VLT/UVES spectroscopy of V4332 Sagittarii in 2005: The best view on a decade-old stellar-merger remnant
aa r X i v : . [ a s t r o - ph . S R ] A p r Astronomy & Astrophysicsmanuscript no. 25592_ap c (cid:13)
ESO 2018July 29, 2018
VLT/UVES spectroscopy of V4332 Sagittarii in 2005:The best view on a decade-old stellar-merger remnant
R. Tylenda , S. K. Górny , T. Kami´nski , and M. Schmidt Department of Astrophysics, Nicolaus Copernicus Astronomical Center, Rabia´nska 8, 87-100 Toru´n, Polande-mail: [email protected] European Southern Observatory, Alonso de Córdova 3107, Vitacura, Santiago, ChileReceived; accepted
ABSTRACT
Context.
The source V4332 Sgr is a red transient (red nova) whose eruption was observed in 1994. The remnant of the eruptionshows a unique optical spectrum: strong emission lines of atoms and molecules superimposed on an M-type stellar spectrum. Thestellar-like remnant is not directly observable, however. It is presumably embedded in a disc-like dusty envelope seen almost face-on.The observed optical spectrum is assumed to result from interactions of the central-star radiation with dust and gas in the disc andoutflows initiated in 1994.
Aims.
We aim at studying the optical spectrum of the object in great detail to better understand the origin of the spectrum and thenature of the object.
Methods.
We reduced and measured a high-resolution (R ≃
40 000) spectrum of V4332 Sgr obtained with VLT / UVES in April / May2005. The spectrum comes from the ESO archives and is the best quality spectrum of the object ever obtained.
Results.
We identified and measured over 200 emission features belonging to 11 elements and 6 molecules. The continuous, stellar-like component can be classified as ∼ M3. The radial velocity of the object is ∼ −
75 km s − as derived from narrow atomic emissionlines. The interstellar reddening was estimated to be 0 . ≤ E B − V ≤ .
75. From radial velocities of interstellar absorption features inthe Na i D lines, we estimated a lower limit of ∼ ∼ V band), which resulted from cooling of the main remnant by 300–350 K, corresponding to its spectral-typechange from M3 to M5-6. The object increased in luminosity by ∼ Key words. stars: individual: V4332 Sgr - stars: emission-line - stars: late-type - stars: activity - stars: winds, outflows - stars:variables: other
1. Introduction
The source V4332 Sagittarii (V4332 Sgr) was discovered asNova Sgr 1994 in February 1994 (Hayashi et al. 1994). Spec-troscopic observations showed narrow Balmer lines in emisionsuperimposed on a K-type (super)giant spectrum, however, thatquickly evolved to M type (super)giant (Tomaney et al. 1994;Martini et al. 1999). This spectral apearance and evolution wasat variance with what is observed in classical novae, but borea resemblance to the luminous red variable observed in M31 in1988 (M31 RV) (Mould et al. 1990). At present, stellar eruptionsof this type are called red transients, intermediate-luminosityred transients, red novae, luminous red novae, or V838 Mon-type objects. The latter name comes from the gigantic erup-tion of V838 Monocerotis observed in 2002 (Munari et al. 2002;Crause et al. 2003), which elicited great interest in astrophysi-cists, as well as public media, partly due to the spectacular light-echo event accompanying the outburst (Bond et al. 2003).In addition to these three objects, the class of red tran-sients in our Galaxy includes V1309 Scorpii (V1309 Sco)(Mason et al. 2010) and OGLE-2002-BLG-360 (Tylenda et al.2013). V1148 Sagittarii (Nova Sgr 1943) probably also be- longed to this class, as can be inferred from its spectral evolu-tion described by Mayall (1949). There is also growing observa-tional evidence that CK Vulpeculae (CK Vul, Nova Vul 1670)(Shara et al. 1985) was a red transient and not a classicalnova (Kato 2003; Tylenda et al. 2013; Kami´nski et al. 2015a).A few extragalactic objects, usually referred to as intermediate-luminosity optical transients, for instance, M85 OT2006(Kulkarni et al. 2007), NGC300 OT2008 (Bond et al. 2009;Berger et al. 2009), and SN 2008S (Smith et al. 2009), might beof a similar nature.Although they are di ff erent in light curve, time scale, andpeak luminosity, red transients always show a similar spectralevolution: in course of the eruption, the objects evolve to pro-gressively lower e ff ective temperatures and decline as M-type(super)giants. Their remnants also resemble a late M-type (su-per)giants with a significant (often dominating) infrared excess.Several mechanisms have been proposed to explain the red-transient events, including an unusual nova (Iben & Tutukov1992), a late He-shell flash (Lawlor 2005), and a stellar merger(Soker & Tylenda 2003). They have been critically discussed inTylenda & Soker (2006). These authors concluded that the onlymechanism that can satisfactorily account for the observational Article number, page 1 of 14 & Aproofs: manuscript no. 25592_ap data of red transients is a merger of two stars. For the case ofV838 Mon they argued that this eruption might have been due toa merger of a low-mass pre-main-sequence star with an ∼ M ⊙ main-sequence star.The source V1309 Sco, which erupted in 2008 (Mason et al.2010), appeared to be a sort of Rosetta stone for understand-ing the nature of red transients. Thanks to the archive data fromthe Optical Gravitational Lensing Experiment (OGLE) (Udalski2003), it was possible to follow the photometric evolution of theobject during six years before the outburst (Tylenda et al. 2011).The result was amazing: the progenitor of V1309 Sco was a con-tact binary that quickly lost its orbital angular momentum andevolved into a merger of the components. Thus V1309 Sco pro-vided strong evidence that the red transients are indeed causedby stellar mergers.After the 1994 outburst, V4332 Sgr was almost forgottenby astrophysicists and observers. No observations of the objectwere made until 2002, when V838 Mon erupted and astronomersrealized that these two objects most probably belong to the sameclass. Then V4332 Sgr regained astrophysical interest. Severalspectroscopic observations made in 2002-2003 revealed an un-usual optical spectrum for stellar objects: strong and numerousemission lines of atoms and molecules were superimposed on aweak, early-M-type stellar spectrum (Banerjee & Ashok 2004;Tylenda et al. 2005; Kimeswenger 2006). In addition, strongbands of AlO in emission were detected in the near-IR spectralregion (Banerjee et al. 2003).A detailed study of the emission-line spectrum and the spec-tral energy distribution (SED) of V4332 Sgr led Kami´nski et al.(2010) to conclude that the main remnant of the 1994 erup-tion is now obscured to us, most probably the central objectis embedded in a dusty disc seen almost edge-on. The ob-served optical spectrum is assumed to be produced by interac-tions of the central star’s radiation with the matter in the discand the outflows originating from the 1994 eruption: the M-type continuum results from scattering on dust grains, while theemission-line spectrum is due to resonant scattering by atomsand molecules. This conclusion was subsequently confirmed bypolarimetric (Kami´nski & Tylenda 2011) and spectropolarimet-ric (Kami´nski & Tylenda 2013) observations, which showed thatthe optical continuum is strongly polarized, while the emissionfeatures are mostly unpolarized.We present an optical spectrum of V4332 Sgr obtained in2005. The data come from the archives of the Very Large Tele-scope (VLT) and were not published before. The quality and res-olution of the spectrum is exceptional; the data present the bestquality spectrum of V4332 Sgr ever obtained. Since 2005, theobject has significantly faded (see Fig. 1) and its spectrum hasconsiderably evolved (see e.g. Barsukova et al. 2014). Thereforethe spectrum described and analysed in the next sections presentsa unique set of data on V4332 Sgr, which was the main reasonfor us to reduce and publish the data.
2. Observations and data reduction
We have found high resolution spectra of V4332 Sgr in the ESOdata archives. They have been carried out with UVES / VLT in2005 on 22 April and 12 May. The observations were made inthe framework of the 075.D-0511 (PI: Banerjee) observing pro-gramme. The spectra were obtained with three di ff erent spectro-graph settings covering the range 3756–10253 Å with two gapsat 5750–5833 Å and 8520–8656 Å. The technical details of theobservations are provided in Table 1. Fig. 1.
Light curve of V4332 Sgr in the V band. Vertical dashed linesindicate the time moments of the VLT spectroscopy we analysed andthe observations obtained with the Subaru telescope in June 2009 thatare described in Kami´nski et al. (2010). Figure 1 shows the light curve of V4332 Sgr in the V photo-metric band since the discovery of the object in February 1994.The data are taken from the compilation of V. Goranskij (seeBarsukova et al. 2014, for the light curves in di ff erent photomet-ric bands). Two vertical dashed lines indicate the time momentsof the VLT spectroscopic observations we analysed and the ob-servations obtained with the Subaru telescope in June 2009 thatwere presented in Kami´nski et al. (2010).One blue and two di ff erent red standard UVES configura-tions were used to observe the target. The observations centredon 8600Å initially split into two exposures were later merged.Since in the case of the red configurations there are two separatelower and upper parts, the final spectra consist of five indepen-dently registered elements. The CCDs were read with 2x2 bin-ning in all configurations. The total useful spectral range covers3760-9500Å, with two small overlapping regions of individualsections and some gaps between lower and upper parts of the redconfigurations.From the ESO data archive we have downloaded the spectrain their raw, unreduced form together with the full set of reduc-tion and calibration files for each of the standard configurations.The reduction and calibration of the spectra was made with theESO-MIDAS reduction software.The background was subtracted from flat-field, arc, and sci-ence frames. The Th-Ar lamps were used for wavelength cal-ibration. Flat-fielding was performed in the pixel-pixel space.The signs of cosmic rays and other defects of the CCD were re-moved from each science frame using the standard MIDAS pro-cedures. The extracted spectra were wavelength calibrated andcorrected for atmospheric extinction. The flux calibration wasperformed using the master response calibration files preparedby ESO for each of the standard UVES configurations. Tel-luric lines were partially removed using Molecfit (Smette et al. http://jet.sao.ru/ ∼ goray/v4332sgr.ne3 Article number, page 2 of 14ylenda et al.: Optical spectrum of V4332 Sgr F l ux ( − e r g c m − s − Å − ) l (Å) 0 2 4 6 810121416 6000 6500 7000 7500 8000Na I Ca I K I Rb I Rb IScO TiO CrO VO TiO CrO TiO VO TiO TiO VO F l ux ( − e r g c m − s − Å − ) l (Å) Fig. 2.
Spectrum of V4332 Sgr obtained in April–May 2005 with the UVES / VLT. The displayed spectrum was smoothed from the originalresolution with boxcar 10. The strongest atomic (red) and molecular (green) spectral features are indicated. A blue horizontal bar indicates thespectral region a ff ected by telluric absorption bands.
3. Spectrum
The final flux-calibrated 1D spectrum of V4332 Sgr is presentedin Fig. 2. The spectrum above ∼ − erg s − cm − Å − . The strongest atomicand molecular features in emission are indicated in the figure.The atomic line identification was based mainly on the NISTAtomic Spectra Database (Kramida et al. 2013) , the AtomicLine List v2.04 by van Hoof , and on the multiplet tables byMoore (1945). All the identified atomic lines seen in emission in http://physics.nist.gov/asd ∼ peter/atomic/ Article number, page 3 of 14 & Aproofs: manuscript no. 25592_ap
Table 1.
Log of observations of V4332 Sgr with UVES / VLT date & time configuration range(Å) resolution exp. time2005.04.22 06:22:01 RED 580 4727 – 6835 42310 3000 sec2005.05.12 06:10:19 RED 860 6650 – 10250 42310 1480 sec2005.05.12 06:10:23 BLUE 437 3730 – 5000 40970 2960 sec2005.05.12 06:43:16 RED 860 6650 – 10250 42310 1100 secthe spectrum of V4332 Sgr are listed in Table 2. The observedwavelengths (in Å) of the lines are given in Col. (1) of the table,while their laboratory wavelengths and ion identification can befound in Cols. (2) and (3), respectively. Column (4) presents themeasured fluxes (in units of 10 − erg s − cm − ) of the lines withtheir estimated uncertainties. Column (5) gives the full widths athalf maximum (FWHM in Å) of the lines. The last column listsnotes on the lines and their measurements. The explanations ofthe symbols used in Table 2 are given in Table 4.All the identified molecular bands seen in emission in thespectrum of V4332 Sgr are listed in Table 3. The method usedto identify the bands with appropriate references for moleculardata can be found in Kami´nski et al. (2009) and Kami´nski et al.(2010). The first column of the table gives wavelengths of thebands. These are mostly values from laboratory measurementsand refer to the heads of the bands. In a few cases, the wave-lengths are results of theoretical estimates. The identification ofthe bands is provided in Col. (2), while Col. (3) displays themeasured fluxes (in 10 − erg s − cm − ) in the bands with theiraccuracies. The last column gives notes on the bands and theirmeasurements. The symbols used in these two columns are ex-plained in Table 4.
4. Spectral classification of the stellar-likecontinuum
To estimate the spectral type of the stellar-like continuum ob-served in V4332 Sgr, we attempted to fit standard and modelspectra to the results of our observations. Standard stellar spectrawere taken from Jacoby et al. (1984), while model atmospherespectra were obtained using the MARCS grid (Gustafsson et al.2008).Figure 3 shows M2 and M4 giant spectra fitted to the ob-servations. An interstellar reddening was applied to the standardspectra to obtain a good fit at the short- and long-wavelengthedges of the spectra. These were E B − V = . ∼ ∼ T e ff = g = . ∼ Fig. 3.
Standard giant spectra fitted to the observed spectrum ofV4332 Sgr (black). Magenta: M2 standard spectrum reddened with E B − V = .
4. Green: M4 standard spectrum reddened with E B − V = . Note that the two reference spectra shown in Fig. 4 were red-dened with E B − V = .
5. Interstellar reddening
The interstellar extinction towards V4332 Sgr was estimated ina number of papers (Martini et al. 1999; Tylenda et al. 2005;Kimeswenger 2006). These results as well as their own deter-minations were discussed in Kami´nski et al. (2010). Most of theestimates relied on a comparison of the observed continuum ofV4332 Sgr with stellar standards. Kami´nski et al. (2010) finallyadopted E B − V = .
32, although the values varied between 0.22and 0.45. This value agrees well with our analysis in Sect. 4,where the best fit to the observed continuum was obtained withan M3 giant spectrum reddened with E B − V ≃ . E B − V = . Article number, page 4 of 14ylenda et al.: Optical spectrum of V4332 Sgr
Fig. 4.
Standard and model giant spectra fitted to the observed spec-trum of V4332 Sgr (black). Red: M3 standard spectrum reddened with E B − V = .
35. Blue: model atmosphere ( T e ff = g = E B − V = . scattering of the central-star radiation by molecules and atomsin a circumstellar matter. They also concluded that the strongestemission lines, such as Na i i i and K i dou-blets was close to 1:1, while in the optically thin case thisratio should be 2:1. The spectrum of V4332 Sgr we discusshere shows much more numerous and strong emission featuresthan that of Kami´nski et al. (2010). In addition to the Na i andK i doublets, we therefore searched for other emission lines,which are likely to be optically thick. These are the Cr i lines at4254,4275,4290 Å. In the optically thin case, the intensity ratioof these lines should be 1.8:1.4:1.0. As can be seen from Ta-ble 2, the observed ratio is 1.02:0.89:1.0. Almost certainly, theCa i i lines at 4031,4033,4034 Å sug-gests that these lines are almost optically thick, at least the firstone. The theoretical optically thin line ratio for these Mn i linesis 2.0:1.5:1.0. The lines are strongly blended, but from the ob-served profile we can estimate a ratio of 1.4:0.8:1.0.In this way, we have a set of emission lines whose monochro-matic fluxes in the central, most opaque regions of their pro-files are expected to measure the monochromatic flux from thecentral star. These lines span a wide spectral region, 4000–8000 Å, which means that they can be used to estimate the inter-stellar reddening when compared to a reference spectrum thatis assumed to represent the spectrum of the central object ofV4332 Sgr. This is done in Fig. 5. The lower part of the figurepresents essentially the same as Fig. 4 (i.e. the standard M3 andmodel, T e ff = E B − V = .
35 ( R ≡ A V / E B − V = . Fig. 5.
Reference (standard and model giant) spectra compared to theobserved spectrum of V4332 Sgr (black). Red: M3 standard spectrum,green: model atmosphere ( T e ff = g = E B − V = .
35. Reference spectrain the upper plots are reddened with E B − V = .
75 and shifted upward tomatch the highest fluxes of the strongest emission lines. to the observed spectrum) but in the logarithm scale of the flux.The upper plots show the same reference spectra, but shifted up-wards to match the highest fluxes of the strongest, that is, op-tically thick, emission lines, namely Mn i i i i i ffi ce to fit the level of the highestfluxes in the lines. It was also necessary to increase the redden-ing to E B − V ≃ .
75. Thus we can conclude that the emission-linespectrum of V4332 Sgr is significantly more reddened than thestellar-like continuum. As discussed above, the di ff erence in thereddening is most likely due to the bluering of the stellar contin-uum that is scattered on circumstellar dust grains.There is a process, however, that can a ff ect the relative fluxesof the lines we have used to estimate the reddening of the emis-sion spectrum. All the lines we have considered above are fromresonant transitions in atoms. If the lines are optically thick, theresonance photons must su ff er from numerous scattering beforethey escape the emitting region. This increases the chances ofbeing absorbed by dust in the region. Since the absorption coe ffi -cent of dust grains is expected to increase with decreasing wave-length, the lines at 4000–4300 Åwould su ff er more from dust ab-sorption than the lines at 7600–7800 Å. Detailed line transfercalculations are neccesary to investigate this e ff ect, which is be-yond the scope of the present study. The process, if e ff ective,would mimic an additional interstellar extinction. Therefore weconclude from the discussion in this section that the interstellarreddening to V4332 is 0 . < ∼ E B − V < ∼ .
6. Radial velocity
All clear absorption features in the spectrum of V4332 Sgr canbe identified as molecular bands (mostly TiO). They are wide,
Article number, page 5 of 14 & Aproofs: manuscript no. 25592_ap
Table 2.
Components of the interstellar absorption in the profile of theNa i designation λ obs (Å) V r (km s − ) distance (kpc)a 5889.779 − < ∼ −
83 and −
68 km s − .The mean value and standard deviation derived from this sam-ple are − . ± . − . This value can be compared with − ± − obtained in Kami´nski et al. (2010), but that wasan estimate based on strong and wide emission lines.We note that these estimates refer to the line-emitting region,and it is not clear to what extent they measure the radial veloc-ity of the main, stellar-like object of V4332 Sgr. If these valuesmeasure the real radial velocity of V4332 Sgr, then the objectdoes not follow the standard rotation of the Galaxy (see e.g.Brand & Blitz 1993) because the heliocentric radial velocity atany distance is expected to be > ∼ −
10 km s − for the position ofV4332 Sgr.
7. Interstellar lines and distance to V4332 Sgr
There is no reliable estimate of the distance to V4332 Sgr.Martini et al. (1999) proposed 300 pc assuming that the ob-ject was a K-type giant at maximum of the 1994 outburst.Tylenda et al. (2005) obtained 1.8 kpc assuming that the progen-itor was a solar-type main-sequence star. Kimeswenger (2006)estimated 2.9 or 10 kpc depending on whether luminosity classV or III was assumed for the progenitor. Kami´nski et al. (2010)derived a kinematic distance > ∼ i i i lines.Figure 6 displays the observed profiles of the Na i i V r ≃
110 km s − in theNa i V r ≃ −
195 km s − in the Na i Fig. 6.
Profiles of the Na i The observed wavelengths of the interstellar componentsgiven in the second column of Table 2 can be used to calcu-late heliocentric radial velocities, which are listed in the thirdcolumn of the table. These values, after being transformed intoradial velocities in the LSR frame (adding 12.0 km s − for theposition of V4332 Sgr) and adopting the Galactic rotation curveof Brand & Blitz (1993), give estimates of distances of the in-terstellar regions that are responsible for the observed features.The results are listed in the last column of Table 2. From thesedata we can conclude that V4332 Sgr is located at a distance > ∼ b = − ◦ .
40) implies that V4332 Sgr is situated > ∼
8. Evolution of the spectrum of V4332 Sgr between2005 and 2009
Four years after the spectrum described in this paper had beenregistered, a high-resolution spectrum of V4332 Sgr was ob-tained by Kami´nski et al. (2010) using the Subaru telescope. Ascan be seen from Fig. 1, the object in meantime faded by ∼ V band. Figure 7 directly compares the spectra obtainedin these two epochs. Two MARCS model spectra fitted to theobserved spectra are also plotted in the figure. Note that bothmodel spectra were reddened with E B − V = .
35. Note also thatthe model fit to the 2009 spectrum is not perfect in the long-wavelength range. This point, a possible source of extra absorp-tion in the 7300–7500 Å range in particular, was discussed inKami´nski et al. (2010).One of the clear di ff erences between the 2005 and 2009 spec-tra displayed in Fig. 7 is that the 2009 stellar-like continuumis definitely of a later spectral type than that of 2005. Indeed,Kami´nski et al. (2010) classified the 2009 spectrum as M5-6,while we have concluded that it is an M3 type in Sect. 4. This hasobvious consequences on the e ff ective temperature of the central Article number, page 6 of 14ylenda et al.: Optical spectrum of V4332 Sgr
Fig. 7.
Comparison of the spectrum of V4332 Sgr in 2005 (black: thispaper) to that observed in 2009 (red: Kami´nski et al. (2010)). Green:model atmosphere ( T e ff = g = E B − V = .
35 and fitted to the 2009 spectrum. Blue: model atmosphere ( T e ff = g = E B − V = .
35 and fitted to the 2005spectrum. stellar-like object in V4332 Sgr. Our MARCS model fitted tothe observed spectra has T e ff ≃ T e ff ≃ ∼ ff ective ra-dius of the object by a factor of ∼ ∼ BVRI ) and near-IR (
JHKLM )bands. In Fig. 8 we plot the results of photometric measure-ments made in 2003 and compiled in Tylenda et al. (2005) (redfilled symbols), as well as those obtained in 2009 and givenin Kami´nski et al. (2010) (blue open symbols). (There were nonear-IR measurements of V4332 Sgr in 2005, but, as can be seenfrom Fig. 1, the object did not significantly evolve photometri-cally between 2003 and 2005.) The data were fitted with stan-dard stellar photometric spectra supplemented with black-bodydust components in the same way as in Tylenda et al. (2005,for details of the fitting procedure see Tylenda (2005)). The fi-nal fits (sum of the stellar and dust components reddened with E B − V = .
35) are shown with the full curves in Fig. 8. Thefit for the 2003 data (red in the figure) is the same as thosein Tylenda et al. (2005, see their Fig. 5), that is, an M2.7 su-pergiant ( T e ff = T e ff = Fig. 8.
Comparison of the photometric measurements of V4332 Sgrin 2003 (red filled symbols: Tylenda et al. (2005)) to those obtained in2009 (blue open symbols: Kami´nski et al. (2010)). Full curves: modelphotometric spectra obtained by summing a standard stellar componentwith a black-body dust component (individual components are shownwith dotted curves) reddened with E B − V = .
35. See text for more de-tails of the modelling. is, the stellar component by 1.48 times, the dust component by1.51 times.Thus we can quite safely conclude that the optical decline ofV4332 Sgr observed between 2005 and 2007 (see Fig. 1) was dueto a decrease of the e ff ective temperature of the central stellar-like object by 300–350 K. The object expanded during this event,however, and its luminosity increased by ∼ ff erence between thespectra in 2005 and 2009. In the 2009 spectrum the emission fea-tures are significantly fainter than in 2005, not only in absoluteflux scale, but also relative to the M-type continuum. The fadingscale is not the same for all the features, however. The strongestatomic lines, for instance those of Na i and K i , show a similarstrength relative to the local continuum in both epochs. On theother hand, many fainter features that were clearly present in2005 completely disappeared in 2009. To investigate the prob-lem more quantitatively, we calculated relative contributions ofthe emission features to the total flux in selected wavelengthranges in both spectra. For a given wavelength range, we de-rived F obs as an integral of the observed flux over the wave-length and F mod as an integral of the MARCS model flux. Weassumed that the MARCS spectrum, when fitted to the observedspectrum as in Fig. 7, is a good representation of the stellar-like continuum of V4332 Sgr. Then we can define a parameter f em = ( F obs − F mod ) / F obs as a measure of the relative contributionof the emission features to the total observed flux.For the whole spectral range common for both spectra, thatis, for 5500–8000 Å, we derive f em = .
43 for the 2005 VLTspectrum and 0.25 for the 2009 Subaru spectrum. Thus theglobal contribution of the emission spectrum to the total ob-served flux decreased by ∼
40% between 2005 and 2009. Fornarrow spectral ranges encompassing the Na i and K i emission Article number, page 7 of 14 & Aproofs: manuscript no. 25592_ap lines, that is, for 5880–5910 Å and 7650–7720 Å, the resultis f em = .
93 and 0.91 in 2005 compared to f em = .
91 and0.89 in 2009. Thus, as stated above, the strongest emission fea-tures remain practically at the same level when compared to thelocal continuum. For the 6560–6580 Å range, which includesthe Ca i line, the figures are f em = .
60 and 0.49 for 2005and 2009, which is a decrease of 18%. More important, fad-ing a ff ected molecular emission features. For series of the TiO γ bands observed within 6590–6760 Å and 7040–7220 Å we de-rive f em = .
27 and 0.57 in 2005, while for the 2009 spectrumthe figures are 0.14 and 0.30. In other words, these emission fea-tures faded by a factor of 2 relative to the continuum between2005 and 2009. Even greater fading a ff ected the VO-B-X(0-0)bands gathered between 7850–8010 Å. The result is f em = . ff ect.This behaviour can be easily explained within the scenarioproposed by Kami´nski et al. (2010), according to which theemission features in V4332 Sgr are produced by radiative excita-tion of atomic and molecular resonant transitions in a circumstel-lar matter by strong radition from the central stellar-like object,which is invisible to us. The fading of the emission features canthen be understood in terms of a decreasing optical thickness ofthe circumstellar matter, for example due to its expansion. Foroptically thin transitions the fading would be proportional to thedecrease of the optical thickness. However, for optically thicklines, as is probably the case of the Na i and K i resonant transi-tions, a modest decrease of their optical thickness would not af-fect their observed strengths relative to the continuum. The fluxin these lines is limited by the available radiative flux from thecentral source. Thus these emission features fade proportionallyto the incident flux, so that their ratio to the observed M-typecontinum remains unchanged.
9. Summary and discussion
We have reduced the best-quality optical spectrum ever obtainedfor V4332 Sgr. The spectrum was recorded in April / May 2005with the VLT / UVES equipment and is unique not only be-cause of its quality, but also because the object has considerablyevolved since that time. Most of the emission features that werenumerous in 2005 disappeared, which means that it will not bepossible, in the future, to repeat the observations of the object ina similar stage. Therefore we decided to reduce and publish thespectrum, so that it is made available to the astrophysical com-munity. We presented the spectrum and results of its measure-ments and general analysis. Forthcoming papers will be devotedto detailed analyses of the emission features, their profiles, in-tensities, and chemical species that produce them.The spectrum is dominated by numerous atomic and molecu-lar emission features that are superimposed on an M-type contin-uous spectrum. Among the emission features we have identifiedover 70 atomic lines belonging to 11 elements (Na, Mg, Al, K,Ca, Cr, Mn, Fe, Rb, Sr, and Ba) (see Table 2) and about 140bands belonging to 6 molecules (AlO, ScO, TiO, VO, CrO, andYO) (see Table 3). There is no other late-type stellar object withan emission spectrum that rich and intense.The underlying stellar-like continuum in the spectrum ofV4332 Sgr can be classified as ∼ M3 (see Sect. 4). A giant spec-trum of this spectral type fits the observed continuum quite well.Moreover, a MARCS model spectrum calculated for an e ff ectivetemeprature characteristic of ∼ M3, that is, ∼ E B − V ≃ .
35 (see Sect. 5). Thisvalue is consistent with previous determinations (Tylenda et al.2005; Kimeswenger 2006; Kami´nski et al. 2010). However,since the observed continuum is assumed to result from scatter-ing on dust grains, this value probably underestimates the inter-stellar reddening. Therefore we attempted to estimate the extinc-tion from the strongest emission lines because the highest fluxin these lines is expected to follow the flux from the central star.Comparing these data with the standard M3 spectrum and theMARCS model, we obtained E B − V ≃ .
75. This value shouldbe regarded as an upper limit to the real reddening because ofpossible dust absorption e ff ects of multiply scattered resonant-line photons. Thus we concluded that the interstellar reddeningof V4332 Sgr is 0 . < ∼ E B − V < ∼ . − . ± . − is consistent with the results fromother determinations (Martini et al. 1999; Tylenda et al. 2005;Kami´nski et al. 2010) in the sense that all of them gave negativevalues, although the values range between −
180 to −
56 km s − .As already discussed in Tylenda et al. (2005), these values showthat V4332 Sgr does not follow the Galactic rotation curve be-cause from the latter one would expect a heliocentric radial ve-locity of > ∼ −
10 km s − for any distance for the object coordi-nates. The problem becomes even more evident if one takes intoaccount the lower limit of the distance derived in Sect. 7, which,assuming the Galactic rotation curve, implies a radial velocity > ∼
60 km s − . One possible explanation is that what we mea-sure in the emission lines is not the radial velocity of the ob-ject, but the expansion of the matter ejected in 1994. It is di ffi -cult to understand, however, why a decade after the eruption, wewould not see the receding part of the ejecta, especially becausethe radial velocity was determined from medium-intensity lines,which are expected to be optically thin. Another possibility isthat V4332 Sgr is not a Galactic disc object, hence its strangeradial velocity. This interpretation would be supported by therelatively large distance of the object from the Galactic plane of > ∼ ∼ M3 spectral type to M5-6, so that the e ff ective temperatureof the central stellar object decreased by 300-350 K. The sameconclusion also results from photometric measurements made in2003 and 2009. This relatively small ( ∼
10 %) decrease in thee ff ective temperature is fully responsible for the optical fadingof the object observed in 2006. The object expanded during thisevent, however, so that its luminosity increased by ∼
50 %. Thisis evident not only from fitting the optical observations, but alsofrom the infrared photometry, which is of particular importancebecause the infrared emission dominates the observed spectralenergy distribution of the object (see e.g. Kami´nski et al. 2010).The origin of this behaviour is not clear, but it might be a man-ifestation of a long-term relaxation of the remnant of the pre-sumable merger event in 1994. The main object is embedded ina massive dusty envelope, probably in the form of a thick disc.An accretion event from the envelope to the central stellar objectcan be invoked here.
Article number, page 8 of 14ylenda et al.: Optical spectrum of V4332 Sgr
The good resolution and quality of the spectrum we anal-ysed allowed us to detect several components of the interstel-lar absorption in the Na i ∼ > ∼ . × L ⊙ , while its e ff ective radius becomes > ∼
450 R ⊙ .These values are not as high as those derived for V838 Mon inits 2002 eruption (see e.g. Tylenda 2005), but are of a similarorder as those of V1309 Sco in its 2008 eruption (Tylenda et al.2011).There are, in fact, other similarities between V4332 Sgr andV1309 Sco. The eruptions of both objects were of a similar timescale, that is, about one month. That of V4332 Sgr was prob-ably slightly longer because its rising part was not observed.The progenitor of V4332 Sgr was variable, and the archive datasuggest that it was slowly rising in brightness on a time scaleof decades (Kimeswenger 2007; Goranskij et al. 2007). This re-sembles the slow systematic rise of V1309 Sco observed duringa few years before its eruption (Tylenda et al. 2011). The rem-nants of both objects are strongly dominated by infrared dustemission (Kami´nski et al. 2010; Nicholls et al. 2013). Finally, anoptical and near-IR spectrum of V1309 Sco obtained in 2012(Kami´nski et al. 2015b) shows an emission-line spectrum simi-lar to that of V4332 Sgr, although not as rich and intense as in thelatter case. Clearly, the remnant of V1309 Sco is heavily embed-ded in dust as was the case of V4332 Sgr. It is therefore temptingto suggest that the nature of the progenitor and the eruption ofV4332 Sgr was similar to those of V1309 Sco. We know that theeruption of V1309 Sco resulted from merger of a contact binary(Tylenda et al. 2011). Acknowledgements.
We are grateful to Dipankar P. K. Banerjee, PI of the ESO-VLT 075.D-0511 observing programme, for the spectroscopic observations ofV4332 Sgr made in 2005, on which this paper is based. Many thanks to thereferee for the comments on an earlier version of the paper.
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Table 2.
Atomic lines of the spectrum of V4332 Sgr. The flux is listed in 10 − erg cm − s − , the wavelengths and FWHM in Å. λ obs λ lab identification flux FWHM notes3860.47 3859.91 FeI 0.37 ( ± ± ± ± ± ± ± ± ± ⌉ B4031.8 4033.06 MnI | B4033.5 4034.48 MnI ⌋ B4076.66 ? 0.20 ( ± ± ± ± ± ± ± ± ± ± ⌉ B4339.72 CrI ⌋ B4343.39 4344.50 CrI 0.34 ( ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 10 of 14ylenda et al.: Optical spectrum of V4332 Sgr
Table 2.
Continued λ obs λ lab identification flux FWHM notes5326.72 5328.04 FeI 0.99 ( ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Article number, page 11 of 14 & Aproofs: manuscript no. 25592_ap
Table 3.
Molecular bands of the spectrum of V4332 Sgr. The flux is listed in 10 − erg cm − s − , the wavelengths in Å. λ lab identification flux notes4470.5 AlO B Σ + -X Σ + (2,0) 0.38 ( ± Σ + -X Σ + (3,1) 0.28 ( ± Σ + -X Σ + (4,2) 0.29 ( ± Σ + -X Σ + (1,0) 9.07 ( ± Σ + -X Σ + (2,1) 6.43 ( ± Σ + -X Σ + (3,2) 2.26 ( ± Σ + -X Σ + (4,3) 0.62 ( ± Σ + -X Σ + (5,4) p4760.9 TiO α (2,0) R ± α (3,1) R p4842.3 AlO B Σ + -X Σ + (0,0) 43.52 ( ± Σ + -X Σ + (1,1) 16.26 ( ± Σ + -X Σ + (2,2) 3.63 ( ± α (1,0) 5.02 ( ± α (2,1) 2.82 ( ± Σ − -X Σ − (3,0) 2.32 ( ± Σ + -X Σ + (0,1) 31.41 ( ± Σ + -X Σ + (1,2) 17.71 ( ± Σ + -X Σ + (2,3) 7.83 ( ± Σ + -X Σ + (3,4) 1.94 ( ± Σ + -X Σ + (4,5) p5166.7 TiO α (0,0) 14.46 ( ± ⌉ B5166.3 FeI | B5168.9 FeI ⌋ B5228.2 VO C Σ − -X Σ − (2,0) 1.15 ( ± Σ − -X Σ − (3,1) p5336.5 AlO B Σ + -X Σ + (0,2) 2.07 ( ± Σ + -X Σ + (1,3) 2.43 ( ± Σ + -X Σ + (2,4) 1.61 ( ± α (0,1) 10.13 ( ± Σ − -X Σ − (1,0) 9.84 ( ± α (1,2) 7.33 ( ± Σ − -X Σ − (2,1) 1.85 ( ± Π − -X Π − (2,0) 3.79 ( ± ⌉ B5565.9 CrO B Π -X Π (2,0) | B5567.7 CrO B Π -X Π (2,0) | B5570.2 CrO B Π -X Π (2,0) | B5576.4 CrO B Π -X Π (2,0) ⌋ B5736.7 VO C Σ − -X Σ − (0,0) 10.47 ( ± γ ′ (1,0) F -F ± γ ′ (1,0) F -F ± γ ′ (1,0) Q ,P p,B5821.2 TiO γ ′ (2,1) F -F ± γ ′ (2,1) R ± γ ′ (2,1) Q ,P ± Π / -X Σ + (0,0) 5.24 ( ± Π / -X Σ + (0,0) 51.23 ( ± Π − -X Π − (0,0) 21.66 ( ± ⌉ B6053.3 CrO B Π -X Π (0,0) | B6054.8 CrO B Π -X Π (0,0) | B6058.5 CrO B Π -X Π (0,0) | B6063.5 CrO B Π -X Π (0,0) | B6064.3 ScO A Π / -X Σ + (0,0) R R G ⌋ B,p
Article number, page 12 of 14ylenda et al.: Optical spectrum of V4332 Sgr
Table 3.
Continued λ lab identification flux notes6072.6 ScO A Π / -X Σ + (1,1) 2.44 ( ± Π / -X Σ + (0,0) 74.26 ( ± ⌉ B6086.4 VO C Σ − -X Σ − (0,1) ⌋ B6116.0 ScO A Π / -X Σ + (1,1) 3.12 ( ± Π / -X Σ + (0,0) 11.16 ( ± ⌉ B6138.8 VO C Σ − -X Σ − (1,2) ⌋ B6148.7 TiO γ ′ (0,0) R ± γ ′ (0,0) F -F ± γ ′ (0,0) Q ± γ ′ (0,0) F -F ± γ ′ (1,1) R ± γ ′ (0,0) F -F ± γ ′ (1,1) R ± γ ′ (1,1) Q , P ± γ ′ (1,1) F -F ± γ (2,0) R ± γ ′ (2,2) Q ± Π − -X Π − (0,1) 33.37 ( ± ⌉ B6396.2 CrO B Π -X Π (0,1) | B6397.8 CrO B Π -X Π (0,1) | B6401.4 CrO B Π -X Π (0,1) | B6407.7 CrO B Π -X Π (0,1) ⌋ B6451.7 CrO B Π − -X Π − (1,2) 14.65 ( ± ⌉ B6451.9 CrO B Π -X Π (1,2) | B6455.2 CrO B Π -X Π (1,2) | B6459.5 CrO B Π -X Π (1,2) | B6465.4 CrO B Π -X Π (1,2) ⌋ B6477.8 VO C Σ − -X Σ − (0,2) 5.46 ( ± Σ − -X Σ − (1,3) 11.48 ( ± γ ′ (0,1) F -F ± γ ′ (0,1) R ± ⌉ B6597.9 TiO γ ′ (0,1) Q ,P ⌋ B6618.0 TiO γ ′ (1,2) F -F ± γ ′ (0,1) Q , P ± γ (1,0) F -F ± ⌉ B6649.8 TiO γ ′ (1,2) F -F ⌋ B6680.8 TiO γ (1,0) F -F ± ⌉ B6681.8 TiO γ ′ (1,2) F -F ⌋ B6714.5 TiO γ (1,0) F -F ± ⌉ B6717.6 TiO γ (2,1) F -F ⌋ B6747.6 TiO γ (2,1) F -F ± Π − -X Π − (0,2) 24.66 ( ± ⌉ B6774.2 CrO B Π -X Π (0,2) | B6775.9 CrO B Π -X Π (0,2) | B6779.6 CrO B Π -X Π (0,2) | B6781.8 TiO γ (2,1) F -F | B6785.7 CrO B Π -X Π (0-2) ⌋ B6815.1 TiO γ (3,2) F -F ± Π -X Π (0,0) R ± ⌉ B6836.5 CrO B Π − -X Π − (1,3) | B6836.5 TiO b Π -X Π (0,0) Q | B6836.6 CrO B Π -X Π (1,3) | B6839.9 CrO B Π -X Π (1,3) | B6844.5 CrO B Π -X Π (1,3) | B6850.9 CrO B Π -X Π (1,3) ⌋ B6919.0 VO C Σ − -X Σ − (0,3) 1.51 ( ± Σ − -X Σ − (1,4) 1.05 ( ± Article number, page 13 of 14 & Aproofs: manuscript no. 25592_ap
Table 3.
Continued λ lab identification flux notes7054.2 TiO γ (0,0) F -F ± γ (0,0) F -F ± γ (0,0) F -F ± ⌉ B7124.9 TiO γ (1,1) F -F ⌋ B7158.8 TiO γ (1,1) F -F ± γ (1,1) F -F ± ⌉ B E,s7196.4 TiO γ (2,2) F -F ⌋ B7270.1 TiO γ (2,2) F -F ± Π / -X Σ − (1,0) 14.38 ( ± Π / -X Σ − (1,0) 17.74 ( ± Π / -X Σ − (1,0) 17.23 ( ± Π − / -X Σ − (1,0) 21.77 ( ± γ F -F (0,1) > ± γ F -F (0,1) p,u,s7665.8 TiO γ F -F (1,2) B,u7671.6 TiO γ F -F (0,1) p,B,u7704.9 TiO γ F -F (1,2) p,B,u7743.0 TiO γ F -F (2,3) p,B7749.5 TiO γ F -F (1,2) 27.48 ( ± γ F -F (2,3) 4.77 ( ± γ F -F (2,3) 5.69 ( ± Π / -X Σ − (0,0) 87.16 ( ± Π / -X Σ − (0,0) 87.96 ( ± Π / -X Σ − (0,0) 100.888: ( ± Π − / -X Σ − (0,0) 158.93 ( ± Π − / -X Σ − (1,1) 19.69 ( ± ǫ (0,0) R ± ⌉ B8442.3 TiO ǫ (0,0) R ,Q ,P | B8451.8 TiO ǫ (0,0) R ,Q ,P | B8462.7 TiO ǫ (0,0) Q ,P ⌋ B8682.6 VO B Π − / -X Σ − (0,1) 33.45 ( ± Π − / -X Σ − (1,2) R ± Π − / -X Σ − (1,2) P ± Table 4.
Meaning of symbols used in Tables 2 and 3. ]B - measurements refer to the whole blendc - uncertain continuum levelf - faint features - shape-fitting was problematicalb - blended featuresB - very strong blendingE - corrected for interstellar and / or telluric absorption: - uncertain > - lower limitp - present or probably present, but unmeasurableu - unmeasurable (cannot be deconvolved)- lower limitp - present or probably present, but unmeasurableu - unmeasurable (cannot be deconvolved)