EVOLUTION OF THE 1919 EJECTA OF V605 AQUILAE^*,†

In the second decade of the 20th century, the variable star V605 Aquilae underwent a remarkable outburst (Clayton & De Marco 1997; Clayton et al. 2006). Its eruption was discovered by Wolf (1920) on plates taken at Heidelberg, with the star reaching a maximum photographic magnitude of 10.2 in 1919 August. Subsequent investigation of Harvard plates by Ida Woods showed that V605 Aql had slowly risen to its maximum over the previous 2 yr (Bailey 1921). The historical visual and photographic photometry has been assembled and discussed by Duerbeck (2002; see also Harrison 1996). Although initially considered to be a very slow nova—it was designated Nova Aquilae No. 4—V605 Aql had a light curve unlike that of any known nova. After fading in 1920, it re-brightened in 1921, faded again, and then brightened yet again in 1923, before finally fading from sight. Fuhrmann (1981) examined over 400 plates obtained between 1928 and 1979, and found that V605 Aql was always below photographic magnitude, 16–17.5, during this interval. Spectroscopy obtained in 1921 September, showed the object to be a cool carbon star, with a spectral type of R0, in stark contrast to the high-excitation emission-line spectra of classical novae at late outburst stages (Lundmark 1921).

In spite of these unique features, V605 Aql was almost entirely ignored by astronomers over the ensuing five decades. Modern interest was rekindled by Bidelman (1971), who reviewed the scanty data and made the novel suggestion that V605 Aql had been an extragalactic supernova (SN). His speculation was based on Lundmark's R0 classification along with the fact that the first SN whose spectrum was photographed, Z Centauri (SN 1895B), had a spectrum that had been considered to resemble that of an R-type star by Cannon & Pickering (1916). However, Bidelman (1973) subsequently was able to examine Lundmark's original 1921 spectrograms, and verified that V605 Aql in outburst indeed was a carbon star—in fact, a hydrogen-deficient carbon (HdC) star, closely similar to the well-known non-variable HdC star HD 182040, as well as other cool R Coronae Borealis (RCB) stars (Clayton 2012). The 1921 spectrum was re-examined and presented in Clayton & De Marco (1997).

Following up on Bidelman (1971), van den Bergh (1971) obtained deep photographs of V605 Aql with the 5 m Hale telescope. These images showed it to be the central star of a large, very faint, old planetary nebula (PN), which had already been cataloged as object 58 in the Abell (1966) list of faint PNe discovered in the course of the Palomar Observatory Sky Survey (POSS). A58 (PN G037.5−05.1) has an elliptical shape, with dimensions of 44'' × 36''. See Figure 1.

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Based on these findings—a cool, hydrogen-deficient star erupting in the center of an old PN—many investigators believe V605 Aql to be a star that underwent a final helium-shell flash (FF), an event first described theoretically by Iben et al. (1983). In the FF scenario, a hot PN nucleus (PNN) re-ignites He shell burning, retraces its evolution to become temporarily a "born-again" red giant whose remaining surface hydrogen has been ingested and burned, and then evolves back to its original high surface temperature. Recent theoretical models indicate that the transition from a hot PNN to a cool giant, and then back again, can occur in as little as 5–10 yr. Such extremely rapid evolution results from a very late thermal pulse, in which the star experiences a FF only when it has already reached the top of the white-dwarf cooling track with a very thin envelope remaining (Herwig 2001; Lawlor & MacDonald 2003; Hajduk et al. 2005).

V605 Aql in 1921 had a spectrum resembling that of a cool RCB star with T_eff ≃ 5000 K (Clayton & De Marco 1997). However, there is evidence that it had reheated no later than the mid-1970s, and likely as early as the early 1950s, if not even earlier. At present, a dusty knot of extremely hydrogen-deficient nebulosity, unresolved or marginally resolved in ground-based optical images, lies at the site of V605 Aql at the center of the large, faint PN, A58 (Clayton & De Marco 1997; Clayton et al. 2006). The spectrum of the compact knot is extraordinary. It shows strong emission lines of [O iii], but no sign of the Balmer lines (Seitter 1985). The surrounding PN, however, has a fairly normal hydrogen-rich composition (e.g., Ford 1971; Wesson et al. 2008). Optical spectra of the knot, obtained in 2001, confirm the presence of broad, high-excitation Wolf–Rayet stellar features, including strong C iv emission, and clearly demonstrate that the star is again very hot, ∼95,000 K (Clayton et al. 2006). Since the knot is visible in blue POSS images obtained in 1951, V605 Aql was most likely already producing ionizing radiation and strong [O iii] emission by that time. A spectrogram obtained in the mid-1970s by P. Osmer with the Cerro Tololo 4 m telescope at the request of H.E.B. showed the nebular emission lines, but unfortunately the unpublished data have now been lost. The compact knot is probably material ejected in the 1919 outburst, which is so dusty that it hides the star from direct view; possibly we see only starlight that emerges after scattering off the dust (Clayton et al. 2006).

The behavior of V605 Aql is very similar to that seen more recently in Sakurai's object (V4334 Sgr), which also lies within a very faint PN. After a rapid brightening in 1996, followed by several fluctuations, Sakurai's object went into a deep decline in 1999, just as V605 Aql did in 1923 (Kerber et al. 1999; Duerbeck 2002). Thick dust formed around both stars, causing their precipitous optical declines. The spectrum of Sakurai's object in 1997 was very similar to the 1921 spectrum of V605 Aql (Kerber 2001). Also, Sakurai's object shows evidence for new circumstellar ionized gas only 10 yr after its outburst (Kerber et al. 2002; Hajduk et al. 2005), suggesting that it is already becoming hot again. Thus, the current behavior of V605 Aql is a template for the future behavior of Sakurai's object. Another FF object, also surrounded by an old PN, FG Sge, is evolving much more slowly than the near-twins V605 Aql and Sakurai's object (Duerbeck 2002; Lechner & Kimeswenger 2004).

Recently, however, the FF scenario for V605 Aql has been questioned by Wesson et al. (2008) and Lau et al. (2011). They carried out a spectroscopic analysis of the V605 Aql central knot, showing a C/O ratio below unity and a remarkably high abundance of Ne, properties that are in stark conflict with theoretical models of FF events. These findings suggest a more complicated evolutionary history, involving a ONe white dwarf involved either in a binary merger or in an unusual classical-nova outburst.

In this paper, we present new polarimetric high-resolution imaging of V605 Aql with the Hubble Space Telescope (HST). The polarimetry was done to investigate whether the recent spectrum of V605 Aql is seen only in scattered light, in which case the knot surrounding the star should be highly polarized. These new images also provide an 18 yr baseline to measure the expansion and evolution of the central knot. We examine the morphology and angular expansion of the compact nebula, estimate its distance and reddening, while discussing the implications for the nature of V605 Aql.

2.1. HST

2.2. Infrared Photometry

The spectral energy distribution (SED) of V605 Aql shows that there is significant dust around the star emitting in the mid- to far-IR. The SED is plotted in Figure 2. The data, summarized in Table 2, consist of JHK photometry (Hinkle et al. 2008), as well as photometry from the Infrared Astronomy Satellite (IRAS; Walker 1985; Rao & Nandy 1986); the AKARI Infrared Camera and Far-Infrared Surveyor (FIS; Murakami et al. 2007; Ishihara et al. 2010); the Wide-field Infrared Survey Explorer (WISE; Wright et al. 2010); and the Spitzer Multiband Imaging Photometer for Spitzer (MIPS; Rieke et al. 2004). MIPS scans at 24, 70, and 160 μm of V605 Aql were downloaded from the archive and photometry was done using point spread function (PSF)-fitting in Starfinder. The other IR photometry of V605 Aql was collected from the literature and from VizieR.

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Table 2. V605 Aql Photometry

Band	Flux
	(Jy)
Ja	3.21e−05 ± 1.5e−05
Ha	2.27e−04 ± 3.0e−05
Ka	8.56e−04 ± 8.0e−05
IRAS/12	4.99e+00 ± 2.0e−01
IRAS/25	2.95e+01 ± 1.18e+00
IRAS/60	4.07e+01 ± 4.1e+00
IRAS/100	1.83e+01 ± 2.0e+00
AKARI/9	2.85e+00 ± 1.68e−02
AKARI/18	1.53e+01 ± 1.20e−01
AKARI/65	2.67e+01 ± 2.50e+00
AKARI/90	2.08e+01 ± 9.87e−01
AKARI/140	6.88e+00 ± 9.03e−01
WISE/3.4	1.44e−02 ± 3.05e−04
WISE/4.6	1.13e−01 ± 2.07e−03
WISE/12	8.56e+00 ± 5.52e−02
WISE/22	2.22e+01 ± 1.02e−01
MIPS/24	1.57e+01 ± 8.01e−02
MIPS/70	1.78e+01 ± 1.85e−01
MIPS/160	2.82e+00 ± 1.12e−01

Note. ^aJHK photometry from Hinkle et al. (2008).

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2.3. MIDI/VLTI

In Figure 3 (top panel), we show the sum of the four polarimetric HST images, taken in the intermediate-width F547M filter, which includes the strong stellar-wind C iv emission feature. This is a very deep image, combining six HST orbits of imaging. Figure 3 is dominated by a slightly non-stellar "hot spot." The FWHM of this bright feature is 033, as compared to 018 for nearby stars in the field. We suggest that this is the central star itself, but still seen through substantial extinction. The surrounding nebula is much fainter; in this filter the nebula is mostly starlight scattered off dust plus a small amount of [O iii] 5007 Å emission leaking in from the short-wavelength tail of the F547M bandpass.

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To illustrate the spatial relationships between the different features, we show in Figure 3 (bottom panel) a false-color image made by registering and combining the [O iii], F547M, and [N ii] frames from 2009. This image demonstrates that our putative central star lies near the geometric center of the compact nebula, including the faint extension to the southwest. Note that the near-IR JHK images presented by Hinkle et al. (2008) show the same nearly stellar hot spot, which they argued is the central star.

Three of the four polarimetric F547M images were summed by combining the various polarization angles into a Stokes I image using linear combinations given by Biretta & McMaster (1997). A "standard" star was chosen in the field which is listed in the GSC2.2 and UCAC4 catalogs as V = 15.57, and in NOMAD at 15.4. We assumed a magnitude of V = 15.5 for the standard star. Photometry was done on this star and the central star of V605 Aql in the F547M image, using standard techniques in DAOPHOT and PSF-fitting. We estimate that the V605 Aql central star is V = 20.2 ± 1.0. The uncertainty is an estimate based on the quality of the profile fitting and the uncertainty in the standard star's magnitude, but does not try to estimate uncertainties introduced by using a medium-pass filter with no color information. If the reddening estimate from Section 5 is used, then the intrinsic apparent brightness of V605 Aql in 2009 was V = 20.2–4 ∼ 16.2. The F547M bandpass is dominated by the C iv Wolf–Rayet emission feature in the spectrum of the central star (Clayton et al. 2006). Therefore, we cannot use the photometry in this filter as an apparent stellar V magnitude which could be used with the estimated distance to calculate M_V.

Figure 4 shows the [O iii] images at the 1991, 2001, and 2009 epochs, and the [N ii] images from 2001 and 2009. The brighter portions of the compact nebula in 2009 have an angular diameter of ∼1''; the total diameter is about 17 if we include the faint extensions on the southwest side. Clayton et al. (2006) suggested that V605 Aql may be surrounded by a thick torus of dust with some stellar light scattering toward the observer from the poles of the torus. This idea was discussed in detail by Hinkle et al. (2008), based on their 2001 HST and ground-based images. This model was suggested by the earlier epochs of imaging which showed two blobs of emission separated by a dark band, thought to be an almost edge-on thick disk of dust obscuring the central star in the visible. Guerrero & Manchado (1996) first noted the asymmetry of the central knot, and the dark band was already seen in the aberrated FOC images from 1991 (Figure 4, top left panel), and then more clearly in the deconvolved versions presented by Bond & Pollacco (2002). Hinkle et al. (2008) suggested that the central star was visible in their near-IR imaging and that the disk was inclined enough to see the star in the center of the disk. This morphology is suggestive of a bipolar structure possibly with a central disk. The orientation of the dark band (P.A. ∼ 150°) is also approximately parallel to the direction of the major-axis of the surrounding PN, as shown in Figure 1.

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Another interpretation of the morphology is that the northeast and southwest emission is coming from material expanding away from the star and its torus or disk at the center. This geometry is shown in Figure 5. It is possible that the dust obscuration of the V605 Aql central star and the nebula is decreasing with time so that in the latest 2009 imaging, the central star is visible for the first time and the southwest component is seen more clearly. The northeast component is on the near side and less obscured than the southwest component which is on the far side of the star. There is a gradient in the ratio of [O iii] to [N ii] in the nebula, in the sense that the ratio is highest on the northeast side and lowest on the southwest. This could represent a range in dust extinction, lower on the northeast, near-side, and higher on the southwest, far-side of the star. The gradient could also be due to a range in excitation level.

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High-resolution (0.35 Å) spectra of the integrated light of the compact nebula were obtained in mid-1991 by Pollacco et al. (1992), using the coudé echelle spectrograph on the Anglo-Australian Telescope. These spectra showed a system of broad and narrow emission-line components in [O iii] and [N ii], as well as Hα from the surrounding PN, A58. Unfortunately, the original data have now been lost (D. Pollacco 2012, private communication), so we reconstructed them by digitizing the figures in the online version of the journal article and converting them to tables of intensity versus wavelength.

Pollacco et al. (1992) noted that Hα and [N ii] 6548–6583 Å have narrow emission-line components with double profiles. For some unknown reason, the [O iii] lines are not split. Attributing these components to the surrounding PN, they derived an expansion velocity of 31 ± 4  km s⁻¹ and a systemic velocity for the large PN of +70  km s⁻¹. The spectra also showed much broader emission features in [N ii] and [O iii] (but not in Hα), attributed to the compact hydrogen-deficient nebula. The peaks of the broad features were blueshifted by ∼100 km s⁻¹ with respect to the systemic velocity, and with FWHMs of 180 and 270 km s⁻¹, respectively. Pollacco et al. interpreted these blueshifted features as arising either from a one-sided collimated flow, or more probably from a dusty, approximately spherical wind, in which the receding (i.e., redshifted) material is obscured by dust within the wind. The latter scenario is bolstered by a subtle indication of a red tail in the broad components, and is the one that we adopt here.

We use a simple model of a spherical nebula of uniform density with radius R, expanding at a constant speed v_exp, and a spherical r, θ coordinate system, where θ is the angle relative to the line of sight. A volume element located at (r, θ) will have radial velocity v_r = −v_expcos θ. If the optical depth to the center of the sphere is τ, the intensity from that volume element will be extinguished by exp (− τ[Z/R]), where $Z=\sqrt{R^2-(r\sin \theta)^2}-r\cos \theta$ is the line-of-sight distance out of the nebula from the volume element toward the observer.

In this model, the emission-line profile, I(v_r), is an integral of extinguished emission over the entire sphere. An asymmetric profile can thus be fit to the observed spectra using standard Levenburg–Marquardt minimization (Press et al. 1992) of the numeric integral as a function of systemic velocity, line width, and the optical depth to the nebula center, τ. We performed such fits on the [O iii] 5007Å and [N ii] 6548–6583 Å line complexes, as shown in Figure 6. The doubled narrow components were forced to have no internal extinction, as they represent emission from the outer PN, A58. The narrow lines yield a mean systemic velocity of +80 km s⁻¹, very close to that found by Pollacco et al. (1992) from these spectra. The best fit model is also shown in Figure 6.

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No single asymmetric profile can account for all of the broad emission; however, we find that the bulk of the emission, including the broad red and blue tails, can be well fit by a single sphere at a systemic velocity of +80 km s⁻¹, and an expansion velocity of 215 km s⁻¹, with τ = 4.0 for [O iii] and 3.0 for [N ii]. These components of the line profiles are shown in green in Figure 6. The ratio of optical depths is that expected for pure carbon dust following a Mathis et al. (1977, hereafter MRN) grain-size distribution, and corresponds to an extinction toward the central star of A_V = 4 mag. Note that this model does not account for profile broadening due to thermal motions, slit width and the line-spread function of the optical system. To fully reproduce the observed profile, we would have to add or subtract some narrow features, which suggests that the outflow has a clumpy substructure, consistent with the narrow-line images presented in Figure 4. The Galactic latitude of V605 Aql is −5° which implies a foreground interstellar extinction of A_V ∼ 1.45–1.75 mag (Schlegel et al. 1998; Schlafly & Finkbeiner 2011). So the total extinction toward V605 Aql is ∼5.5 mag. The uncertainty in this estimate is substantial as the extinction estimates from both the emission line fits and the foreground component contain simplifying assumptions.

A primary goal of this investigation was to measure the angular expansion of the compact nebula. Visual inspection of the [O iii] and [N ii] images in Figure 4 does confirm noticeable expansion.

We chose the [O iii] images from 2001 and 2009 (separated by 7.8 yr) for detailed analysis. As described in Section 2, we resampled the 2009 image to the (lower) resolution of the 2001 image. We then magnified the earlier image by various factors and subtracted it from the later one, and sought to minimize the residual fluxes. Figure 7 shows PSF-matched difference images for magnifications of 0%, 8%, and 16% between the 2001 and 2009 epochs. The 0% and 16% difference have much larger residuals due to the expansion of the nebula, and the 8% image is the best fit.

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The magnification residuals were minimized as follows. Assuming a homologous expansion by a magnification factor a, and denoting the reference and input images by R and I, and the magnified input image by aI, we calculate the sum of the squared residuals, ρ(a) = ∑_pix(R − aI)². The most likely expansion factor is then the value a_min that minimizes ρ.

The uncertainty can be directly computed as $\sigma =\sqrt{\langle a^2 \rangle - \langle a \rangle ^2}$, where 〈x〉 is the expectation value of the parameter x, i.e., 〈a〉 = ∫aρ(a)da, if ρ has been normalized. When nebulae are better resolved, one can perform this task across various wedges or quadrants in position angle, in order to determine the local average expansion, or one can compare the expansion of individual features to measure the microscopic expansion properties (Li et al. 2002; Ueta et al. 2006). However, since the V605 Aql nebula is only marginally resolved, we opted for a global measurement in which the entire nebula is assumed to expand homologously.

As a note of caution, when the input and reference images have PSFs that differ significantly, the residual ρ may not be minimized for the actual best-fit magnification factor, even if that factor is exactly known ahead of time. We attempted to minimize this problem by PSF-matching the WFPC2 images before computing the difference residuals, and by masking out pixels surrounding the bright, quasi-stellar source at the center of the nebula.

Figure 8 plots the normalized residuals for the [O iii] filter against magnification factor, for the 2001–2009 pair (solid line) and the 1991–2009 pair (dot-dash line). For the 2001–2009 WFPC2 images, the derived expansion factor is 8% ± 2%. We verified this numerical result in a "double-blind" fashion, by visually examining residual images produced by scaling and subtracting the two epochs, and selecting the magnification factors that appeared best, without knowledge of the results plotted in Figure 8. The flattest-appearing residuals were indeed close to the 8% value. Not shown here is a similar exercise for the 2001–2009 [N ii] images, which gave an expansion factor of 5% ± 3%.

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Under the assumption of homologous expansion, the 2001–2009 [O iii] data yield an ejection epoch of 1903 ± 49. Performing the same exercise between the 2009 WFPC2 image and the FOC data from 1991 yields an expansion factor of 27% ± 9%, as shown in Figure 8, and an ejection epoch of 1926 ± 22. Both results are thus consistent, within the errors, with an ejection of the compact nebula at the time of the 1919 optical outburst.

We estimated the distance to V605 Aql using the expansion-parallax method. In this technique, the distance is determined by equating the velocity implied by the angular expansion rate in the plane of the sky, derived in the previous two sections, to the expansion velocity in the line of sight.

In order to estimate a distance to V605 Aql based on the expansion parallax, we need to convert the fractional expansion rate of the compact nebula into angular units. We measured this by considering each of the four quadrants of the image, and determining the radial pixel shift needed to match features between epochs. We found that a radial shift of 008 ± 001 minimizes the subtraction residuals.

We also used a "by-eye" approach, in which we estimated the location of the outer edge of the nebula from high-contrast renditions of the images. This visual method also yielded an estimated expansion of 008 over the 7.8 yr epoch difference. The angular expansion rate is thus 10.2 ± 1.3 mas yr⁻¹.

A rough geometric distance, d, can be estimated from the observed angular radius, a ≃ 085 at epoch t = 2009, as discussed in Section 4, combined with the expansion velocity of v_exp = 215 ± 20  km s⁻¹. We assume a simple spherical expansion (in spite of the obviously non-spherical appearance of the compact nebula), and that it was ejected in 1919. The result is

$$d \simeq 4.8 \, {\rm kpc} \left({0{\buildrel{\prime\prime}\over{.}} 85}\over {a}\right) \left({v_{\rm exp}\vphantom{d}}\over {215 {\:\rm km\:s^{-1}}\vphantom{p}}\right) \left({t-1919\vphantom{d}}\over {2009-1919\vphantom{p}}\right),$$

but this is somewhat uncertain, since the nebular radius is poorly defined.

An alternative calculation, using instead the angular expansion rate of $\dot{\theta }=10.2\pm 1.3\,\rm mas\,yr^{-1}$, yields

$$d \simeq 4.4 \, {\rm kpc} \left({v_{\rm exp}\vphantom{d}}\over {215 {\:\rm km\:s^{-1}}\vphantom{p}}\right) \left({10.2\,\rm mas \, yr^{-1}}\over {\dot{\theta }}\right),$$

with a formal error of about ±0.6 kpc. There is likely a significant systematic error because of our assumption of a simple spherical expansion. The images of the V605 Aql knot clearly show that it is not spherical so the expansion velocity in the plane of the sky may be larger than in the radial direction, in which case our distance estimates would be too low.

Other distance estimates, based on various methods applied to the large PN A58, have been reviewed by Clayton & De Marco (1997) and Hinkle et al. (2008), and range from 2.7 to about 6 kpc. The result estimated here, ∼4.6 kpc, is consistent with the previous estimates, given the large uncertainties and systematic errors in all of the methods.

Koller & Kimeswenger (2001) modeled the dust emission from V605 Aql using a radiative transfer code which fit ISO and IRAS photometry of the star. They assumed that d = 5 kpc, T_eff = 100,000 K, and L_⋆ = 5800 L_☉. They also assumed carbon dust with A_V = 7 and found a dust mass of 8 × 10⁻³ M_☉. Hinkle et al. (2008) fit IR photometry from the J-band to 14.5 μm with two dust components emitting as blackbodies with temperatures of 350 and 1500 K, but do not estimate the total dust mass.

Wesson et al. (2008) calculated a gas mass for the ejected knot of 5 × 10⁻⁵ M_☉ assuming a distance of 3.5 kpc, n_e = 2100 cm⁻³, and an angular radius of 038. The angular radius corresponds to 2 × 10¹⁶ cm at 3.5 kpc. This agrees well with the size of a nebula expanding at 215 km s⁻¹ for 90 yr (6 × 10¹⁶ cm). They estimate the extinction from c(Hβ) = 2.0 for the inner knot which corresponds to A_V ∼ 3 mag.

The SED of V605 Aql, shown in Figure 2, was fit using various dust components at different temperatures. As described in Sugerman et al. (2012), a zeroth-order estimate of the emitting dust mass can be produced by fits to the observed SED. Each dust component has a single temperature, and emits as a blackbody modified by the mass-absorption coefficient. In addition, each was assumed to be optically thin, 100% amorphous carbon with an MRN distribution, and optical constants from Weingartner & Draine (2001). Figure 2 shows the best fit which required dust with three different temperatures and masses (810 K, 1 × 10⁻¹¹ M_☉; 235 K, 9 × 10⁻⁶ M_☉; 75 K, 2 × 10⁻³ M_☉). There is a range of models that fit the data roughly equivalently, with masses that are consistent to within a factor of two of those presented here. Also, these dust masses are lower limits to the values one would find when running a Monte Carlo radiative-transfer model.

We have estimated in Section 5 that V605 Aql suffers from ∼4 mag of internal extinction at V. Let's assume a spherical shell of dust with a radius of 2 × 10¹⁶ cm around the star as did Wesson et al. (2008). Then, looking through the dust to the star, the optical depth is, τ = κρL, where ρ is the dust mass density, κ is the mass absorption coefficient, and L is the path length. V605 Aql is carbon rich (Clayton et al. 2006), and the Spitzer/IRS spectrum of shows no silicate or polycyclic aromatic hydrocarbon (PAH) features so the dust is likely to be amorphous carbon (Evans et al. 2006). A typical value for the absorption coefficient is κ ∼ 5 × 10⁵ cm² g⁻¹ (Zubko et al. 1996; Draine & Li 2007). If L = 2 × 10¹⁶ cm and τ = 4, then ρ = 4 × 10⁻²² cm² g⁻¹. Then, the mass of dust in the shell is just 4/3 πL³ρ or ∼7 × 10⁻⁶ M_☉.

This implies a gas-to-dust ratio of ∼7 which is very small compared to the typical gas-to-dust-ratio for the Galaxy of 100 (Draine et al. 2007). However, the gas-phase abundances in the central knot of V605 Aql are very different than those seen in a typical interstellar cloud. Wesson et al. (2008) measured the fractional abundances of various elements in the V605 knot, H (0.019), He (0.250), C (0.021), N (0.043), O (0.323), and Ne (0.345). Firstly, H, which is typically 75% of interstellar gas, and does not deplete into dust grains, is missing, which will lower the gas-to-dust ratio significantly. Among the other abundant elements in the central knot, Ne and He will also be undepleted as they do not participate in the grain chemistry. In the interstellar medium (ISM), C and O are heavily depleted into carbon grains and silicate grains. However, the MIDI, ISO, and Spitzer Infrared Spectrograph (IRS; Evans et al. 2006) observations of V605 Aql indicate that its mid-IR spectrum is a featureless continuum. This is similar to what is seen in the RCB stars (Kraemer et al. 2005). There are no silicate features at 10 and 20 μm, and no sign of PAHs or SiC. The featureless spectrum indicates that the dust around V605 is probably amorphous carbon, i.e., 100% carbon. So, it is likely that carbon is highly depleted into grains, but that the other elements, measured by Wesson et al. (2008), are not. If carbon is depleted at a level that we see in the ISM, then probably 90% of the carbon is in dust. In the gas phase in the V605 knot, carbon is ∼1 × 10⁻⁶ M_☉. If carbon is 90% depleted, then the total mass of carbon (gas phase and in dust) in the knot is 1 × 10⁻⁵ M_☉. Then the gas-to-dust ratio is ∼5, which is remarkably low. This estimate for the dust mass (1 × 10⁻⁵ M_☉) fits in very well with the estimates above from the extinction and likely size of the shell (7 × 10⁻⁶ M_☉), and from the hot and warm dust components of the SED fit (9 × 10⁻⁶ M_☉).

Assuming that the mass of C is actually 1 × 10⁻⁵ M_☉, increases the mass of the central knot to 5.9 × 10⁻⁵ M_☉, and changes the ratio C/O from 0.06 to 0.6 but still well below unity and far below the large C/O predicted for an FF.

Wesson's estimate of the gas mass measures only hot gas near the star. The second of the mass estimates in Wesson et al. (2008) used a density of 2100 cm⁻³, derived from collisionally excited lines, so that could underestimate the gas mass if there were regions of very high density to quench the collisionally excited lines (CEL) emission. Also, the apertures for the near- and mid-IR photometry, plotted in Figure 2, have sizes of a few arcseconds and are measuring the hot gas and dust in the central knot. But, the apertures for the far-IR (>20 μm) photometry are much larger, IRAS (spatial resolution ∼4') and AKARI/FIS (spatial resolution ∼35''), and may be including colder dust further from the star, possibly associated with the PN which is 44'' × 35'' in size (Abell 1966). This idea is supported by the fact that the flux density in the Spitzer/MIPS images, which have a better spatial resolution, is slightly lower. So, the value estimated above, ∼1 × 10⁻⁵ M_☉, is probably consistent with the dust mass in the central knot. The large mass of cold dust, 2 × 10⁻³ M_☉, estimated from the SED fit, likely lies outside the knot. A significant mass of cold dust was also recently found around R CrB, which is a possible FF star (Clayton et al. 2011).

New polarimetric imaging with HST/WFPC2 in nebular [O iii], [N ii], and stellar C iv provides a baseline of 18 yr to investigate the expansion of the ejecta from the 1919 event, and allow us to investigate the position and extinction of the central star. The new measurements of the rate of the expansion, reported here, are consistent with an ejection in 1919. When combined with velocities measured from spectra of V605 Aql, a distance of 4.6 kpc is estimated, which is in the middle of the range of previous distance estimates.

The morphology of the 1919 ejecta is highly non-spherical. It seems to be either a tilted dusty torus with a dark band, or a disk surrounding the central star with bipolar ejections on either side. Further imaging as the expansion continues may clarify this. The asymmetry of the knot is aligned with the asymmetry of the much larger surrounding PN. A disk and PN asymmetry could result if V605 Aql is a binary system (Lau et al. 2011). There is early evidence that there is a disk around Sakurai's object (Chesneau et al. 2009).

The new HST images of V605 Aql, presented here, provide the first direct detection of the central star in the visible since 1923. In the stacked F547M image, the star is V = 20.2 ± 1.0 mag. We estimate that the intrinsic apparent magnitude today should be V ∼ 16.2 mag, assuming the extinction is A_V ∼ 4 mag which is estimated from the spectral modeling. This modeling estimates a systemic velocity of +80 km s⁻¹ and an expansion velocity of 215 km s⁻¹. Polarimetric imaging was obtained to investigate whether the 2001 spectrum of V605 Aql (Clayton et al. 2006) was obtained primarily in scattered light from dust in the central knot, but the signal-to-noise in the data was not sufficient to measure the level of polarization.

It is found that the SED of the dust around V605 Aql, consisting of Near-IR, IRAS, WISE, AKARI, and Spitzer photometry, can be fit using carbon dust with three different temperatures and masses. Using three different methods, we estimate a mass of warm dust in the central knot of ∼10⁻⁵ M_☉. A significant mass of cold dust (10⁻³ M_☉) may be spread over a larger volume perhaps associated with the surrounding PN.

V605 Aql is one of only three stars, identified with a final helium-shell flash, to be observed from their initial brightening. It is very similar to the more recent FF star, Sakurai's object. Monitoring of V605 Aql is important because it may allow the future behavior of Sakurai's object to be predicted and better observed. In V605 Aql, the composition of the hydrogen-deficient ejecta (the C/O ratio below unity and a Ne mass fraction of 35%), however, leaves little doubt that the scenario that formed this object is far more complex than an FF (that predicts C/O far larger than unity and an Ne mass fraction of 2% (Werner & Herwig 2006)). V605 Aql must have passed through a phase involving the dredge-up of Ne from the core of an ONe WD. Another object currently listed in the FF class, A30, also has hydrogen-deficient knots with similar abundance trends (C/O below unity and a Ne mass fraction between 8% and 20% (Wesson et al. 2003)). The ONe nova scenario, proposed by Lau et al. (2011) to explain these abundance patterns, may not be viable because it does not include the compulsory series of carbon–oxygen nova outbursts that precede the ONe nova detonation. If so, then the only viable scenario is that of a merger between a massive ONe white dwarf and a companion. At the present time, it is not possible to resolve whether V605 Aql is the result of a final helium-shell flash, an ONe nova, or a merger event.

We thank the anonymous referee for several helpful suggestions. Support for this work was provided by NASA through grant number HST-GO-11985 from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555; and by the Goucher College Office of the Provost. Some of the data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST). We would like to acknowledge James Liebert, Alvio Renzini, and Nye Evans for assisting with this work. This research has made use of the VizieR catalog access tool, CDS, Strasbourg, France.

Facility: HST (WFPC2, FOS, FOC) - Hubble Space Telescope satellite