the spatially resolved bipolar nebula of sakurai's object - IOPscience

The Astrophysical Journal, 785:146 (12pp), 2014 April 20  C 2014.

doi:10.1088/0004-637X/785/2/146

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THE SPATIALLY RESOLVED BIPOLAR NEBULA OF SAKURAI’S OBJECT Kenneth H. Hinkle and Richard R. Joyce National Optical Astronomy Observatory, P.O. Box 26732, Tucson, AZ 85726, USA; [email protected], [email protected] Received 2013 November 19; accepted 2014 March 10; published 2014 April 4

ABSTRACT Sakurai’s object (V4334 Sgr), the final flash object discovered in the mid-1990s, underwent rapid cooling during the first decade of the 21st century becoming as faint as K ∼ 25. This stage of evolution has ceased. Between observations in 2010 September and 2013 April V4334 Sgr brightened >2 mag to K = 14.2 and the effective temperature increased to ∼590 K. AO images show a central source and two extended globules defining a 13◦ position angle. The globules span a spatial extent of ∼0. 3 in 2013. This spatial extent is consistent with sizes derived from spectral energy distributions taken over the previous decade and a debris cloud expanding at 0.055 mas d−1 since late 1998. Near-simultaneous 0.85–2.5 μm spectra reveal helium lines attributed to a wind-interaction shock. The He i 1.0830 μm emission has a spectral width of ∼1000 km s−1 and a spatial extent of ∼1. 4. The helium shell is fragmented, spatially asymmetric, and five times larger than the dust debris cloud. [C i] and [N i] forbidden lines are present in the 1 μm region spectrum. The forbidden line spectrum is similar to that of proto-planetary nebulae. The [C i] 9850 Å line is spatially extended. The expansion velocity and change of angular size limit the distance to 2.1–3.7 kpc. Key words: infared: ISM – planetary nebulae: general – stars: evolution – stars: mass-loss – stars: AGB and post-AGB – stars: winds, outflows Online-only material: color figure (VLTP) which occurs in about 10%–15% of all AGB stars (Lawlor & MacDonald 2003). Termination of AGB evolution occurs when the hydrogen-rich envelope drops below a critical value of ∼10−3 M (Lawlor & MacDonald 2003). In a VLTP the helium flash convection zone ingests this hydrogen-rich envelope (Herwig 2001b). Zijlstra (2002) estimates one final flash star should occur per decade in the galaxy so while this is a common process the number of observable bright events will be small. The discovery of a new final flash object in the 1990s provides the opportunity for observation at wavelengths other than the visual and in much more detail than was possible in the early 1900s. These observations provide critical feedback to refine the theory of stellar evolution. Since its discovery V4334 Sgr has been observed by numerous groups over much of the spectrum and extensively in the optical and near-IR. Early observations of brightness and abundance changes were well matched by theory (Herwig et al. 2011). However, models have not been successful at predicting evolutionary times for V4334 Sgr (van Hoof et al. 2008). van Hoof et al. (2007) notes that the time scale for the reheating of the remnant is a critical comparison with models. The need to understand the time scale and nature of the continuing evolution of V4334 Sgr was the primary motivation for the current work. We review the observations of V4334 Sgr below. We will then discuss our near-infrared observations of the expanding remnant. Finally we will discuss these observations in the context of the previous data on V4334 Sgr, other final flash objects, and evolutionary models.

1. INTRODUCTION On 1996 February 20 an 11.4 mag “nova-like object” was discovered in Sagittarius by the amateur astronomer Y. Sakurai and reported by Nakano (1996a). Sakurai’s object would later be given the variable star name V4334 Sgr. In the same IAU Circular as the discovery announcement Benetti & Duerbeck (1996) found this variable to be a reddened G-type star with no emission lines, characteristics strongly suggesting that V4334 Sgr was not a nova. Within a week Duerbeck (1996) and Pollacco (1996) had independently discovered that V4334 Sgr was at the center of an old planetary nebula (PN). The PN required that V4334 Sgr had previously entered the white dwarf sequence. This confirmed a claim made a few days earlier by Benetti et al. (1996) that V4334 Sgr was a rare example of a star undergoing final helium shell flash burning. Benetti et al. (1996) had noted similarities to the 1919 eruption of the final flash star V605 Aql (see also Clayton & de Marco 1997; Clayton et al. 2013). Calculations by Paczy´nski (1970) first showed that it is possible for a stellar remnant to undergo a post-asymptotic giant branch (AGB) shell flash. Further work by Renzini (1979) and Iben et al. (1983) resulted in the association between postAGB shell flashes and the H-deficient central stars of PNe. The models suggest that the final flash process takes decades (Iben & MacDonald 1995; Bl¨ocker 1995) to centuries (Herwig et al. 1999; Lawlor & MacDonald 2003). However the time scales of both observed events, V4334 Sgr and V605 Aql, are shorter. Solutions to the time scale problem invoke large deviations from mixing length theory to suppress convection (Herwig 2001b) and additional burning episodes (Lawlor & MacDonald 2003). A relatively low mass remnant is required, ∼0.5–0.6 M (Lawlor & MacDonald 2003; Hajduk et al. 2005; Miller Bertolami et al. 2006; van Hoof et al. 2007), which evolves from a zero-age main sequence star of mass near 2.5 M (Miller Bertolami et al. 2006). In both V605 Aql and V4334 Sgr the existing old PNe imply that the helium shell flash is a very late thermal pulse

2. A BRIEF HISTORY OF V4334 SGR 2.1. Progenitor The progenitor for V4334 Sgr was identified in a 1976 ESO sky survey image as a 21st magnitude blue star (Duerbeck & Benetti 1996). The PN that surrounds this star (Duerbeck 1996; Pollacco 1996) was an early argument for final flash status. The properties of the PN show that the post-AGB object had a 1

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temperature of 98,000 ± 7000 K, implying that it had reached the top of the white dwarf cooling track (Kerber et al. 1999). This observation supports VLTP evolution. From the PN spectrum the radial velocity was found to be +113 km s−1 , which matched the velocity of the central star, +115 km s−1 , confirming the connection between these objects. The expansion velocity of the PN is 25 km s−1 (Duerbeck & Benetti 1996). While the nebula is roughly spherical, a hole is visible in the northeast part, with the brighter region providing a P.A. ∼ 130 ◦ ± 8◦ (Chesneau et al. 2009). V4334 Sgr is a member of small family of final flash objects. These are characterized by old hydrogen-rich PNe and compact hydrogen-poor inner nebulae. The class has four other widely recognized members, V605 Aql, FG Sge, and the central stars of A30 and A78 (Hinkle et al. 2008). V605 Aql has been mentioned above. FG Sge is undergoing much slower evolution and is likely a late-thermal pulse object (Jeffery & Sch¨onberner 2006). A30 and A78 underwent a final flash event in the last few thousand years and now appear as hot dust-free Wolf–Rayet (WR) central stars of compact hydrogen-deficient nebulae. Duerbeck & Benetti (1996) estimate the time since the ejection of the hydrogen-rich PN around V4334 Sgr as 17,000 yr. The ages of the V605 Aql, A30, and A78 PNe are similar, 23,000 yr, 13,000 yr, and 34,000 yr, respectively (Guerrero & Manchado 1996).

to 8 kpc have been found based on Galactic structure (Duerbeck et al. 1997) and the nature of the old PN (Duerbeck & Benetti 1996). From the magnitude of the pre-final flash star the distance can be compared to that of similar central stars of PNe. Duerbeck pc. et al. (2000) found that this implied a distance of 1900 +1400 −800 They, however, suggested a distance of 3–5.4 kpc based on the stellar luminosity. Duerbeck et al. (2000) also adopted a mean value of the reddening at EB−V = 0.8. From Hα/Hβ in the old PN van Hoof et al. (2007) determined the extinction, EB−V = 0.86. Kimeswenger & Kerber (1998) previously had shown that the field stars have increasing extinction up to 1.8 kpc with constant EB−V = 0.9 ± 0.09 between 1.8 and 5 kpc. Based on this van Hoof et al. (2007) concluded that a lower limit distance to V4334 Sgr is 1.8 kpc. Chesneau et al. (2009) found that evolution models required a distance in excess of 1.5 kpc. van Hoof et al. (2007) note that the line of sight to V4334 Sgr reaches the scale height of the old disk at a distance of about 4 kpc. 2.4. Enshrouded by Dust Duerbeck et al. (2000) report that V4334 Sgr had an infrared excess, probably due to free–free emission in the stellar wind, in 1996. However, following the transformation of V4334 Sgr to a cool carbon star in 1997, an infrared excess attributable to carbon-rich dust was widely reported (Eyres et al. 1998; Kamath & Ashok 1999; Kerber et al. 1999). At that time the derived dust temperatures range from 680 to 1800 K with the hotter temperatures implying dust as close as 7 R∗ above the photosphere. As V4334 Sgr became more obscured by dust in 1998, temperature estimates converged to values near 1000 K (see for example Lynch et al. 2002). By early 1999 V4334 Sgr was no longer detectable at optical wavelengths (Duerbeck et al. 2000; Tyne et al. 2000). Bond & Pollacco (2002) provide an Hubble Space Telescope (HST) upper limit on the V magnitude of 24.2 by 2000 August. From mid-1999 onward the dust temperature decreased with time, reaching a lowest measured value of 180–210 K in 2005 April (Evans et al. 2006). Table 1 summarizes temperatures from blackbody fits to the IR colors. These values are plotted in Figure 1. Worters et al. (2009) have figures of the 3–5 μm spectra in 2005 August and 2007 June. While these spectra cover insufficient range to fit a blackbody, the continuum in 2007 is formed at a lower temperature than in 2005. The blackbody angular diameter,1 θ , can be computed using the relation

2.2. Return to the AGB Possible prediscovery observations reported in Duerbeck & Benetti (1996) and Duerbeck et al. (1997) suggest that the final flash for V4334 Sgr was already underway in 1994 June. If this observation is correct V4334 Sgr brightened from 15.5 in 1994 June to 12.4 at the time of the first definite pre-discovery observation on 1995 February. The discovery magnitude was 11.4 in 1996 February. Thus the outburst of V4334 Sgr occurred at least a year before the 1996 discovery. Following 1996 February V4334 Sgr evolved very rapidly. The effective temperature was 8000 K in 1996 March (JD 2450144), dropping to ∼6000 K 400 days later (Duerbeck et al. 1997). Between 1996 May and October Asplund et al. (1997) observed the surface hydrogen abundance decrease by 0.7 dex. This was accompanied by increases in the abundance of Li, Sr, Y, and Zr. The temperatures in a white dwarf merger would destroy Li, so the observation of Li provides a strong constraint against this type of event (Clayton et al. 2011). The 12 C/13 C ratio measured in 1996 October and confirmed later by Pavlenko et al. (2004) is very low, in agreement with ingestion of hydrogen. Duerbeck et al. (1997) report that in early 1997, less than two years after discovery, the spectrum was that of a cool carbon star. This was confirmed by the IR spectroscopy of Eyres et al. (1998) that showed photospheric lines of the carbonrich diatomic molecules CN, C2 , and CO. The spectrum of V605 Aql was similarly hydrogen deficient and carbon rich in 1921 a few years after its discovery (Clayton et al. 2013). In 1998 V4334 Sgr started to undergo R CrB-type fading events. Simultaneously He i 1.0830 μm was detected with a line width requiring a high velocity wind with an expansion velocity of ∼500 km s−1 (Duerbeck et al. 2000).

−2 θ (arcseconds) = 2.0 × 1011 (λFλ )1/2 , max T

(1)

where (λFλ )max is the apparent maximum flux observed in watts cm−2 and T is the blackbody temperature in kelvins (Gallagher & Ney 1976). This relation was applied by Duerbeck et al. (2000), Tyne et al. (2002), and Evans et al. (2006) to infrared spectrophotometry of V4334 Sgr. Evans et al. (2006) showed that the angular diameter of the dust shell grew linearly at the rate 0.0586 mas d−1 from the start of the expansion in 1997 November (JD 2450763) through 2003 September (JD 2453500). Comparison with the visual light curve in Tyne et al. (2002) shows that the origin of the expansion coincides closely with the time when the light curve began its deep minimum between 1997 November and 1998 February. As noted

2.3. Distance From observations of the old PN at 4.86 GHz Eyres et al. (1998) derived a statistical distance of 3.8 ± 0.6 kpc with a possible range of distances of 1.9–5.3 kpc. Other distances up

1

2

In this relation θ is the diameter. In Gallagher & Ney (1976) θ is the radius.

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Table 1 V4334 Sgr Blackbody Temperatures Date 1997 Feb 1997 May 1997 Apr 1998 Mar 1998 May 1998 Sep 1999 Apr 1999 May 1999 Jun 1999 Jul 1999 Jul 1999 Sep 1999 Sep 2000 Apr 2000 Apr 2000 Jun 2000 Jun 2000 Jul 2000 Oct 2001 Feb 2001 May 2001 Jun 2001 Sep 2003 Sep 2005 Apr 2007 Jun

Ta (K)

TBB b (K)

Reference

1500 1800 680 1170 1030 960 1285 700 ··· ··· ··· ··· 740 ··· 670 680 620 25. However, observations suggesting an emerging bipolar nebula imply that the blackbody cooling of V4334 Sgr has terminated. To distinguish between these models we imaged the field of

3.1.1. Coordinates

The coordinates determined for V4334 Sgr by Y. Kushida (Nakano 1996b) are α2000 = 17h 52m 32.s 69 δ2000 = −17◦ 41 07. 7. 4

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Eyres et al. (1998) give a radio position centroid of the old PN as α2000 = 17h 52m 32.s 75 ± 0.02 δ2000 = −17◦ 41 06. 0 ± 0.3 in reasonable agreement with the Kushida position but differing by more than 1 arcsec in declination. Fragmentation of the old PN results in the uncertainty in the PN center derived from this data. The 2MASS position α2000 = 17h 52m 32.s 694 δ2000 = −17◦ 41 08. 01 has a quoted uncertainty of 0.06 arcsec, although the survey pixel scale was 1.0 arcsec and V4334 Sgr was extremely bright (Ks = 6.3) at the time of the survey. Images of the V4334 Sgr field were obtained by Bond & Pollacco (2002) using HST WFPC2. Images taken when V4334 Sgr had become very faint allowed precise determination of the image centroid relative to other faint stars in the field. Using the world coordinate system (WCS) from HST images taken in 1999 and 2001 through the F814W filter, we determined the position of V4334 Sgr to be α2000 = 17h 52m 32.s 703 ± 0.001 δ2000 = −17◦ 41 07. 88 ± 0.01.

Figure 3. 2013 Jun 14 GNIRS spectrum of V4334 Sgr with a 590 K blackbody fit to the K band continuum.

3.1.2. Re-detection

3.2. Spectroscopy

Eight stars within 3 arcsec of V4334 Sgr in the HST images were also visible in the 2010 September Gemini AO images, permitting a correlated astrometric solution. These stars are sufficiently close to the center of the field that the radial barrel distortion of the f/32 NIRI/Altair configuration can be ignored.3 A faint, slightly extended object was indeed present in the Gemini AO images very close to the HST coordinates of V4334 Sgr. Encouraged by this finding, we obtained a second Ks image of the field with NIRI/Altair on 2013 April 26 (JD 2456408) using the same configuration, observing protocol, and reduction techniques as in 2010 (Figure 2). The 2013 image showed a bright (Ks ∼ 14.2) pointlike source located at the northern periphery of the extended source seen in the 2010 Ks image. Using the WCS parameters of the 2001 August 29 HST image, a grid of 73 stars visible in both the 2001 HST and 2013 NIRI images was used to recalculate the WCS parameters of the NIRI image using the IRAF task ccmap. In this system, we were able to determine the coordinates of the infrared point source in the 2013 image, which we believe to be associated with V4334 Sgr, as α2000 = 17h 52m 32.s 703 ± 0.002 δ2000 = −17◦ 41 07. 92 ± 0.02. The uncertainties given are those of the coordinate fitting in the ccmap task, which is probably limited by small proper motions of the field stars used to calculate the coordinate transformation over the intervening 12 yr interval. Several stars had obvious proper motion, ∼0. 1, over this period and one star was resolved as a binary in the NIRI image. These stars were rejected from the fit. Comparison of the 2010 and 2013 Ks images confirms that the two sources near the HST position of V4334 Sgr in 2010 are ejecta moving radially outward at P.A. 13◦ from that position, with the southern object becoming noticeably more diffuse in 2013. The elongated southern object in the 2010 Ks image could be a spatial blend of ejecta and the brightening point source. The pointlike source in the 2013 image is more than 2.5 mag brighter than the southern source in the 2010 image, setting a lower limit on the brightening of the point source in this time interval. We observed the V4334 Sgr field on 2013 September 20 (JD 2456555) using WHIRC on the WIYN 3.5 m telescope in the H and Ks bands. V4334 Sgr was clearly visible in both filters at H = 16.7, Ks = 14.0 but was unresolved at the ∼0. 9 delivered image quality of these images. The magnitudes confirm both the continuing brightening of the source in the near-IR and its very red color. 3

3.2.1. GNIRS

The 2013 image showed that V4334 Sgr had brightened sufficiently so that medium-resolution near-IR spectroscopy would be possible at least in the K band region. We obtained a spectrum using GNIRS (Elias et al. 2006) on the Gemini North telescope on 2013 June 14 (JD 2456457). The standard medium-resolution broad wavelength configuration was employed with the short camera and 31.7 line mm−1 grating used in combination with a cross dispersing prism giving wavelength coverage from ∼0.85 to 2.5 μm. A 0.45 arcsec wide, three pixel slit was used for resolving power λ/Δλ ≈ 1100. The slit was rotated to match the 13◦ P.A. of the ejecta. The spectrum was wavelength calibrated to angstroms in vacuum using the Ar hollow cathode spectrum obtained as part of the science observation. The typical rms accuracy of the identification fit was 0.4–0.6 Å, depending on the number of Ar lines within a given spectral order; over the 0.9–2.5 μm range, this is equivalent to ∼6–15 km s−1 uncertainty in the velocity. The A1V star HD 155427 was observed at a similar airmass to provide telluric line correction, after removal of the hydrogen lines from its spectrum. This was effective for all but the most opaque telluric absorption line regions. The telluric standard was also used to provide an approximate flux calibration, assuming similar slit losses and the use of a 9850 K blackbody for a constant magnitude. The resulting spectrum is shown in Figure 3. The regions obscured by optically thick telluric absorption are not plotted. The spectrum is characterized by a nearly featureless continuum in the K and H bands that is well fit by a 590 K blackbody (Figure 3). The only spectral feature in this region is an emission line at 2.059 μm. Near ∼1 μm no continuum was measurable, but several emission lines are apparent in the spectrum. These are listed in Table 2, with the flux in units of watts cm−2 . The FWHM is the Gaussian FWHM from the IRAF “splot” task. The deblending function provided in this routine was used to separate the close line pairs at 0.98 μm and at 1.1 μm. The 2.059 μm line is in the middle of a strong telluric CO2 band. The line was obvious after the initial telluric correction. Other suspect lines in the K band correspond to poorly corrected strong telluric lines. All but the very strongest telluric features near 2.059 μm could be corrected by small adjustments in wavelength and differential airmass in the IRAF “telluric” task, and both of the broad CO2 bands at 2.00 and 2.06 μm were well

http://www.gemini.edu/sciops/instruments/niri/imaging?q=node/10058

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Figure 4. 1.1 μm region of the spectrum in Figure 3 shown at higher dispersion. Table 2 V4334 Sgr Near-IR Spectral Lines Wavelengtha (μm) 0.9824 0.9851 1.0394 1.0831 1.0942 1.1005 2.0593

Wavenumbera (cm−1 )

Flux (w cm−2 )

FWHM (μm)

Identa (μm)

10179 10151 9621 9233 9139 9087 4856

5.9E−23 2.4E−22 2.4E−23 2.5E−22 3.4E−23 2.2E−23 4.1E−23

0.0016 0.002 0.0026 0.0024 0.0039 0.0041 0.0065

[C i] 0.9826 [C i] 0.9852 [N i] 1.0398 He i 1.0833 unid unid He i 2.0587

Figure 5. Dust mass in the ejecta around V4334 Sgr, computed from blackbody relations and assuming a distance of 3 kpc, as a function of time. The fit (dash line) has a slope of 0.00128 log(M ) d−1 .

4. DISCUSSION 4.1. The Debris Cloud Gehrz et al. (2005) provide a blackbody relation for deriving the mass in grams, Md , of a circumstellar dust shell from the temperature and luminosity:

Note. a Units in vacuum.

Md =

corrected, so we are confident in the reality of the emission feature at 2.059 μm. Depending on whether a Gaussian or Lorentzian is used to fit to the line the measured wavelength is 2.0593 or 2.0591 μm respectively, although the steep continuum makes the fitting procedure difficult.

(3)

where agr is the grain radius, ρgr is the density of the optically active grains, σ = 5.7 × 10−5 erg s−1 cm−2 K−4 is the Stefan–Boltzmann constant, Qe is the mean thermal emissivity of the optically active grains, and Tgr is the grain temperature. For amorphous carbon grains ρgr = 2.25 gm cm−3 . Values for Qe are given by Draine & Lee (1984). The total dust luminosity, Ld in erg s−1 , can be found for an assumed distance, D in cm, from the properties of the spectral energy distribution (SED):

3.2.2. Line Identifications

The He i 3 S–3 P line at 1.0830 μm is clearly present in our spectrum (Figure 4) and had been previously detected in V4334 Sgr. We identify the 2.059 μm line with He i 1 S–1 P 2.058129 μm, the singlet equivalent of the 1.0830 μm He i line. Three other lines in the spectrum are nebular forbidden lines, [C i] 0.9826 μm, [C i] 0.9852 μm, and [N i] 1.0398 μm. The three forbidden lines as well as helium lines have been identified in proto-PNe (Kelly et al. 1992). The lines at 1.0942 and 1.1005 μm remain unidentified. The line at 1.0942 μm closely matches the location of a He ii and/or hydrogen Paschen γ feature. However, the He ii/Paschen β line at 1.281 μm is not present in the spectrum. No suitable forbidden transitions at 1.094 and 1.100 μm appear in the NIST line compilation.4 No matching WR lines are present from the Mauerhan et al. (2011) list. Assuming that the two unidentified lines are real, the most likely identification is He i although the wavelengths do not match as well as might be expected. Hydrogen is ruled out because hydrogen lines are not seen at other wavelengths in the spectrum and because V4334 Sgr is known to be hydrogen poor. The same reasoning eliminates H2 . The [C i] 0.9852 μm and He i 1.0830 μm lines are clearly spatially extended on the raw spectral images, even under the ∼0.6 arcsec natural seeing conditions at the time of the GNIRS observations. 4

agr ρgr Ld , 3σ Qe Tgr4

Ld = 5.44π D 2 (λFλ )max .

(4)

Table 3 summaries the dust luminosity and mass computed from the above relations, assuming that the distance is 3 kpc and that (λFλ )max is covered by the partial SEDs available in literature spectrophotometry. From 1999 through at least 2006 the temperature of the V4334 Sgr debris cloud was decreasing (Figure 1). During this period the luminosity of the stellar remnant was roughly constant at ∼5600 L for a distance of 3 kpc. (Section 2.4 and Table 3). If the underlying assumptions in the blackbody relations have not been violated, constant luminosity requires that the dust mass increases as 1/T 4 . The blackbody dust mass (Equation (3)) ranges from ∼10−7 M to 10−4 M during the period from 1999 May to 2005 April (Table 3 and Figure 5). The large gas plus dust ejecta mass, 10−4 to 10−3 M , and extremely high mass loss rate implied by standard blackbody relations was noted previously by van Hoof et al. (2007). They find that applying standard gas-to-dust ratios is questionable at the very low temperature of the dust shell. They also question the validity of the dust relations given the extremely high

Interactive version available at http://www.nist.gov/pml/data/asd.cfm.

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Hinkle & joyce Table 3 Dust Luminosity and Mass

Date

λmax (μm)

(λFλ )max (W cm−2 )

θ (mas)

Ld a (L )

Qe

Md a (M )

References

1999 May 1999 Jun 1999 Jul 1999 Jul 1999 Sep 2000 Apr 2000 Jun 2001 May 2001 Sep 2003 Sep 2005 Sep 2005 Sep

4.2 4.3 4.5 4.6 4.6 5.4 4.3 5.0 5.5 8.0 11.3

1.51E−15 1.54E−15 1.61E−15 1.50E−15 1.71E−15 4.25E−15 1.43E−15 2.81E−15 3.07E−15 1.26E−15 1.46E−15 1.98E−15

16.5 17.4 19.5 19.7 21.0 45.6 16.5 31.8 40.3 55.3 117.0 203.0

5600 5600 5900 5600 6500 15750 5300 10100 11250 4800 5600 7300

0.2 0.2 0.2 0.15 0.15 0.1 0.2 0.13 0.1 0.04 0.02 0.015

6.2E−08 6.8E−08 9.0E−08 1.2E−07 1.4E−07 9.6E−07 6.2E−08 3.6E−07 7.3E−07 3.5E−06 3.2E−05 1.2E−04

Tyne et al. (2002) Tyne et al. (2002) Tyne et al. (2002) Tyne et al. (2002) Tyne et al. (2002) Tyne et al. (2002) Hinkle & Joyce (2002) Tyne et al. (2002) Tyne et al. (2002) Evans et al. (2004) Evans et al. (2006) Evans et al. (2006)b

Notes. a Computed assuming D = 3 kpc. b T dust from fit.

agreement with the expansion rate, 0.059 mas d−1 (for the diameter), found by Evans et al. (2006). The implication is that the two well separated dust clouds in the images are debris ejected in 1998/1999. If we assume that the debris cloud is expanding uniformly at 0.055 mas d−1 at a distance of 3 kpc, the debris radial expansion velocity is ∼150 km s−1 . The linear momentum in the debris exceeds L/c by orders of magnitude. There has been ongoing speculation that the Eddington limit, i.e., radiation pressure high enough to balance gravity, is exceeded in final flash objects and separates the envelope from the star (Faulkner & Wood 1985; Duerbeck et al. 2000; Miller Bertolami et al. 2006). Van Hoof et al. (2007) also found that the energy in the V4334 Sgr ejecta completely overwhelms the luminosity and suggest that the mass was lost in a brief event. 4.2. He i 1.0830 μm In 1998 March Eyres et al. (1999) first observed He i 1.0830 μm in V4334 Sgr. At that time the line was strong with a P Cyg profile. The He i 1.0830 μm line had not been present at all in 1997 July. The P Cyg profile required line formation in an extended region relative to the continuum. He i 1.0830 μm appeared at approximately the same time as the infrared excess due to dust (Duerbeck et al. 2000). A high-resolution spectrum in 1998 June (Eyres et al. 1999) showed blue shifted absorption with a full width of ∼550 km s−1 . Measured from the photospheric velocity the outflow velocity was ∼670 km s−1 . The absorption shifted blueward an additional ∼50 km s−1 from 1998 June to 1998 September (Eyres et al. 1999). By 1998 September the P Cyg profile had red emission and blue absorption separated by ∼800 km s−1 (Lynch et al. 2002). Tyne et al. (2000) and Geballe et al. (2002) report that the line profile changed from P Cyg to entirely emission by 1999 April. The full width of the emission was ∼1000 km s−1 in 1999 May and ∼1500 km s−1 in 1999 June (Geballe et al. 2002). In 1999 continuum could be detected under the emission line but the P Cyg absorption component was absent. The lack of absorption requires that the emission originated in a thin, semi-transparent shell of much larger surface area than the continuum forming region (Duerbeck et al. 2000). By 2000 April Geballe et al. (2002) report that continuum was no longer present in the 1 μm region. The mean velocity of the emission is always blue shifted

Figure 6. Dust shell diameter in mas. The filled circles are blackbody diameters from the SED. The filled triangles are diameters measured from the AO images. The dashed line fit shows a growth of 0.05542 mas d−1 . The fit goes to zero on 1998 November during the deep R CrB phase.

optical depth of the dust shell and the possibility that mass loss has ceased and the shell is hollow. Indeed, interferometric observations show the interior of the ejecta to be a compact disk (Chesneau et al. 2009). Starting in 2004 (Section 2.5) the V4334 Sgr debris were reported to contain reionized material with bipolar characteristics. Since then the blackbody assumptions may have been violated and it is likely that the debris cloud parameters cannot be computed simply. A SED with broad wavelength coverage would be of interest to look for multiple temperatures, as are seen in V605 Aql (Hinkle et al. 2008). Nevertheless, the dust cloud diameters resolved in the AO images are consistent with the angular diameters derived from blackbody relations before the nebula was resolved (Equation (1)) and a constant expansion rate (Table 3 and Figure 6). A least squares fit to the angular diameter as a function of time gives an expansion rate of 0.055 ± 0.008 mas d−1 . The angular diameter goes to zero at the time of the deep R CrB decreases in visual magnitude and is in good 7

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Figure 7. Emission line fluxes for He i 10830 Å and [C i] 9824, 9851 Å as a function of Julian date. Figure 8. Long slit image of He i 1.0830 μm. The peak emission defines the spatial origin at a heliocentric velocity of −240 km s−1 . The parallel lines define the extent of the 0.45 arcsec slit, with positive spatial offsets at a P.A. of 13◦ on the sky. One angstrom corresponds to about 27 km s−1 .

Table 4 He i 1.0830 μm Line Intensity History Date 1998 Mar–Aug 1999 Apr 1999 May 2000 Apr

Intensity (W cm−2 )

Reference

10 ± 0.3 × 10−20 11 ± 0.3 × 10−20 7.7 × 10−20 4.2 × 10−20

Tyne et al. (2000) Tyne et al. (2000) Figure 8 of Geballe et al. (2002) Geballe et al. (2002)

The spatial extent of the emission is ∼1. 4. At zero intensity the extent is closer to 1. 9. The dust continuum has a 0. 3 diameter, roughly a factor of five smaller than the He i 1.0830 μm nebula. Thus the dust does not obscure emission from the backside of the nebula. Helium emission is mainly seen from the front side of the nebula demonstrating that the helium emission does not originate from a symmetric shell. The He i long slit spectrum (Figure 8) shows that the velocities are skewed to the north, from an asymmetric shell or disk. The spatial information in the spectrum has a spatial resolution of ∼0. 6 arcsec. Thus the 0. 3 diameter dust forming region is unresolved and it is not possible to separate the He i contribution from this region from more spatially extended He i emission. Two stars in Figure 2, which we have named star 9 and star 12, lie within the He shell diameter. Due to the brightness of V4334 Sgr in 2013 and the processing of the AO image accurate photometry of these stars is difficult. To within the uncertainties of our measurements neither star has changed in brightness. However, given the multiple kiloparsec distance to V4334 Sgr it is likely that these stars are in the foreground. The integrated He i 1.0830 μm flux from the GNIRS spectrum has been included in Figure 7. The integrated flux from the 2010 NIRI He i 1.0830 μm image is also included. The image was faint and this value is uncertain but, nonetheless, in reasonable agreement with the 2013 spectrum. Since 1998 the flux has decreased by a factor of 400 while at the same time the velocity of the helium line has been roughly constant. If the He i flux arises from wind interaction then the area of the interaction has decreased at the same time the wind interaction front has moved away from the central object. This requires that the emitting area is fragmented.

by several hundred km s−1 relative to the stellar velocity. This requires an asymmetry in the emitting region with the far side of the emitting region either not present or obscured. Historical measurements of the pure emission He i 1.0830 μm line fluxes are given in Table 4 and plotted in Figure 7. He i 1.0830 μm is a triplet transition to a metastable lower level 20 eV above the ground state. LTE population of this level requires temperatures of ∼20,000 K. As noted by Tyne et al. (2000) the lack of He ii in the spectrum of V4334 Sgr requires that the He i 1.0830 μm line not be formed in LTE. Eyres et al. (1999) suggests that the He i originates in shocked gas dragged outward by the accelerating dust. However, the He i 1.0830 μm line is seen after dust obscured the stellar flux which requires that the P Cyg line be formed in a shell above the dust formation region. If, as seems appropriate for the metastable He i 1.0830 μm transition, the line originates in a wind interaction shock (Hinkle et al. 2013), the He i emission marks the boundary where the fast wind is colliding with the existing circumstellar material. Kerber et al. (2002) speculate that UV radiation related to the shock triggers dust formation through the ionization of key molecular species. The He i 1.0830 μm nebula is resolved in our long slit spectra. Figure 8 shows the spatial-spectral image of the He i 1.0830 μm line. The measured wavelength of the line peak (1.08241 μm) corresponds to a heliocentric velocity of −240 km s−1 which is −360 km s−1 relative to the stellar center-of-mass velocity. The integrated He i 1.0830 μm emission line (Figure 3), when corrected for the R ∼ 1100 intrinsic line width of ∼9 Å yields a FWZI of ∼1000 km s−1 , spanning the range ∼+200 to ∼−800 km s−1 relative to the star’s velocity.

4.3. CO The CO molecule is a ubiquitous monitor of cool gas in stellar atmospheres and circumstellar environments. CO was not reported in the near-IR spectrum of V4334 Sgr prior to 1997. 8

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CO Δv = 2 lines from both the 12 CO and 13 CO isotopologues are obvious in spectra from 1997 April and 1997 July (Eyres et al. 1998). The conspicuous bandheads in these observations require photospheric excitation temperatures. Pavlenko et al. (2004) find that CO Δv = 2 lines from 1998 July are blue shifted relative to the stellar velocity by 11 ± 3 km s−1 . Eyres et al. (2004) reports strong CO fundamental (Δv = 1) lines present by 2000 April. Unlike the CO first overtone lines, the low excitation CO fundamental lines were detected after dust obscured the photosphere and must be circumstellar in origin. This was clearly the case in 2003 September when Eyres et al. (2004) measured an excitation temperature of 400 ± 100 K and an outflow speed of 290 km s−1 . A blackbody fit to the 1–5 μm continuum on this date is 350 ± 30 K so the gas and dust are at similar temperatures with the CO in absorption against a dust continuum. Eyres et al. (2004) notes that the outflow velocity of CO was approximately constant from 2000 April through 2003 September and consistent with the 670 km s−1 broad He i 1.0830 μm profiles seen by Eyres et al. (1999). Worters et al. (2009) reports on observations spanning 1.5–5.3 μm. CO 1-0 lines are clearly present in 2003 and 2004 and likely present in 2005. A high-resolution spectrum taken in 2004 June shows that the lines are broad with FWZI = 400 km s−1 . The blueward edge of the CO had an expansion velocity of 500 km s−1 , similar to the He i velocity. The CO gas temperature was 320 ± 30 K, which is again in agreement with the dust temperature (Worters et al. 2009). The 12 C/13 C +2.0 ratio was 3.5 −1.5 . This isotopic ratio matches the previously determined isotopic ratio in the V4334 Sgr pseudo-photosphere and requires that the CO originate in AGB ejecta (Worters et al. 2009). Throughout the period when CO has been observed the lines are in absorption with rotational levels thermally populated near the dust temperature, as expected for CO originating near the continuum forming region. The absence of P Cyg profiles requires that the CO be formed in a region thin compared to the diameter of the dust shell. The CO velocities are similar to those seen in He i requiring that a significant section of the circumstellar shell is moving at this velocity. However, the dust expansion velocity is about five times less than the CO velocity. As discussed above a wind interaction model (Kenny & Taylor 2005) seems applicable. The He i shell is formed in the outer shock where the fast wind interacts with the old circumstellar environment. The CO and dust, which are physically in close proximity but have velocities differing by >100 km s−1 , mark the inner shock. The inner shock could also contribute to the He i 1.0830 μm spectrum but this is not spatially resolved in our data. The flux suggests that the He i 1.0830 μm emitting area is fragmented. The images of the V4334 Sgr debris clouds have globules similar to the infrared clouds observed in images of V605 Aql (Hinkle et al. 2008). Clayton et al. (2011) find similar globules in the circumstellar environment of R CrB. They propose that these are “cometary globules” which are slower moving ejecta interacting with a fast wind (Garc´ıa-Segura et al. 2006). Wind-shell interaction simulations show that secondary lobes in PNe can be formed through the interaction of a fast lowdensity wind with a complex filamentary shell (Steffen et al. 2013).

Figure 9. Long slit image of [C i] 9851 Å. The slit is shown as in Figure 8.

N i, N ii, O i, O ii, and S ii were detected in optical spectra of V4334 Sgr observed in the period 2001 through 2006 (Kerber et al. 2002; van Hoof et al. 2007). The emission line strength and the level of excitation were decreasing with time during this period. van Hoof et al. (2007) note that the constant decline is not consistent with a photoionization origin for the nebular lines and requires that these lines result from recombination following shock excitation. Our spectra include two lines, [C i] 9824 and [C i] 9850 Å previously observed by van Hoof et al. (2007). [C i] flux as a function of time is shown in Figure 7. The [C i] flux is now increasing. The [C i] 9850 Å line is spatially extended (Figure 9). The radial velocity shows a shifted outflow from the He i. Assuming the flow is radial the flow from the front of the nebula is slower, −400 km s−1 versus −800 km s−1 for the He i, while the flow in the back is faster, +300 versus +200 for the helium. As for helium, the flow in the back of the nebula is offset spatially. As noted above the spatial resolution leaves the contribution from the central region poorly defined and the line likely originates in both the extended wind interaction region and in the central region. The forbidden [N ii] 6583 Å transition in the spectrum of V4334 Sgr was previously found to be spatially extended (Kerber et al. 2002; van Hoof et al. 2007). van Hoof et al. (2007) measures a diameter for the [N ii] emitting region of 0. 3–0. 5 in 2003. The [C i] emitting region has a diameter of ∼1. 5 in 2013 (Figure 9). Assuming that the [N ii] and [C i] are formed in the same extended region the expansion rate of the gas is ∼0.28 mas d−1 , five times the expansion rate of the dust. Van Hoof et al. (2007) report that the radio flux increased significantly in 2006. They attributed this to photoionization of major constituent elements with the lowest ionization potential. The radio flux and the [N ii] flux were consistent with the presence of a 10,000–12,000 K Teff star driving the photoionization. Carbon was predicted to be photoionized. However, the [C i] flux is similar in 2013 to the 2003 measurement of van Hoof et al. (2007). [C i] appears to originate throughout the nebula and multiple excitation mechanisms are likely.

4.4. Forbidden Lines The GNIRS near-IR spectrum includes forbidden lines of [C i] and [N i] in the region near 1 μm. Forbidden lines from C i, 9

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As noted above, the 0.9–1.2 μm spectrum of V4334 Sgr is similar to that of proto-PNe (Kelly et al. 1992). Schmidt & Cohen (1981) found that in bipolar proto-PNe the forbidden line radiation originates in the lobes. Kelly et al. (1992) concluded that the forbidden lines are formed in proto-PNe in a region shielded from direct starlight that has an electron temperature of 10000 K and hydrogen density of ∼105 cm−3 . Schmidt & Cohen (1981) found that the central star exciting the forbidden lines in proto-PNe was a 30,000 K B star. Clumping allows narrow beams of ionizing radiation into the lobes but the bulk material is cooler and neutral. The absence of ionized lines in our spectra implies the electron temperature is less than about 10,000 K. However, Kelly et al. (1992) attribute the overall emission line spectrum of proto-PN to a combination of shock excitation and photoionization from the central star. Ratios of line fluxes can provide the nebular density, temperature, and extinction. It would be insightful to have this information across the resolved V4334 Sgr nebula. Unfortunately our wavelength coverage does not extend far enough to the blue to utilize this diagnostic.

N/O mass fractions for both V605 Aql and V4334 Sgr showing that both are helium rich but that V605 Aql is very oxygen rich. Oxygen-rich debris in the V605 Aql ejecta is surprising given the carbon-rich spectra of the pseudo-photosphere taken in 1921 (Clayton et al. 2013). If these abundance results are confirmed the processes creating V4334 Sgr and V605 Aql are very different. The evolution of V4334 Sgr is frequently discussed as faster than that of V605 Aql although this is difficult to quantify due to the sparse V605 Aql data. Arguably the evolution has not been slower. The 2 arcsec diameter V605 Aql nebula is ∼80 years old (Hinkle et al. 2008), giving an expansion rate of 0.034 mas d−1 versus 0.055 mas d−1 for V4334 Sgr. V605 Aql was found by Clayton et al. (2013) to have a geometrical distance of 4.6 kpc. The distance to V4334 Sgr is known with much less certainty but 2.8 kpc, which would result in equal physical expansion rates, is a good match to the geometrical bounds set by the velocity and expansion rate. V605 Aql has a FWZI for the He i 1.0830 μm line of ∼1000 km s−1 (Hinkle et al. 2008) similar to values found in V4334 Sgr. However, the dust mass in the cloud around V4334 Sgr is now 10−4 M (Section 4.1) as compared to ∼10−5 M for V605 Aql (Clayton et al. 2013). The larger dust mass and the comparable if not faster evolution require a more energetic mass loss event for V4334 Sgr. The central object of V605 Aql is now a hot dwarf with WR spectral features requiring an effective temperature of ∼95,000 K (Guerrero & Manchado 1996; Clayton et al. 2006, 2013). In the near-IR we detected only neutral forbidden lines from V4334 Sgr. While Mauerhan et al. (2011) catalog near-IR WR/WC features, the strong 5801 Å C iv line in the optical remains the most diagnostic feature (Guerrero & Manchado 1996). As the debris cloud around V4334 Sgr continues to expand the optical spectrum should be monitored for WR lines. Of the five well studied objects in the final flash class only V4334 Sgr and the slow final flash object FG Sge do not have a WR component to their spectra (Hinkle et al. 2008). The mass of the debris ejected from V4334 Sgr approximates the mass of the preflash envelope. Thus as the debris expands the intershell surface of the remnant will be exposed and V4334 Sgr should become a WR object. van Hoof et al. (2007) and Sch¨onberner (2008) compare the evolutionary predictions for V4334 Sgr from three models. The models of Herwig (2001a) assume a luminosity of 10,000 L . This implies a 4 kpc distance (Table 3) which is at the limit of the geometrical distance from the outflow rate. No model predicted the 10,000 K implied by the forbidden line spectra. However, as noted by Schmidt & Cohen (1981) the excitation temperature of these lines is lower than the temperature of the central star. In proto-PNe the exciting central star has an effective temperature near 30,000 K. The rapid time scale of Herwig (2001a) is closest to the observed changes.

4.5. Comparison with other Objects and Models Recent changes in the understanding of the central stars of PNe have altered the paradigm for the creation of these objects. As reviewed by De Marco (2009) the long running debate over the shaping mechanism for PNe is now reaching a consensus that single stars cannot trivially manufacture nonspherical PNe. Binary companions, possibly even substellar ones, are being invoked in theoretical models. For instance, Soker & Kashi (2012) argue that bipolar PNe are formed by the accretion disk around the companion to the AGB star. Both the He i cloud and the ejecta from V4334 Sgr are nonspherical. The globules ejected from V4334 Sgr define a 13◦ position angle. If we assume this is the disk axis then the position angle of the disk plane is 103◦ . As the ejecta disperses it should be possible to define the P.A. with more precision. The old PN position angle, defined by a gap, is 130◦ which is in excellent agreement with the position angle of 132◦ for the compact disk detected by Chesneau et al. (2009). V605 Aql has a position angle for the new debris cloud similar to that for the old PN (Hinkle et al. 2008). If, as is reviewed in De Marco (2009), rotation or magnetic fields cannot orient the debris over the required time spans to define the position angle then a companion object is required. The presence of this companion does not necessarily require that it be responsible for the eruption. As is the case in the central stars of many PNe Kerber et al. (2002) find no indication of binarity in V4334 Sgr. V605 Aql, which underwent an outburst in 1919, is generally considered a nearly identical final flash object to V4334 Sgr. Several authors have recently questioned the nature of the V605 Aql outburst. Wesson et al. (2008) measured a Ne mass fraction of 35% for V605 Aql inner ejecta which is much different from the 2% predictions of final flash modeling (Clayton et al. 2013). Lau et al. (2011) attempted to explain these abundance patterns with an ONe nova following the helium shell flash. Another explanation advanced is a merger of a ONe white dwarf and a companion. Hajduk et al. (2013) review the evolution of CK Vul, another of the possible final flash objects, as a white dwarf binary mass transfer. However, none of these scenarios can fully explain all the observations. V4334 Sgr has a much lower Ne mass fraction of 5.9% (Wesson et al. 2008) which is, however, a factor of three larger than the prediction. Wesson et al. (2008) also provides H/He/C/

5. CONCLUSIONS For the first time since the eruption of V4334 Sgr in the mid-1990s the circumstellar debris cloud has been imaged. Images were taken on two dates separated by 2.5 yr in the near-infrared Ks band using adaptive optics. The second image was followed by near-simultaneous 0.9–2.5 μm spectroscopy. The spectroscopy demonstrated that the Ks band, and hence the images, is dominated by continuum from a cool, ∼590 K, debris cloud. The 2010 image shows two sources connected by faint continuum. By 2013 a K ∼14.2 point-like source appeared at the northern periphery of the extended source. The northern part of 10

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The Gemini Observatory is operated by the Association of Universities for Research in Astronomy, under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Particle Physics and Astronomy Research Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), CNPq (Brazil), and CONICET (Argentina). The images were obtained as part of programs GN-2010B-Q-12 and GN-2013A-Q-26. The spectroscopy was made possible by time granted us through director’s discretionary time, program GN-2013A-DD-4. We are very appreciative of this special allocation of telescope time. This research would not have been possible without the SIMBAD database, operated by CDS in Strasbourg, France, and NASA’s Astrophysics Data System Abstract Service. We thank Lori Allen, Charles Corson, Meredith Drosback, and Doug Williams for obtaining WHIRC images of the V4334 Sgr field. Thomas Matheson provided useful comments on the GNIRS spectrum line identifications. Susan Ridgway provided help with the Gemini phase II process. Tom Geballe was very helpful with final adjustments to the Gemini phase II plan.

the doubled image seen in 2010 is now a well-separated source and a fan-like extension appears to the south of the bright source. The 0. 3 diameter across these clouds agrees with the diameters derived from by blackbody relations based on the SEDs with an expansion rate of 0.055 mas d−1 . The spectrum also reveals that the He i 1.0830 μm transition comes from an extended region. The He i 1.0830 μm has a typical expansion velocity of 500 km s−1 . The spectral image subtends ∼1.4 arcsec or more than four times the extent of the dust debris image. The line shape and excitation temperature of CO Δv = 1 lines reported in the literature require that these lines be formed just outside of the dust debris. The CO lines share the velocity of the He i 1.0830 μm lines. We propose that the V4334 Sgr circumstellar shell contains two wind interaction shocks, one exterior to the He i 1.0830 μm line formation region and a second exterior to the dust debris. Forbidden lines are present in the 1 μm spectral region. [C i] 9850 Å is spatially resolved, of similar extent to the He i 1.0830 μm line. Following the discussion of van Hoof et al. (2007) the extended [C i] 9850 Å emission originates from recombination following a shock. The forbidden lines also have a spatially unresolved component. Bipolar proto-PNe have a similar photoionization spectrum arising from a 10,000 K region near the central object. The observations reported here provide limits to the distance to V4334 Sgr. The angular size of the debris cloud is ∼20%–25% that of the gas cloud. The helium expansion velocity is in the range 500–700 km s−1 . Thus the dust expansion velocity is 100–175 km s−1 . The change of angular size of the dust shell is 0.055 mas d−1 which is equal to the expansion velocity divided by the distance (Equation (2)). Hence the distance is in the range 2.1–3.7 kpc. Previous distances discussed in the literature converge to 2–4 kpc. A distance of 2.8 kpc results in the same dust expansion velocity for V605 Aql and V4334 Sgr. This suggests the same ejection process for both objects although measurements suggest a dust mass for V4334 Sgr 10 times larger than that of V605 Aql. There are multiple indicators showing that the cloud of ejecta is asymmetric. These include the bipolar distribution of dust with debris clouds visible at opposed sides of the central nebula. The He i 1.0830 μm spectrum shows an asymmetric flow toward the observer. This was previously thought to result from the backside being blocked by dust but the very different dust and gas expansion rates do not support this explanation. The spatialspectral images of both He i 1.0830 μm and [C i] 9850 Å show spatially extended lines that are spatially asymmetric. Mid-IR imaging reveals a compact disk (Chesneau et al. 2009). While the old PN and the abundances of the debris require VLTP, the asymmetric nebula appears not to be in accord with single star evolution and suggests that V4334 Sgr, as previously suggested for V605 Aql (Lau et al. 2011), may result from a more complex evolution path involving a companion. Considerable change has taken place in the morphology of the images over the 2.5 yr interval between the observations. The mass of the ejecta is equal to that of the initial hydrogen-rich envelope. This implies that the remnant has been totally stripped of the envelope. Thus as the debris cloud expands the stellar interior intershell region will be exposed as a [WR] object as it has for the older final flash objects.

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