 850–2000 Å, while the binned

GALEX spectrum has a range of 1456–2828 Å, so the overlap forcomparison is fairly small from 1456 to 2000 Å. Over this range,the observed spectrum is featureless (with the two emissionlines not being relevant for the disk models) so all we can getis a spectral slope. Over this same range, for low-inclinationsystems, the models all are largely featureless (certainly forfeatures large enough to be visible in the GALEX spectrum) andhave essentially the same spectral slope. In this situation, thereis no way for the GALEX spectrum to distinguish the disk modelwith the best accretion rate. This same analysis and conclusionalso hold for the IUE spectrum. In all, this hopeful method doesnot work in practice. And in principle, the likelihood that T Pyxis a magnetic system (with no disk or a truncated disk) and/orhas non-disk light sources dominating in the ultraviolet impliesthat the disk models are not applicable anyway.

Method 5. The accretion luminosity dominates the light fromT Pyx, and this is proportional to the accretion rate, so inprinciple we can take the observed flux from the system andderive M . Typically, this will involve making some estimate ofthe unobserved flux in the FUV, then performing an integralover the SED (as in Section 3.4), and finally multiplyingby 4πD2 for some assumed distance. Selvelli et al. (2003)report such a calculation based on their IUE spectrum, derivingaccretion rates of (2.2–4.6) × 10−8 M� yr−1. Selvelli et al.(2008) used a similar method to determine the accretion rates as(1.1 ± 0.3) × 10−8 M� yr−1, while they also report on variousrelations between M and the V-band absolute magnitude tocome up with similar accretion rates. Unfortunately, all suchcalculations have real uncertainties of over four orders ofmagnitude in M . One reason is that the extrapolation to theunobserved ultraviolet leads to a range of a factor of 104×depending on where the break in the SED is in the range from2 × 1015 Hz to 2 × 1017 Hz (see Section 3.4). And there isanother independent factor of 100× uncertainty simply dueto the distance to T Pyx being uncertain by a factor of 10.So, for acceptable input, we could get a luminosity as low as2×1034 erg s−1 and as high as 2×1040 erg s−1. Any M derived inthis way will have a real uncertainty of six orders of magnitude,and that makes this method useless.

Method 6. T Pyx is at best only a faint X-ray source, andthis can be used to place limits on the accretion rate. Selvelliet al. (2008) report that their observations with XMM-Newtonfor 22.1 ks in 2006 show a faint source visible over the0.2–8 keV energy range. The signal was too faint to allow fora reliable fit, but they demonstrate that a 200,000 K blackbodywith 1037 erg s−1 source at 3.5 kpc is rejected because thespectrum below 0.5 keV would be roughly 300 times brighterthan observed. For the soft X-ray component, this would implyan X-ray luminosity of roughly 1034.5 erg s−1. Balman (2010)used the same XMM-Newton data set from 2006 (except shereports on effective exposure times of 30.6, 39.7, and 38.8 ksfor the three detectors). She found that an X-ray image showsan extended profile, with the unstated fraction of light abovea single point source certainly coming from the shell aroundT Pyx. A spectral fit gives a luminosity of 6.0 × 1032 erg s−1

at a distance of 3.5 kpc over 0.3–10.0 keV. If this X-rayluminosity is taken to be the light from the accretion onto thewhite dwarf ((1/2)GMWDM/RWD), then the accretion rate is8 × 10−11 M� yr−1. If the X-ray luminosity is taken to be fromsteady hydrogen burning of the freshly accreted material onthe surface of the white dwarf (0.007XMc2 if all the hydrogenis burned, where X is the usual hydrogen fraction), then theaccretion rate is around 2 × 10−12 M� yr−1. But this statement

15


is made with the assumption that all of the accretion luminosityis coming out in the X-ray, whereas this is clearly not true.Most of the accretion energy is coming out in the ultraviolet orthe FUV. So this method really is only producing a limit, withM 2 × 10−12 M� yr−1.

We have given a full discussion for all six methods becausethey have already appeared in the literature. We have foundmethods 4–6 as being useless or inapplicable. Methods 2 and 3are based on theoretical models, which have to be right, althoughthe real uncertainties are larger than we would hope for. Onlymethod 1 returns a reliable and robust measure of M , and eventhis has an uncertainty of roughly a factor of two. So, finally,we conclude that M = (1.7–3.5) × 10−7 M� yr−1.

5. ANALYSIS

In the following subsections, we address a variety of questionsand topics related to a model for T Pyx.

5.1. The Rise to Peak

T Pyx has a well-defined light curve with essentially continu-ous coverage and BVI colors starting at the time of the discovery(see Figures 2 and 3 and Table 2). The nova rose from V = 13.0(discovery) to V = 10.0 in a time of 8 hr, for an average rate of9.0 mag day−1. This fast rise slows within the first day, rising toV = 7.9, 2.0 days after discovery. Then, T Pyx fades by a thirdof a magnitude over the next day, followed by a very slow riseback to V = 7.9 over the next five days. Soon after this, the riseincreases to a minor peak of V = 6.7 on JD 2,455,685 (19 daysafter discovery) and then up to the highest peak of V = 6.4 onJD 2,455,693 (27 days after discovery).

The initial rise of a nova can be simply modeled by auniformly expanding shell. The observed flux from the systemwill equal

F = F0 + (vτ/D)2Fν(T ). (6)

Here, F0 is some constant flux from the disk and companionstar, v is the expansion velocity of the outer edge of the shell(corresponding to the effective photosphere), τ is the time ofexpansion, D is the distance to the T Pyx system, T is theeffective temperature of the photosphere, and Fν(T ) is the fluxfrom a unit area of the photosphere. For the case of bolometricflux, Fν(T ) = σT 4 for a blackbody photosphere. Equation (5)can be converted into AB magnitudes as

m = −48.60 − 2.5 log10[F0 + C(t − t0)], (7)

withC = (v/D)2Fν(T ). (8)

Here, τ is expressed in the more practical form of the time (t)minus some fitted zero time for the start of the rise (t0). C is acombination of values that presumably remain nearly constantthroughout the fast initial rise. Before t0, the magnitude shouldbe a constant corresponding to the flux F0 alone. We now havea simple model with three unknowns (F0, C, and t0) that shoulddescribe the initial rise.

We have fit our T Pyx light curve from the time of discoveryuntil JD 2,455,668.0 with a chi-square analysis. This includes2151 magnitudes. To get a reduced chi-square of unity, wehad to adopt a systematic uncertainty of 0.25 mag, whichis a measure of systematic errors in the measurements plusdeviations from the model. Our best-fit model is shown inFigure 4. The constant flux corresponds to a magnitude of 14.63,

t0 = 2,455,665.967, and C = 1.56×10−23 erg s−2 cm−2 Hz−1.This model fits the initial fast rise well, and gives someconfidence that the simple model is close to the real case.The sudden deviation two days after the initial rise would thenindicate some significant departure from the simple model inthe sense that the expanding shell appears fainter than expected.We suggest that this deviation is caused by the outer portion ofthe ejecta expanding to become optically thin enough so that thephotosphere can shrink significantly in size.

We had originally hoped that this rise could be used todetermine an accurate distance to T Pyx by the Baade–Wesselinkmethod. Our light curve fit gives a value for C, the spectralline widths will give V, and we can use theory or colorsto estimate Fν(T ). Then the distance can be directly derivedfrom Equation (8). Unfortunately, the inputs are known onlypoorly. For the velocity, we do not know what position in aline profile (e.g., HWZI or HWHM) corresponds to the velocityat the photosphere. For Fν(T ), the blackbody approximation isexpected to be good for the photosphere, but the early colors areinconsistent with a blackbody. As such, the temperature cannotbe determined with any useful confidence from observation ofthe colors. Also, no early spectrum exists, so we cannot get thetemperature from spectral lines. From theory, the temperaturevaries by a factor of ∼3, with a strong dependence on the whitedwarf mass and other model details. With no way to get thetemperature to any useful accuracy, we cannot get Fν(T ) even ifwe know the emission mechanism. Nor can we turn this aroundto derive Fν(T ) from C, D, and V because again the uncertaintiesof the inputs would be too large to be useful. In all, even thoughC is fairly well measured, we cannot get useful information onany of D, V, or Fν(T ).

5.2. The Nature of the Emission from Near the White Dwarf

The secondary star is a small red dwarf with a negligibleluminosity, and the white dwarf itself is so small that it cannotprovide any significant emission. So all the observed flux fromT Pyx can only come from the accretion luminosity by somemeans. This emission can only come from near the white dwarf,perhaps as an accretion disk, an accretion stream, or some morecomplicated structure. This subsection gives our analysis as tothe nature of this emission.

The mass falling onto the white dwarf has a lot of energythat will be radiated at some wavelength. We have a fairly goodidea of the accretion rate and the white dwarf properties, sowe can make a good estimate of the accretion luminosity asGMWDM/RWD. With (1.7–3.5)×10−7 M� yr−1 falling onto thewhite dwarf, the luminosity would be (2.7–5.5) × 1036 erg s−1.This luminosity is independent of the large uncertainty inthe distance, rather it comes from knowing the accretion ratedynamically from the steady P . This accretion luminosity couldcome out with any of a variety of possible SEDs, and with someunknown fraction in the unobserved FUV.

5.2.1. Not a Single Blackbody

Selvelli et al. (2003) and Gilmozzi & Selvelli (2007) pointout how the IUE spectrum can be well fit with a singleblackbody spectrum of 34,000 K. This temperature explainswhy essentially no X-ray flux is visible with XMM. However, asthey point out, this blackbody cannot account for the flux outsidethe ultraviolet regime, because the optical and infrared flux isfar above this single blackbody. This can be seen in Figure 7,where the optical and infrared fall off far from the fν ∝ ν2 of

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the Rayleigh–Jeans tail. No single blackbody can account forany significant fraction of the T Pyx SED over any but a narrowfrequency range.

Patterson et al. (1998) and Knigge et al. (2000) postulatedan emission source near the white dwarf that emits a largeluminosity of FUV and supersoft X-ray photons so as to irradiatethe companion star and drive the required high accretion rate inthe T Pyx system. They took this source to be something like ablackbody with a temperature of 240,000 K and a luminosity of1037 erg s−1. Selvelli et al. (2008) have confidently proven thatany such source would have been seen with the XMM-NewtonX-ray images in 2006.

5.2.2. Not an Accretion Disk

Based on the strong precedent of other novae and cataclysmicvariables, the common expectation would be to have the fluxcoming from some sort of an accretion disk with reprocessing.Indeed, Gilmozzi & Selvelli (2007) find that the ultraviolet SEDis well fit by fν ∝ ν1/3, and this is exactly the slope expected forthe continuum from disks. (But this is not a strong argument dueto their ultraviolet spectrum also fitting a blackbody. Further,Frank et al. (2002) point out that for the case of T Pyx witha hot outer edge, the ν1/3 regime “may be quite short and thespectrum not very different from a blackbody” and “disks aroundwhite dwarfs with Rout ∼ 102Rin [as for T Pyx] do not possessan obvious power-law character.” We know that the T Pyx ν1

spectrum has to turn over, so there must be some region aroundthe ultraviolet where the SED slope will be ∼ν1/3 whether ornot there is any disk.) An additional argument for a disk isthat Gilmozzi & Selvelli (2007) found absorption lines in theultraviolet that are characteristic of disks. (This last point isonly weak evidence for a disk, because it really only says that acool gas of ordinary composition is just outside the continuumemission region, and there are many structures that can yieldthe same absorption lines. In addition, their detailed comparisonwith the accretion disk model grids of Wade & Hubeny (1998)showed no satisfactory match.) So the default idea is that thelight from T Pyx is dominated by an accretion disk.

But this default idea cannot be correct because the SED isgreatly different from that of any accretion disk. The generalshape of accretion disk continua is a fν ∝ ν1/3 power law overa region (perhaps quite narrow) corresponding to temperatureswithin the disk, with a Rayleigh–Jeans slope (fν ∝ ν1/3) tolower frequencies (see Figure 5.2 of Frank et al. 2002 and Wade& Hubeny 1998). The ultraviolet portion of the T Pyx SED isconsistent with a ν1/3 power law, but the entire UBVRIJHK partof the SED has fν ∝ ν1. The UBVRIJHK region is much toolarge a range to be the rollover from ν1/3 to ν2, and we see zerocurvature in Figure 7. With this, we know that the UBVRIJHKemission is not from an accretion disk.

For any disk around T Pyx, even the outer edge wouldbe very hot. Detailed calculation for an α-disk and M =3 × 10−7 M� yr−1 give a temperature of the outer edge (takento be at 90% of the Roche lobe radius) to be Touter = 20,000 K,or 42,000 K if reprocessing of the accretion luminosity fromthe central source is included. The turnover to a Rayleigh–Jeansslope in the SED occurs at a frequency of kTouter/h, with thisbeing either 2.6 × 1015 Hz or 5.5 × 1015 Hz. (This temperatureof the outer edge of any disk in the T Pyx system is substantiallyhotter than is common for cataclysmic variables, simply becauseT Pyx has such a high accretion rate and because the shortorbital period requires a very small disk.) Thus, the entireobserved portion of the ultraviolet should already be in the

Rayleigh–Jeans region. This means that the observed fν ∝ ν1/3

from IUE has no causal connection with any accretion disk.

5.2.3. Not a Superposition of Blackbodies

Is it possible to superpose a number of blackbody spectraso as to produce the observed fν ∝ ν SED? That is, just as anaccretion disk SED is a superposition of blackbody spectra fromdifferent temperatures within the accretion disk so as to producea power-law SED, perhaps we can find some other distribution oftemperatures over the emitting surface area so that the sum of allthe thermal spectra will produce the observed SED. The surfacearea of the photosphere of the thermal emission must be entirelycontained within the Roche lobe of the white dwarf, because thewhite dwarf motion is visible in the radial velocity curve (nothidden inside some enveloping photosphere) and because anylarger structure would be rapidly destroyed by churning fromthe companion star. The radius of the white dwarf Roche lobeis close to 0.5 R�.

Can the near-infrared flux come from any superposition ofblackbody spectra? The most efficient case (while avoidinga Rayleigh–Jeans SED) is to have enough surface area witha temperature such that the peak of fν , at a frequency ofνmax, is in the near-infrared. For the K-band flux, we haveνmax = 1.36 × 1014 Hz and fν = 0.0012 Jy. The νmax valuecorresponds to a temperature of 2300 K, for which the blackbodyequation gives a flux of Fν = 7.42 × 10−6 erg cm−2 s−1 Hz−1.These quantities are connected as 4πD2fν = 4πR2

sphereFν ,where Rsphere is the radius of a sphere that has the requiredsurface area. The distance D is unknown with only poor limitsof 1000 pc � D � 10,000 pc, even though T Pyx has a“traditional” value of 3500 pc or so. For the minimal distance of1000 pc, the K-band flux can only be produced by a region whosesurface area equals that of a sphere 1.80 R� in size. The inclusionof any hotter surface area would add flux to the K band onlywith this hotter blackbody imposing a Rayleigh–Jeans slope onthe SED, and such is not observed. This required surface area,so as to produce the K-band flux, is greatly larger than that ofthe Roche lobe, so there is no possibility that any significantamount of the flux comes from any superposition of thermalspectra. This same conclusion holds true for the R-, I-, J-, H-,and K-band fluxes. For the “traditional” distance of 3500 pc, theflux with wavelengths longer than 2000 Å cannot come froma superposition of blackbodies. Thus, the red and near-infraredflux from T Pyx certainly cannot come from any superpositionof thermal spectra, and hence there must be a bright nonthermalsource in the system.

We can place a distance-independent limit on the presenceof any superposition of blackbodies by realizing that theaccretion luminosity (with M = (1.7–3.5) × 10−7 M� yr−1)must come out in the SED. This accretion rate gives a totalaccretion luminosity of (2.7–5.5) × 1036 erg s−1. The mostefficient way to get this luminosity out with a superpositionof blackbodies is to put all of the surface area to the highesttemperature. (This would immediately violate the observedSED shape, but such a calculation is good to show an extremelimit.) Gilmozzi & Selvelli (2007) fit the ultraviolet SED to atemperature of 34,000 K, and the temperature cannot be muchhigher without violating the X-ray limit. A 34,000 K blackbodyemitting 5 × 1036 erg s−1 has a surface area corresponding toa sphere with a radius 1.0 R�. That is, even pushed past whatis possible, the accretion luminosity requires a surface area thatis much larger than that of the white dwarf Roche lobe. Withmost of the luminosity coming out in the ultraviolet, we see

17


that the accretion luminosity is too large to come out as anysuperposition of blackbodies.

5.2.4. Nonthermal Emission

With all possibilities of thermal emission strongly rejected,we can only conclude that the T Pyx SED is primarily fromnonthermal emission. Nonthermal emission, say, from free-freeor cyclotron processes, can easily produce a high luminositywithin the Roche lobe. Nonthermal emission can producefν ∝ ν over a wide range of frequency. So we know that thelight from T Pyx is nonthermal.

Nonthermal emission must come from a region that is notoptically thick, because if it were optically thick, then thephotons would rapidly thermalize, form a photosphere, and besent into space with a blackbody spectrum. This means thatwe are seeing a volume of emission, not a surface of emission.Likely, any such volume will have increased density near thewhite dwarf, and some central core might have high enoughoptical thickness so as to appear as a blackbody, with this thermallight being seen simultaneously with the nonthermal emissionfrom the surrounding volume. It is unclear what fraction of thisvolume is optically thick.

Some fraction of the light in the SED can still come from ablackbody or a disk, but any such contribution must be smallso as to beat the limits above. The most likely spectral regimefor such an exception would be in the ultraviolet, where evena body with a small surface area can provide much luminosity.Thus, a disk can still be hiding in the SED, with its light beingswamped by nonthermal light, but the above limits mean thatonly a small fraction of the SED and of the accretion luminositycan come out as disk light.

At least three nonthermal emission mechanisms can providepower-law SEDs with roughly the correct slope. The first is free-free emission. The second is cyclotron emission by thermalelectrons in a hot plasma, presumably near the white dwarfand fed by the accretion. The third is from Compton scatteringof low-energy seed photons by hot electrons in a surroundingplasma.

Nonthermal processes lack the simplicity of blackbody emis-sion. This means that any calculations (e.g., energy budgets,SED shape) depend on the details of the model, with thesealways being very complex, dependent on various unknown pa-rameters, and largely unsolved. As such, this paper cannot makeany useful predictions for the nonthermal emissions, and it iseasy to imagine a variety of scenarios.

5.3. The White Dwarf Magnetic Field

Williams (1989) claims that the He ii λ4686 line is a signatureof the magnetic field strength on the white dwarf, with theemission line coming from accretion columns attached to themagnetic poles. He also points out that T Pyx and U Scoare the two cataclysmic variables with the strongest He iistrength (relative to Hβ). Putting these two points together, wesee that T Pyx would have a very strong magnetic field andaccretion columns from its white dwarf.

Oksanen & Schaefer (2011) and Schaefer & Collazzi (2010)have pointed to the T Pyx system as having a highly magnetizedwhite dwarf that channels the accretion flow onto a polar cap.The first reason is that T Pyx shows photometric oscillationsthat vary in time with a period near to, but not equal to,that of the orbital period. This is the hallmark of magneticsystems, and this one property has been used to identify many

systems as magnetic. For examples amongst novae, see RWUMi (Tamburini et al. 2007), V4633 Sgr (Lipkin & Leibowitz2008), V4745 Sgr (Dobrotka et al. 2006), V4743 Sgr (Kanget al. 2006), V697 Sco (Warner & Woudt 2002), and V1495Aql (Retter et al. 1998). (But this is not a forcing argument,as there are other means to make non-orbital photometricperiodicities, for example, superhumps.) The second reasonis that T Pyx is a quintessential example of the V1500 Cygstars, wherein post-eruption novae are greatly brighter than pre-eruption novae while fading slowly for the next century or so.All eight known cases have relatively short orbital periods, long-lasting supersoft emission, and highly magnetic white dwarfs;with all these uncommon shared properties likely forming asuite of requirements to get the bright post-eruption. Schaefer &Collazzi (2010) present a mechanism by which nuclear burningis sustained on the surface of the white dwarf (hence irradiatingthe companion star and driving the high accretion rate) by themagnetic field channeling the accretion onto a polar cap wherethe local accretion rate is high enough to sustain the nuclearburning. (But this is not a strong reason, because there areother possible mechanisms that can produce the post-eruptionirradiation-induced high-mass transfer that characterizes theV1500 Cyg behavior.) The third reason is that T Pyx has oneof the highest fluxes in the He ii λ4686 line, and this has beentaken as pointing to magnetic accretion as the only way to geta high excitation luminosity (Williams 1989; Diaz & Steiner1994). (But this is not a strong reason because little studyor explanation of the connection, bright helium emission tohigh magnetic fields, has been made, much less proven to beunique.) The fourth reason is that we have already proven thatthe emission from T Pyx cannot be from a disk of any type(see Section 5.2.4), and the only way to get the material to thesurface of the white dwarf without a disk is through some sortof a magnetic accretion column. (This is not a strong argumentbecause perhaps the disk light is being swamped by nonthermallight from some other source in the system.)

The existence of a strong magnetic field would be helpfulfor explaining how the optical light curve of T Pyx can havea 0.15 mag amplitude modulation where the minimum coversabout half an orbit, despite having an inclination of roughly10◦. With T Pyx being seen essentially face-on, there canbe no eclipses, while any illumination or irradiation effectsmust be small (cos[10◦] = 0.985). So there must be someasymmetry in the system that is tilted with respect to the orbitalplane. The obvious possibility, and we cannot think of anyother possibility, is to invoke a tilted magnetic field in thewhite dwarf. Nonthermal cyclotron emission can easily havea mild beaming pattern that changes by 15% for a 20◦ shiftin viewing angle. For example, maybe the accretion column ispointing far away from the normal to the orbital plane, or themagnetic field suffuses some cloud near the white dwarf withcyclotron emission preferential in some direction. For this ideato work, the rotation period of the white dwarf must be exactlytied to the orbital period, because the photometric period istied to the radial velocity period. Again, any real calculations(e.g., beaming patterns and degree of polarization) are modeldependent, complex, and vary with unknown system parameters.

The reality of the high magnetic field for T Pyx is notproven, and indeed there are troubles with the basic idea. First,a plausible standard of proof is to measure significant circularpolarization that changes with the orbital period, and T Pyx hasno measured polarization. (But this is not evidence, becauseno reports have been published yet with circular polarization

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measures around an orbit. The linear polarization measuresreported by Eggen et al. (1967) are for the time around peakbrightness when the outer shell was shrouding the inner binary.)Second, Patterson et al. (2012) have argued that the original non-orbital periodicity reported by Patterson et al. (1998) has neverbeen seen to re-appear, so “without confirmation, its evidentiaryvalue must be reckoned as weak.” (This argument is weakbecause they are not denying their own data, and there is norequirement for the signal to be frequently present.) Third, theyfurther argue that the V1500 Cyg stars with orbital periods belowthe period gap (GQ Mus, V1974 Cyg, CP Pup, and RW UMi)have their non-orbital photometric periodicities of low phasestability, hence pointing to a superhump phenomenon insteadof a magnetic white dwarf. (This is a weak argument becauseif the earlier non-orbital signals had been from superhumping,then superhumps should have been repeatedly seen and clearlyidentified in their long photometric records. This argument isalso weak because these same stars have been identified asbeing highly magnetic systems on independent grounds; seeDiaz & Steiner 1994; Diaz & Steiner 1991; Chochol 1999.)Pointedly, Patterson et al. (2012) agree with Schaefer & Collazzi(2010) that the V1500 Cyg stars with orbital periods abovethe period gap have high magnetic fields, so this discussionis just about the short-period systems like T Pyx. Fourth, thespectra of T Pyx show double-peaked emission lines, with thisshape characteristic of disks with a bright inner region. (But thisargument is weak because other structures, like pairs of accretioncolumns, can also produce the observed line profiles.) Takingall the arguments pro-and-con together, we see no decisive caseeither way. Until resolved, we judge that the existence of ahigh magnetic field on the T Pyx white dwarf is reasonable yetundecided.

Fortunately, we already have excellent polarization data forT Pyx, taken in 2013 February with FORS on UT1 at the VeryLarge Telescope on Cerro Paranal in Chile. This observing runhas S. Katajainen as P.I. plus B. E. Schaefer and A. Oksanenas Co.I.s. The answer to whether T Pyx is polarized, and hencewhether its white dwarf is highly magnetized, is not yet known,but an answer will come soon.

5.4. The 44 yr Inter-eruption Interval

Schaefer (2005) and Schaefer et al. (2010a) made a quanti-tative physics-based prediction that the date of the next T Pyxeruption would be in the year 2052 or much later. The first keyidea is that the accretion rate can be measured from the blueflux, so a B-band light curve between eruptions can be usedto determine the average accretion rate. This is certainly truefor all emission mechanisms and structures. For any particularsystem over a relatively restricted range of M , we can representor approximate the relation as a power law. The general relationwill be

M = M16FEB = M1610−0.4(B−16)E, (9)

where M is the accretion rate, M16 is the accretion rate whenT Pyx has a B magnitude of 16.0, B is the B-band magnitude, FBis the B-band flux in units of the flux when B = 16, and E is anexponent representing the physics of the light sources in T Pyx.The E value will vary from source to source, depending on thestructure of the accretion flow, the emission mechanism, andthe wavelength of observation. The use of the B = 16 level isarbitrary (any other magnitude level could have been adopted),but this choice is convenient for T Pyx. The second key idea isthat the total mass accreted between eruptions can be estimated

from the accretion rate averaged over the inter-eruption intervalof ΔT , with

ΔMacc = 〈M〉ΔT . (10)

Here, the accretion rate values from Equation (9) are averagedover all available blue magnitudes. The third key idea is thatthe trigger mass will be close to the mass accreted betweeneruptions:

Mtrigger = ΔMacc. (11)

The fourth key idea is that a given RN system should havea unique trigger mass that will be the same for all eruptions.The trigger mass can only be a function of the white dwarfmass, radius, magnetic field, and the composition of the accretedmaterial, and these vary from RN to RN, but not from eruptionto eruption. This can be expressed as

Mtrigger1 = Mtrigger2 = Mtrigger3 = · · · , (12)

where the numerical subscript indicates different eruptions.Putting these equations together, we see that 〈FE

B 〉ΔT shouldbe a constant from eruption to eruption. For intervals where theRN is bright on average, the accretion rate will be high, thetrigger mass will be accumulated relatively fast, and the inter-eruption interval will be relatively short. For intervals wherethe RN is faint on average, the time between eruptions willbe long. For both T Pyx and U Sco (Schaefer 2005; Schaeferet al. 2010b), the blue flux was seen to vary with the short inter-eruption intervals having a bright quiescence (and thus a highaccretion rate) and with the long inter-eruption intervals havinga faint quiescence (and thus a low accretion rate). We now candetermine the current inter-eruption time interval as

ΔT3 = ΔT2〈10−0.4(B−16)E〉2/〈10−0.4(B−16)E〉3. (13)

The subscripts on the brackets indicate an average over theindicated inter-eruption interval.

We now have a means of predicting the date of the nexteruption by simply measuring the historical light curve alongwith some assumption for the dependence of the quiescentbrightness on the accretion rate. By calibrating with previousinter-eruption intervals, the various constants of proportionalitycancel out, and we need never know the trigger mass, accretionrate, or the distance. Schaefer (2005) predicted that the nexteruption of U Sco would be in 2009.3 ± 1.0, and this formed thebasis for a large international collaboration of researchers at allwavelengths designed to discover and then intensively observethe eruption. With the discovery of the eruption in 2010.1 by B.Harris and S. Dvorak, this advance preparation allowed the USco eruption to become the all-time best-observed nova event.Schaefer (2005) also predicted that T Pyx would next erupt in2052 ± 3, while Schaefer et al. (2010a) extended this date tothe year 2225 or much later.

This prediction for T Pyx was refuted by the 2011 eruption.So the task now is to understand why the prediction went sofar wrong. The basic logic of the prediction is simple andeasy. The light curve is confident and the mass for the triggerthreshold really has to be constant. About the only place thatcould be wrong is the conversion from the B-band magnitude toan accretion rate. Schaefer (2005) calculated that the accretionrate is proportional to the square of the blue flux (i.e., E = 2.0),with this exponent being derived from a model for a simpleaccretion disk with the parameters for T Pyx. (For U Sco withits much larger disk, the exponent is 1.5.) So this exponent isthe likely cause of the wrong prediction.

19


The exponent E is merely a description of the accretion-fluxrelation for a relatively small range of M for a given star andwavelength. The E value varies substantially. We have alreadyseen how E = 2.0 and E = 1.5 for the standard disk models ofT Pyx and U Sco, respectively (Schaefer 2005). For T Pyx,we have repeated the calculations for a wide range of diskmodels, where a variable accretion rate goes through a diskwith the inner and outer radii of the disk widely varying withand without irradiation on the disk, while we find that E for theB band varies from 1.0 to 3.1, with the higher values going forthe higher accretion rates and the lower values going for diskswith severe inner truncation. The 1.21 M� models of Wade &Hubeny (1998) have exponents of 1.31 and 1.41 for 1500 Å and2000 Å at the highest range of accretion. Retter & Leibowitz(1998) give a simplistic theoretical model (Lbol ∝ M , F ∝ Lbol,so F ∝ M) that gives E = 1, but this has not included thecritical variation of the bolometric correction as the accretionrate changes. With this wide variation in models, it is clear thatthe Schaefer (2005) value of E = 2.0 could well be wrong. Inaddition, it is now clear that we can expect no confident guidancefrom models or theory, even within the framework of a diskmodel, and especially since the disk does not dominate the bluelight. Thus, we can only turn to empirical evidence regarding Efor T Pyx.

If the exponent is around 1.0, then the predicted date wouldhave been close to 2011. This is the case where the accretionrate is simply proportional to the B-band flux. The quiescentmagnitude from 1944 to 1967 was 14.72, while from 1967 to2005 it was 15.49 (Schaefer 2005), for a magnitude differenceof 0.77 mag, corresponding to a ratio of B-band flux of close to2.0. For the accretion rate being proportional to the B-band flux,the average accretion rate must then also have a ratio of closeto a factor of 2.0, and the post-1967 interval must be 2.0 timesthat of the pre-1967 interval. This would imply an eruption2.0×22 yr after 1967, which predicts the next eruption in 2011.So it seems that, somehow, T Pyx effectively has the B-bandflux being proportional to the accretion rate (i.e., E ≈ 1).

This fast and easy analysis can be substantially improved.Using the quiescent light curve and the inter-eruption timeintervals, the value of E that makes for the most constant〈FE

B 〉ΔT from interval to interval can be found. To determinethe best-fit exponent, we have taken our exhaustive light curvesfrom each of the intervals (see Table 2) and formed averages ofFE

B (see Table 3). The only free parameter in this analysis is E.The best estimate for E will be that value which returns the mostconstant values for 〈FE

B 〉ΔT across all measured time intervals.The values and the scatter will scale as E changes, so we findthe value of E that has the least fractional scatter. The smallestfractional scatter is near E = 1.13, with 13% rms variationsacross the five intervals with measures. With this, Table 3 hasthe value of 〈F 1.13

B 〉ΔT in the last column, with a straight averageof 71. Table 3 also lists the values of 〈FB〉ΔT (with a straightaverage of 64).

The scatter in 〈FEB 〉ΔT is much larger than the formal

uncertainties, and this indicates that our basic model is not exact.So, at the 13% level, the imperfect sampling of the light curvemight be misrepresentative, the value of E varies, or the triggermass might vary from eruption to eruption. For whatever reason,this means that the total uncertainty in 〈FE

B 〉ΔT is around 13%,so for five independent measures the uncertainty on the averageis around 6%. For predicting any single inter-eruption interval,the uncertainty will be 13%. Adopting a 13% uncertainty, wecan make a chi-square fit for E, with this giving a value of 1.13 ±

0.14. As such, we see that the five inter-eruption time intervalsare consistent with E = 1.

So we now have a plausible explanation for why T Pyx eruptedin 2011. The observed flux from T Pyx has been declining forthe last 121 yr, so the accretion rate has been falling and it takeslonger times between eruptions to accumulate the nova triggermass. In the 22 yr inter-eruption interval before the 1967 event,T Pyx was nearly a factor of two times brighter on averagecompared to the next inter-eruption interval, so the next inter-eruption interval should be roughly twice as long (i.e., around44 yr). Therefore, the next eruption after 1967 should occur in1967+44, which is to say in the year 2011. A more accuratepostdiction is to use the average value of 〈FB〉ΔT = 64 ± 9and 〈FB〉 = 1.47 ± 0.011 to get ΔT = 43.6 ± 5.9, whichputs the eruption peak in 2010.7 ± 5.9. The uncertainty on thispostdicted date is not small, so the close coincidence betweenthe postdiction and the real eruption date is fairly surprising.

5.5. The Nature of the Luminous Hot SourceIrradiating the Companion Star

Four different groups of researchers have concluded, ongreatly different grounds, that T Pyx harbors a very hot andluminous source irradiating the companion star and surroundingshell. Contini & Prialnik (1997) require that the entire surfaceof the white dwarf emit with a temperature of 250,000 K (for aluminosity of over 1036 erg s−1) so as to have the irradiation ofthe surrounding shell produce the observed spectrum. Pattersonet al. (1998) needed a supersoft source to produce the highobserved luminosity as driven by the high observed accretionrate. Knigge et al. (2000) invoked a wind-driven supersoftsource so as to explain the central T Pyx mystery of how abinary with such a short orbital period (so that only generalrelativistic effects should be driving a very low mass transfer)can possibly have high enough of an accretion rate to make thesystem recurrent with a 12–44 yr recurrence timescale. Schaefer& Collazzi (2010) needed the 1866 classical nova eruption toturn on the irradiation so as to make the high accretion rate thatis secularly declining, both to explain the fading in quiescentbrightness from B = 13.8 in 1890 to B = 15.7 in 2010 as wellas to explain why the inter-eruption intervals have been gettinglonger and longer since 1890. Thus, we see that the existenceof a hot and luminous source would explain many of the centralmysteries of T Pyx, while no alternative idea has even beenproposed.

Four mechanisms have been proposed to power the hot-luminous source near the surface of the white dwarf: Contini &Prialnik (1997) envisioned simply the surface of the white dwarfbeing very hot for unspecified reasons, likely due to residual heatfrom the eruptions. Patterson et al. have their supersoft sourcepowered by nuclear burning of the accreted material on thesurface of the white dwarf. Knigge et al. (2000) have the masstransfer rate sufficiently high that the accretion energy powersthe source. Schaefer & Collazzi (2010) have the accretion beingfunneled onto the polar cap of a highly magnetic white dwarfso that the local accretion rate is sufficiently high to makefor steady hydrogen burning, and this nuclear power is whatruns the hot-luminous source. Importantly, uncertainties in themechanism should not hide the broad agreement about what isgoing on in T Pyx (where the hot-luminous source is irradiatingthe companion, driving the high accretion rate, and determiningthe evolution of the system). Also importantly, the mechanismoperating on T Pyx need not be the same mechanism operating

20


on either the other short-period V1500 Cyg star systems or thelong-period V1500 Cyg star systems.

Despite the strong needs for a central hot-luminous sourcein T Pyx, the idea has the substantial problem that it has neverbeen directly seen, and the X-ray limits put strong constraintson possible sources. Selvelli et al. (2010) demonstrate that ablackbody with a temperature of 240,000 K (kT = 0.020 keV)at 3 kpc with 1037 erg s−1 luminosity is nearly two ordersof magnitude above the threshold for the XMM limit. Thisexplicitly excludes the default cases proposed by Contini &Prialnik (1997), Patterson et al. (1998), and Knigge et al.(2000). (Schaefer & Collazzi (2010) did not propose any specificparameters for their hot-luminous source.) So for the hot-luminous source idea to pass, there must be some device tohide the X-ray flux. Here are four possibilities: First, the veryhot nuclear burning region can be covered with a shroud of gasthat is substantially colder, so that the luminosity comes out inthe unseen FUV, with essentially no flux above 0.2 keV. Second,the feedback between the accretion luminosity and the accretionrate is necessarily unstable, and the accretion rate is seen to befalling, so the current luminosity might well be below the XMMthreshold. Third, perhaps most of the heating of the companionstar atmosphere occurred soon after the 1866 classical novaeruption, and the high accretion rate has been driven by thislatent heat, so the hot-luminous source has long been turned off.(Determining the timescale for restoring the companion star toits pre-1866 atmosphere is a complex unsolved problem whichdepends greatly on the assumed model parameters.) Fourth andsimplest, maybe the hot-luminous source is simply cooler thanan effective temperature of ∼50,000 K (e.g., the 34,000 K givenby Gilmozzi & Selvelli 2007), and all would be well. Given thesemany reasonable resolutions, we judge that the X-ray limit isnot a serious problem for the hot-luminous source idea.

Nevertheless, the particular mechanisms have substantialproblems.

The possibility that the hot-luminous source is powered bysteady nuclear burning of accreted material (Patterson et al.1998; Schaefer & Collazzi 2010) has the problem that theburning itself gets rid of the very fuel needed for the RNeruptions. That is, with the hydrogen being burnt up by thesupersoft source, there will be no hydrogen left to power thelater nova events. (Importantly, this problem does not apply tothe other V1500 Cyg stars, because they have no nova eventuntil long after the supersoft source has turned off and normalhydrogen-rich material has had plenty of time to accumulate onthe surface of the white dwarf. So this is a problem for T Pyxalone.) A possible dodge to this conclusion is that the hydrogenmight only partially burn, with enough hydrogen accumulating,perhaps away from a magnetic pole cap, to allow for a normaleruption.

The possibility that the hot-luminous source is powered bythe accretion energy (Knigge et al. 2000) has the problemthat the current accretion rate is very close to the limit forsteady hydrogen burning, so that the rate over a century agowould certainly be above the threshold. The accretion rate is(1.7–3.5) × 10−7 M� yr−1, as measured from the steady periodchange (P ) in the O − C diagram for the time interval 1968 to2011. But back before the 1890 eruption, the 〈F 〉 value was5.7× brighter, so with E = 1 the accretion rate must have been5.7× larger with M = (2.6–5.5) × 10−6 M� yr−1. For a 1.3 M�white dwarf, the threshold for steady hydrogen burning is(2–3)×10−7 M� yr−1 (Nomoto 1982; Yaron et al. 2005; Nomotoet al. 2007; Shen & Bildsten 2007). Thus, before the 1966

eruption, T Pyx was certainly accreting at a rate that places it inthe steady hydrogen burning regime. But T Pyx was undergoingRN eruptions from 1890 to 1966, so there is something wronghere. A possible solution to this dilemma is that the highlyvariable and fast changing accretion history makes the thermalstructure of the white dwarf far from equilibrium and colder thanexpected for its current accretion rate, and this will effectivelyshift the threshold for steady hydrogen burning to some largerlimit. That is, for a million years before 1866, the white dwarf’sthermal structure equilibrated to a very low M state with noeruptions, and then this equilibrium will be greatly perturbedby the various eruptions and the very high accretion rate, thusresulting in an out-of-equilibrium white dwarf that is muchcolder than expected for its current state as an RN. It will takedetailed physics calculations to test the size of the effect of thevariable accretion rate and eruption history.

In all, we judge the case for the existence of the hot-luminouscentral source to be persuasive but not proven, while the detailsof the mechanism that power it are not known, even whilereasonable possibilities are known.

6. A FULL MODEL FOR THE T Pyx SYSTEM

T Pyx has a fairly complicated history that has been piecedtogether from a wide variety of data. It might be useful to haveone consistent picture explicitly collected in one place. Here,we present a largely chronological summary based primarily onthis paper, Schaefer (2010), Schaefer & Collazzi (2010), andSchaefer et al. (2010a).

For an estimated interval of ∼750,000 yr before 1866, T Pyxwas a nondescript short-period cataclysmic variable (possiblywith a highly magnetic white dwarf) with an accretion rate ofaround 4 × 10−11 M� yr−1 and B = 18.5 mag. After this longand slow accumulation of mass on the white dwarf, a normalclassical nova eruption occurred in the year 1866 ± 5, reachingperhaps B = 5 mag. This eruption sent out a smooth shellwith a mass of ∼10−4.5 M� (larger than the accumulated masson the white dwarf) at a velocity of 500–715 km s−1 (for thetraditional distance of 3500 pc). This eruption irradiated thenearby companion star (which is very close due to the shortorbital period) so as to drive a high rate of mass being pushedthrough the L1 point in the Roche lobe. With some combinationof the high accretion energy and steady hydrogen burning(perhaps on a magnetic polar cap), the white dwarf remainedas a hot high-luminosity source to continue the irradiation andkeep the high M at a nearly self-sustaining rate. After the 1866eruption was long over, T Pyx stayed as a bright high-accretionsystem, much in contrast to its state a few years previously,which is to say that it is a V1500 Cyg star. With its newfound highaccretion rate, the white dwarf accumulated material quickly,and with it being near the Chandrasekhar mass, the systemstarted producing RN events. Each successive RN eruption sentout a fast (2000 km s−1) shell of relatively low mass, with theseshells ramming into the slower 1866 shell causing thousands ofknots to form by Rayleigh–Taylor instabilities. The feedbackloop of irradiation and accretion is not stable, so the accretionrate must be declining, with the typical timescale for such casesbeing a century or so. With this, the accretion rate must decline,the brightness of the system must decline (as observed to fadefrom B = 13.8 in 1890 to B = 15.7 in 2011), and the overallRN phase (starting in 1866) must have a brief duration. Anotherconsequence of the falling accretion rate is that the duration ofthe inter-eruption intervals must increase.

21


Two or three requirements are needed for this history to applyto T Pyx: a white dwarf near the Chandrasekhar mass, a veryshort orbital period, and possibly a significant magnetic field onthe white dwarf. The combination of the first two requirementsis rare, so there will be few such stars in our Milky Way and thedurations of their RN phase will be short, perhaps one to a fewcenturies out of millions of years. Nevertheless, their frequentand bright eruptions will call attention to these systems from faracross the Galaxy.

Although T Pyx is a system with material falling at a veryhigh rate onto a white dwarf near the Chandrasekhar mass,it is not a progenitor of a Type Ia supernova. First, the massgain/loss in the system is dominated by the classical novaevents, and these always eject more mass than has been accreted.Second, the RN events also eject more mass than they accrete, aswe have just learned from the measured period changes acrossthe 2011 eruption of T Pyx, the 2010 eruption of U Sco, andthe 2000 eruption of CI Aql. Third, systems like T Pyx aretoo rare and their lifetimes too long to provide enough Type Iasupernovae. (This statement is not true for the long-period RNe.)Fourth, half of all RNe are neon novae, with this making Type Iasupernovae impossible. So, not only is T Pyx not a progenitor,but none of the RNe are progenitors.

The irradiation of the low-mass companion star causes ahigh accretion rate, which is measured to be (1.7–3.5) ×10−7 M� yr−1 as based on the dynamics of the mass transfer.This rate is close to the threshold for which a white dwarf willsteadily burn the accreted hydrogen. The mechanism to powerthe hot-luminous source is not known, but published possibilitiesinclude using the gravitational energy of the accreted materialor the nuclear energy from steady hydrogen burning, while thematerial may or may not be funneled onto a magnetic pole of thewhite dwarf. The hot-luminous source cannot have an effectivetemperature over ∼50,000 K so as to avoid the limits fromthe X-ray flux. The ultraviolet–optical–infrared SED followsclosely to a power law with fν ∝ ν1.0, although the relativelynarrow ultraviolet region is tilted to more like fν ∝ ν1/3. Thebulk of the flux cannot come from an accretion disk, or anysuperposition of blackbody spectra, so most of the accretionlight is coming out from some nonthermal emission process(free–free, cyclotron, or Compton scattering) out of a volumewithin the Roche lobe of the white dwarf that is largely opticallythin. With the realization that the optical light is not from a disk,we can break out of the idea that leads to the incorrect predictionthat the next eruption was far into the future. Instead, we havesomething close to FB ∝ M . As such, the 22 yr inter-eruptioninterval before 1967 (with 〈B〉 = 14.70 mag) has a flux roughlya factor of two brighter than the inter-eruption interval after1967 (with 〈B〉 = 15.59 mag), so its duration should be roughly2 × 22 yr and the next eruption around the year 1967+44 =2011.

Based on the prior average inter-eruption interval, manypeople expected an imminent nova event starting in the mid-1980s. No eruption could have been missed (say, due tosolar gaps) because the T Pyx events have a long duration(roughly 230 days, with t3 = 62 days) and a bright peak (B =6.7 mag) such that the nightly observations throughout the entireobserving season would certainly have caught any eruption.Nevertheless, many observers in the professional and amateurcommunities persevered by taking many time series and nightlymonitoring for decades on end. From 1986 to 2011, the orbitalperiod had a steady increase at P/P = 313,000 yr, although asmall deviation appears from 2008 to 2011 as expected for the

small observed drop in accretion luminosity over this time. Inthe several years before 2011, T Pyx averaged V = 15.5 magand was never seen brighter than V = 15.0.

It came as a surprise to Stubbings when T Pyx appeared atV = 14.5 on UT 2011 April 5.51, and he realized that thismight presage an eruption. Instead, he saw a unique rise and fallin brightness with a peak at V = 14.4 over several days. Thispre-eruption rise started just 18 days before the eruption. Onlyfour anticipatory rises have ever been seen (T Pyx, V533 Her,V1500 Cyg, and T CrB), with all being greatly different fromeach other, and all with zero theoretical understanding.

As part of his long series of monitoring of RNe, Linnoltsaw T Pyx at V = 14.5 on UT 2011 April 13.36 and realizedthat this was unusually bright. On the next night at UT 2011April 14.29, Linnolt saw T Pyx at V = 13.0 and immediatelyrecognized the eruption. Within 15 minutes the world had beenalerted, Plummer and Kerr soon confirmed the eruption, andNelson, Carstens, and Streamer started time series in BVI. Forthe first two days, the magnitude during this initial rise is closelyrepresented as that of a uniformly expanding shell. The steadyand frequent monitoring, principally by amateur observers, fastrecognition and notification, and speedy response around theworld all make for T Pyx being the first and only nova eruptionwhere the light curve was followed throughout the entire initialrise. Indeed, while all prior nova events have returned only afew single magnitudes in the late part of the fast initial rise, forT Pyx, we now have a complete record of magnitudes and colorsover the entire initial rise with high time resolution.

Throughout this work, the AAVSO provided charts, compar-ison sequences, data archiving, and alert notices. This work issupported under grants from the National Science Foundationand from NASA as part of the GALEX program. The NASAGalaxy Evolution Explorer, GALEX, is operated for NASA bythe California Institute of Technology under NASA contractNAS5-98034. A.U.L. acknowledges support from NSF grantsMPS 75-01890 and AST 0803158, and from AFOSR grants77-3218 and 82-0192. A.P. is supported by the Edward CorbouldResearch Fund and the Astronomical Association of Queens-land. M.S. thanks the AAVSO for the loan of the SBIG 402camera. W.G.D. thanks Arne Rosner at Global Rent-A-Scope inMoorook, South Australia.

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