{'fraction_non_alphanumeric': 0.037612294899398616, 'fraction_numerical': 0.03976538718538867, 'mean_word_length': 4.1864458991143625, 'pattern_counts': {'":': 0, '<': 10, '<?xml version=': 0, '>': 8, 'https://': 0, 'lorem ipsum': 0, 'www.': 0, 'xml': 0}, 'pii_count': 0, 'substrings_counts': 0, 'word_list_counts': {'cursed_substrings.json': 0, 'profanity_word_list.json': 0, 'sexual_word_list.json': 0, 'zh_pornsignals.json': 0}}
The spectacular evolution of Supernova 1996al over 15 years: a low energy explosion of a stripped massive star in a highly structured environment

Spectrophotometry of SN 1996al carried out throughout 15 years is presented. The early photometry suggests that SN 1996al is a Linear type-II supernova, with an absolute peak of Mv ~ -18.2 mag. Early spectra present broad, asymmetric Balmer emissions, with super-imposed narrow lines with P-Cygni profile, and He I features with asymmetric, broad emission components. The analysis of the line profiles shows that the H and He broad components form in the same region of the ejecta. By day +142, the Halpha profile dramatically changes: the narrow P-Cygni profile disappears, and the Halpha is fitted by three emission components, that will be detected over the remaining 15 yrs of the SN monitoring campaign. Instead, the He I emissions become progressively narrower and symmetric. A sudden increase in flux of all He I lines is observed between 300 and 600 days. Models show that the supernova luminosity is sustained by the interaction of low mass (~1.15 Msun) ejecta, expelled in a low kinetic energy (~ 1.6 x 10^50 erg) explosion, with highly asymmetric circumstellar medium. The detection of Halpha emission in pre-explosion archive images suggests that the progenitor was most likely a massive star (~25 Msun ZAMS) that had lost a large fraction of its hydrogen envelope before explosion, and was hence embedded in a H-rich cocoon. The low-mass ejecta and modest kinetic energy of the explosion are explained with massive fallback of material into the compact remnant, a 7-8 Msun black hole.


INTRODUCTION
The aftermath of a supernova explosion is a bubble of gas expanding at very high velocity. Eventually, the ejecta may impact on the pre-existing circumstellar material (CSM), generating a shock in which a fraction of the kinetic energy of the ejecta is converted into radiation. The intensity of the resulting emission mostly depends on the density of the CSM and the velocity contrast between the ejecta and the CSM. If the density of the CSM is low, the emission from the CSM-ejecta interaction becomes visible only after the SN luminosity has faded, sometimes several years after the explosion. However, occasionally, the CSM near the SN is so dense that the CSM-ejecta interaction dominates the SN emission even at early phases.
With improved statistics and quality of observations, we are now observing counterparts for the different scenarios, and a new taxonomy for type II supernovae (SNII), based on the mass of the residual hydrogen envelope and the strength of the CSM-ejecta interaction signatures, can be proposed. The emission of classical SNII is determined by the thermal balance in the ejecta, and CSM-ejecta interaction is negligible (at least at early phases). Depending on the shape of the light curve, they are sub-classified into plateau (IIP) and linear (IIL) types. The current interpretation is that different light curves correspond to different envelope masses, ranging from ∼ 10 M in SNIIP (cf. Smartt 2009) to ∼ 1 M in SNIIL, due to different progenitor initial masses and mass loss histories. When soon after the explosion the SN ejecta interact with a dense CSM, the SN emission itself is overwhelmed by the emission arising from the interaction (Chugai 1990;Chevalier & Fransson 1994;Terlevich 1994). Best examples of this class are SN 1988Z (Turatto et al. 1993;Aretxaga et al. 1999), with a long plateau-like light curve, and SN 1998S (Liu et al. 2000;Fassia et al. 2000) with a linear-like light curve. These SNe are named as type IIn.
In some SNII, the CSM-ejecta interaction contribution emerges only when the 56 Co decay energy input fades. The fact that these are mostly SNIIL is consistent with our understanding that they experienced stronger mass loss during their evolutions (e.g. SNe 1980K, 1979C, and 1990K;Uomoto & Kirshner 1986;Fesen et al. 1999;Cappellaro et al. 1995, respectively), hence they are expected to have higher CSM density.
Somewhat in between, a few SNe have emission line profiles which reveal the presence of a CSM around the exploding stars from early-on, though the ejecta-CSM interaction becomes preponderant some months later. These are sometimes labeled SNIId, where 'd' stands for 'double' profile because of the simultaneous presence of broad profiles from the ejecta and narrow ones from the CSM. SN 1994aj (Benetti et al. 1998, Paper I) and SN 1996L (Benetti et al. 1999, Paper II) belong to this group.
In this paper we present the study of a new SN (SN 1996al) with spectroscopic properties very similar to those of SNIId (Benetti & Neuhauser 1996). SN 1996al was visually discovered on July 22.71 UT by Evans et al. (1996), about 30" north of the center of its host galaxy, NGC 7689. The measured position of the SN is R.A.= 23 h 33 m 16 s .29±0.03; Decl.= −54 o 04'59".69±0.05 (J2000.0), in agreement with that reported by Evans et al. (1996). In view of its relative closeness and apparent brightness, an extensive observational campaign was set-up with the ESO-Chile telescopes. Thanks to this massive effort, we have been able to follow the SN for more than 15 years, and it has become one of the best monitored SNIIL-SNIId. In this paper we will describe the physical properties of the explosion and constrain the nature of SN 1996al progenitor star.

OBSERVATIONS
Spectroscopy and imaging observations were carried out at ESO La Silla and Paranal Observatories using a number of different telescopes and/or instruments (see Tables A1 &  A2).

Optical Photometry
The CCD frames were first bias and flat-field corrected in a standard manner. Since a fraction of data was obtained under non-photometric conditions, relative photometry was Figure 1. SN 1996al in NGC 7689 and reference stars. The main image is an R frame taken on 1996 October 1 with the D1.54m+DFOSC. The seeing is 1.4" and the field of view is 5.9'x5.9'. The zoomed image was obtained with VLT+FORS2 using a redshifted Hα filter on 2002 June 16. The seeing was 0.7" and the field of view is 25". The SN is the source at the center of the image.
performed with respect to a local sequence of field stars (see Figure 1). Twelve photometric nights were used to calibrate this sequence against Landolt standard stars (Landolt 1992). The magnitudes and estimated errors of the local standards are shown in Table 1. These magnitudes were obtained by summing the counts through an aperture, whose size varied according to the seeing.
Ideally, one would like to remove the galaxy background by subtracting a galaxy "template" where the SN is absent. However, a suitable template image was not available. Therefore, the SN magnitudes were measured using the IRAF 1 point spread function fitting task Daophot (Stetson 1992) and/or ROMAFOT package (Buonanno et al. 1983) under MIDAS 2 . These procedures allowed the simultaneous fit and subtraction of the galaxy background. While the pixel scales changed from one instrument to another, they were always sufficiently small to provide good sampling of the PSF (see last column, and caption of Table A1). The supernova magnitudes with estimated internal errors are listed in Table A1, which also gives the seeing for each epoch, averaged over the observed bands. On Dec. 8, 2007 an Hα frame of SN 1996al was also secured with VLT+FORS2, and the frame was calibrated observing in the same night the spectrophotometric standard G158-100 (Oke 1990). The derived Hα flux inside the SN 1996al  (Barbon et al. 1982;Buta 1982;Tsvetkov 1983). The distance to the SN 1980K host galaxy, NGC 6948, is taken from Karachentsev et al. (2000). The SN 1980K B, V light curves are reported to the distance and reddening of SN 1996al with a +0.3 mag shift, and labelled Br and Vr in the figure. The dashed lines represent the B and V light curves of SN 1998S (Liu et al. 2000;Fassia et al. 2000). The distance for the SN 1998S host galaxy, NGC 3877, is taken from Willick et al. (1997). The SN 1998S B, V light curves are reported to the distance and reddening of SN 1996al with a +1.5 mag shift, and labelled Br and Vr in the figure The starry symbol is the visual SN estimate at the discovery from Evans et al. (1996). The best match with the template V light curve of SNIIL (from Doggett & Branch 1985), along with the comparisons with those of SNe 1980K and 1998S, suggest that the maximum occurred on MJD ∼ 50265 ± 2 days at a magnitude V ∼ 13.9 ± 0.3.
PSF was found to be in excellent agreement with that derived from the subsequent FORS2 spectrum (Sect. 2.4).

Light curves
The U BV RIJH light curves are plotted in Figure 2. Patat et al. (1994) showed that SNII can be characterised on the basis of the β100 parameter, which quantifies the luminosity decline rate in the first 100 days after maximum light. For SN 1996al, we measure β B 100 = 4.84 mag (100d) −1 , which is typical of the type IIL subclass, very similar for instance to SN 1980K. However, the post maximum decline shows a short plateau-like break (see inset of Figure 2) in all bands. This flattening is reached at day +30 in U and B bands, and is very short (2 − 3 days). In V RI it is reached a few days later (+34d), and it lasts about one week.
The comparisons with the light curves of SNe 1980K and 1998S V , and the V -band template light curve of SNIIL from Doggett & Branch (1985), suggest that SN 1996al reaches the V maximum on MJD=50265±2 (1996 June 30), i.e. 23 days prior to our first observation, and likely with a V peak magnitude of about 13.9 ± 0.3 (B = 13.7 ± 0.3). In the following, we will adopt this estimate of the V maximum epoch as reference time. The epoch of the B-band peak is coincident with that of the V maximum. However, we stress that these estimates are subject to relatively large uncertainty, since the maxima themselves were not observed and Type IIL SNe show a variety of photometric behaviours around the maximum epochs (Faran et al. 2014). In Table 2, the main parameters of the U BV RI light curves of SN 1996al are summarised.
An inflection point in the linear decline is observed at about day +150. Between +150 and about +320 days the light curves, though not well sampled, show declines that match the 56 Co decay rate (0.98 mag(100) −1 days (Table  2). From +320 to +450 days the light curves level-out to a new plateau, where the V light curve slightly increases in luminosity by +0.2 mag, and R by ∼ +0.1 mag. After this plateau, the BV RIJ light curves start a monotonic decline (with some modulation in the R-band) with rates very similar to those observed between 150 and 300 days (Table  2).

Colour curves
In Figure 3, the extinction corrected 3 (U − B)0, (B − V )0 and (V − R)0 colour curves of SN 1996al are plotted along with those of two bright SNIIL, SNe 1980K and 1990K, and SN 1998S. The early (U − B)0 colour curve of SN 1996al is consistent with that of SN 1980K. It shows a smooth increase reaching a value of -0.1 mag at about 70 days. From day +140, it experiences a linear increase from -0.30 to +0.13 mag at phase 450 days, although the colour curve sampling is poor. However, it never reaches the red colours observed in SN 1980K (+0.5 mag). The (B − V )0 colour curve shows an overall similar behaviour as the (U − B)0 colour curve. It is initially blue (+0.07 mag), then it smoothly increases to a maximum of +0.54 mag at ∼ 70 days, followed by a slow decline. At late times, the (B − V )0 colour curve remains substantially constant at +0.34 mag with a small decrease to about 0.24 mag at about 460 days after maximum.
The overall appearance of the (B − V ) colour curve is very similar to that of SN 1998S, and not very discrepant from that of SN 1980K, although somehow displaced in phase. Also in this colour, SN 1996al never reaches the red values ((B − V )∼ 1 mag) shown by SN 1980K. SN 1990K, another bright SNIIL studied by Cappellaro et al. (1995), showed very red colours already at maximum. At phases close to ∼ +160 days, all four SNe have similar (B − V ) colours. Finally, the SN 1996al (V − R) colour curve experiences a gradual "reddening" with early values similar to those seen in SN 1990K. After +100 days, the V − R colour curves of the two SNe show a plateau-like feature, with a (V − R)∼ 0.41 mag for SN 1996al. Starting at about day +120, (V − R) further increases to ∼ 0.9 mag, and remains thereafter almost constant.

Reddening and distance to NGC 7689
The Galactic reddening towards NGC 7689 is AB=0.051 mag (Schlegel et al. 1998). The higher S/N ratio spectra of SN 1996al show narrow absorptions due to interstellar Na ID, suggesting that some light extinction may occur within the parent galaxy. The mean EW of this doublet averaged over 10 measurements is 0.60 ± 0.08Å. Using the relation E(B − V ) ∼ 0.16 × EW (Na ID) of Turatto et al. (2003) yields E(B − V ) host ∼ 0.10 ± 0.05 mag. In the following, we will adopt for SN 1996al a total reddening of E(B − V )tot = 0.11 ± 0.05 mag. From the Na ID interstellar lines we also derive a mean recessional velocity of 1835 ± 68 km s −1 for NGC 7689, to be compared with the optical (1974 ± 6 km s −1 ) and radio (1970 ± 5 km s −1 ) velocities given by NED 4 . The difference (-135 km s −1 ) is likely due to the host galaxy rotation, projected along the line-of-sight (LOS).
Adopting a value of H0=73 km s −1 Mpc −1 along with the radio heliocentric radial velocity of 1967 km s −1 , and correcting it for Local Group infall onto the Virgo cluster, NED gives for NGC 7689 a distance modulus of µB = 31.98 (24.89 Mpc).
On the other hand, using the relative distance of NGC 7689 from the Virgo cluster (of 1.37, Kraan-Korteweg 1986), and assuming a Virgo cluster distance of 15.3 Mpc (Freedman et al. 2001), a distance modulus of µB = 31.61 mag (20.96 Mpc) is derived for NGC 7689.
The mean value of the above estimates, µB = 31.80 ±  Evans et al. (1996) 0.20 mag (22.91 Mpc) is adopted throughout the paper as distance modulus for NGC 7689.

Absolute magnitudes and bolometric light curve
With the above assumptions for extinction and distance (using the Cardelli et al. (1989) extinction law), we obtain MB ∼ −18.6 ± 0.4 and MV ∼ −18.2 ± 0.3 for the SN absolute B and V magnitudes at maximum. These are similar to those of the "bright" SNIIL sub-class (which have a < MB >= −19.0 ± 0.6 mag, rescaled to H0=73 km s −1 Mpc −1 ; Patat et al. 1994).
In order to compute the energy budget of SN 1996al and to simplify the comparison with theory, it is useful to determine the bolometric light curve. Unfortunately, no observations at UV, X-ray or radio wavelengths are available. We are then limited to derive a pseudo-bolometric (U BV RIJHK, or uvoir) light curve. This was obtained first adding the U BV RIJ contributions and then calculating the HK contributions in the phase range 153-408 days, assuming during this period a constant H − K colour. The HK contribution resulted in logLUtoK −logLUtoJ = 0.08 ± 0.01 dex for this period. We made the assumption that the IR contribution to the total luminosity remained constant at earlier epochs. The bolometric curve shown in Figure  4 shares all relevant features already noted in chromatic curves: the overall linear decay up to 150 days after explosion, including a hint of a plateau around day +30. It shows a break in the decay at around day +150, then a 56 Co like decay up to about day +300. This is followed by a plateau-like phase with an almost constant bolometric luminosity (+300 to +450 days) and, finally, a new decay with a rate close to that of 56 Co.
In Sect. 3, we will argue that most probably in SN 1996al, the CSM/ejecta interaction is the primary energy source for a large part of its evolution. However, comparing the early part of the linear tail of SN 1996al, between ∼ 150 and 300 days after explosion, with the radioactive tail of SN 1987A, one can derive an upper limit for the 56 Ni ejected in the explosion. Assuming a 56 Ni mass of 0.075 ± 0.005 M for SN 1987A (Danziger 1988;Woosley, Hartmann, & Pinto 1989), and adopting for SN 1996al a scaling factor of 0.24 as suggested by Figure 4, we derive that the 56 Ni ejected in the explosion of SN 1996al is 0.018 ± 0.007 M , where the error budget is equally dominated by the uncertainties on the SN 1996al distance, the epoch of the explosion, and the SN 1987A 56 Ni mass. We can also estimate the 56 Ni mass from the observed luminosity assuming full γ-rays trapping and a rise time to maximum of 10 days. Hence, the amount of 56 Co necessary to fit the tail from 100 to 300 days is 0.017 ± 0.003 M , fully consistent with that derived from the SN 1987A scaling.  . Bolometric uvoir light curve of SN 1996al. A zoom of the uvoir curve in phases near maximum is shown in the upperright box. For comparison, the 56 Ni+ 56 Co decay model in the hypothesis of full γ-rays trapping is also reported. The rise time to maximum was assumed to be 10 days.
A summary of the main parameter values inferred for SN 1996al and its host galaxy is provided in Tab. 3.

Spectroscopy
Spectroscopic observations spanned from days +23 to +5542, with an excellent temporal coverage during the first two months. Table A2 lists the date (col. 1), the Modified Julian Day (col. 2), the phase relative to tV max (col. 3), the wavelength range (col. 4), the instrument used (col. 5), and the resolution as measured from the FWHM of the night-sky lines (col. 6). At some epochs, almost contemporary spectra were obtained. These were merged to produce higher S/N spectra or wider wavelength ranges.
The spectra were reduced following standard MIDAS and IRAF routines. The XShooter (Vernet et al. 2011) spectra were reduced using the ESO-XShooter pipeline version 1.1.0. One dimensional spectra were extracted weighting the signal by the variance based on the data values and a poisson/CCD model using the gain and readout-noise parame-ters. The background to either side of the SN signal was fitted with a low-order polynomial, and then subtracted. The flux calibration and telluric absorptions modelling were achieved using spectra of spectrophotometric standard stars. Most spectra were taken through a slit aligned along the parallattic angle, or normalised to spectra taken with a wider ( 5 arcsec) slit. The flux calibration of the spectra was checked against photometry (using the IRAF task stsdas.hst calib.synphot.calphot) and, if discrepancies occurred, the spectral fluxes were scaled to match the photometry. On nights with fair sky transparency, the agreement with photometry was within 15%. The spectra of SN 1996al are shown in Figure 5.

A brief description of the overall spectroscopic evolution
The earliest spectra show a relatively blue continuum (T bb ∼ 12000 o K), once corrected for reddening and redshift. The continuum becomes progressively redder with phase, reaching a temperature of about 6500 − 7000 o K at phases later than ∼60 days (see also Sect. 3.4). Early spectra are dominated by broad Hα and Hβ, and the blend of He I 5876Å-Na ID lines. The Balmer lines show complex profiles (see Figures 6  and 7), and the decomposition of the early Hα profile is detailed in Figure 8a. Hβ has a very similar profile. The detailed evolution of the Balmer lines (in particular Hα) is discussed in Sect. 3.2.
Another intense, broad emission detected in our first spectrum is a blend of the He I 5876Å-Na ID lines. It has a FWHM of ∼ 260Å centered at ∼ 5906Å, though it shows a double-peaked profile (see Figure 6). The detailed evolution of this and other unblended He lines will be discussed in Sect. 3.3.
A broad absorption is observed with a minimum at 5079 A. Assuming it is FeII 5169Å, it gives an expansion velocity of ∼ 5200 km s −1 . Also, our early spectra show narrow Fe II (5018-5169Å) lines with P-Cygni profiles. These are clearly visible in the SN spectra up to phase +43 days, with a mean velocity from the absorption minima of 480 ± 150 km s −1 .
The narrow lines likely arise from a dense CSM, which was ionised either by the X-UV flash shock break-out, or by the UV photons emitted by the progenitor star (as suggested by the Hα emission seen around the progenitor before the explosion, cf. Sect 3.1). The CSM, which is recombining at Figure 5. Spectral evolution of SN 1996al. Wavelengths are in the observer's frame. The ordinate flux scale is optimised to the first spectrum, and all others are arbitrarily shifted downwards. In a few cases, noise spikes have been manually removed. From blue to red wavelengths, the dotted (blue) vertical lines mark the positions of Hγ, Hβ, and Hα, respectively; while the dashed (red) lines mark the 5876Å, 6678Å, 7065Å and 7281Å He I transitions. The positions of the most intense telluric absorption bands are marked. c 0000 RAS, MNRAS 000, 000-000 Figure 6. Earliest spectra of SN 1996al corrected for redshift and reddening. The main features are identified, together with the different components seen in Hβ and Hα. The narrow absorption (marked "int.") on the top of the He I λ5876-Na ID emission is due to Na ID Galactic absorption. The dashed (red) lines are blackbody fits of the spectral continuum (T=10100 o K and 8600 o K for +23 and +29 days, respectively; see Sect. 3.4).
these phases, has a bulk velocity of ∼ 250 km s −1 , as deduced from Hβ narrow emission, but extends up to ∼ 2000 km s −1 , as inferred from the blue edge of the narrow Balmer absorptions.
The spectrum does not change significantly during the first few weeks, with the continuum becoming progressively redder. At phase +29 days, we are able to monitor for the first time a more extended wavelength range (see Figure 6). Starting from the blue-end of the spectrum, we see Hδ and Hγ. They have complex profiles similar to those of Hβ and Hα. Redward to Hα, there is possible evidence for the He I 7065Å transition (see discussion in Sect. 3.3), while the He I 6678Å might be embedded in the Hα red wing. The +41 day spectrum shows broad emissions (centered at 8531 and 8677Å) due to the CaII IR triplet. The FWHM of the total emission is about 3500 km s −1 . The He I 5876Å-Na ID feature still has a very broad profile similar to the emission red-ward of Hα, which we interpret as mostly due to He I 7065Å.
The subsequent spectrum (phase +53 days) shows the appearance of other important features. The broad Hα profile is now better reproduced with a Gaussian profile (having a FWHM∼ 5700 km s −1 , see Table A3), and its red wing is better fitted adding a broad (FWHM∼ 4500 km s −1 ) emission centered at 6699Å (rest frame) in Figure 8b. We identify this emission as the He I 6678Å transition.
In general, the He I lines become more symmetric and narrow with time, with their flux sharply rising after +300 day (cfr. Sect. 3.3). The He I transitions show peculiar line ratios (see Figure 5 and Tables A3, A4) and then disappear at about phase +700 days (see Sect. 3.3 for a detailed discussion on the He I lines evolution). As at these early epochs the FWHM velocities of the Hα broad component are similar to those of the He I lines, it is plausible that these lines originate in the same region of the ejecta and are probably powered by the same mechanism. The evolution of the Hα/Hβ ratio tells that this mechanism is not photoionisation. In fact, starting from phase +43 days this ratio starts a linear increase from a value close to the case B recombination value of 2.8, to ∼ 10 − 12 about 1000 days after the V -band maximum (see Sect. 3.2).
At epochs +142 days, the Balmer lines undergo another profound change in their profiles (see Figure 8c). The narrow lines arising in the un-shocked CSM are not visible anymore. This could either imply that the ejecta have overtaken the bulk of the denser CSM or, more likely, that the CSM has by that time completely recombined. The Hα profile is now well fitted with three Lorentzians of comparable FWHM centered at 6539Å, 6569Å and 6597Å (see Table A4 and Sect. 3.2). These components will be visible along the residual SN lifetime. We will refer to them as Blue, Core and Red components. With time, the Core and Red components become weaker and weaker in comparison with the Blue one (see Figure 7). As already stressed, Hβ becomes weaker with time with respect to Hα (see Sect. 3.2), and the Hα/Hβ ratio reaches values exceeding about 15 at the latest epochs.

The progenitor star
A search for archival material concerning the parent galaxy of SN 1996al returned a deep Hα image performed for a portion of NGC 7689 with the Rutgers Imaging Fabry-Perot instrument + CCD Camera attached at the Cassegrain focus of the CTIO 4-m Blanco telescope. The image was taken in the context of a work aimed at studying the star formation in the disk of a sample of Sa galaxies (Caldwell et al. 1991). Luckily enough, the portion of NGC 7689 imaged on September 1988 included the zone where SN 1996al appeared eight years later. Unfortunately, these were the first digital observations performed with modern CCDs and at that time CTIO was not yet equipped with a digital archive (CTIO director, private communication). For this reason, the original frames stored on a 9-track tape by the observer were lost. In order to properly place the SN location on the map we digitalised Plate 29 of Caldwell et al. and aligned it with the VLT+FORS2 Hα frame obtained under good seeing conditions on 2002, June 16 (see caption of Figure  1). After the alignment, performed with about 40 sources around the SN positions, we found that a faint Hα region is coincident with the SN position, within an uncertainty of ±0.75". Given the relatively clean environment at that position (cfr. Figure 1), we believe that the Hα emission visible in the Caldwell et al. image was indeed associated with the progenitor star of SN 1996al. The Hα luminosity of this source measured by Caldwell et al. but not reported in their original paper, given the distance and reddening assumed in this work (see Table 3), is logLHα ∼ 37.28 ± 0.10 dex. For consistency, we checked the Caldwell et al. Hα measurements of few H II regions close to the SN 1996al position with those derived from the FORS2 Hα frame, and found a very good agreement. The Hα pre-explosion luminosity is compared with those c 0000 RAS, MNRAS 000, 000-000 of some possible progenitor systems for SN 1996al in Table  4. The comparison indicates that the Hα emission coming from the SN 1996al progenitor is similar to that expected from hot LBV stars, and lower than that measured in SNimpostor/LBVs prior to a major outburst. This is consistent with the finding of Kennicutt & Hodge (1980), that H II regions in NGC 628 with logLHα < 37.90 dex have a radius < 200 pc, and are supported by the radiation field of a star with a spectral type later than O4 Ia. After profile decomposition (see Figure 8 and Table A3), the narrow Hα component visible in the first SN spectrum has a flux of ∼ 9 × 10 −14 erg s −1 cm −2 , which provides a line luminosity of logL∼ 39.85 dex. This is more than two dex higher than the pre-explosion Hα luminosity (see Table 4). The abrupt increase of the narrow Hα luminosity is a clear indication that the SN shock-break out or ejecta/CSM interaction ionised a much larger fraction of CSM around the progenitor star than the SN progenitor alone.

The evolution of the Balmer lines
The profile of Balmer lines, in particular Hα, shows a multi-component, complex evolution which allows us to track the restless mass loss activity of the progenitor star. It also indicates an asymmetric distribution of the emitting gas. In the following, we will have a close look into the Hα evolution in four temporal intervals.

Days +23 to +43
The Hα profile consists of a broad Lorentzian emission with, super-posed, a narrow P-Cygni like component. The Hα profile does not change (see Figure 7) until phase +40 days, which coincides with the end of the "short plateau" (see Figure 2). At the same time, the expansion velocity obtained from the Hα FWHM of the broad component decreases from ∼ 6500 km s −1 to ∼ 5600 km s −1 , while the maximum velocity deduced from the Hα blue wing remains constant at about 12000 km s −1 (see Table A3 and Figure 9). The average wind velocity as given by the width of the narrow emission shows some scatter around the value of ∼ 440 km s −1 . The mean FWHM of the narrow absorption is about 890 km s −1 (see Table A3 and Figure 9). The Hα luminosity evolution is negligible (see Figure 10), and the Hα/Hβ flux ratio, which remains constant at about 2.68 (see Figure 11), is close to case B recombination (Osterbrock 1989).

Days +53 to +101
The broad Hα is now best fitted by a Gaussian profile with a FWHM decreasing from 5700 km s −1 to 3250 km s −1 , though faint, extended wings are still visible (see Figure 8b). The maximum velocity deduced from the Hα blue wing remains, instead, constant at about 10500 km s −1 (see Table A3 and Figure 9). Again, these velocities are similar to those of the He I lines, hence we suppose that the lines originate in the same region of the ejecta. Since in the next Section we will demonstrate that at these phases the He I lines are powered by ejecta-CSM interaction, we believe that the same mechanism also supports the Hα flux. At day +100, the Hα luminosity is about 0.3 dex lower than the Hα luminosity for the twin Figure 7. Hα profile evolution of SN 1996al. Wavelengths are in the observer's frame. The feature near 6875Å in the +43 days spectrum is a residual from an poorly removed telluric band.
SN 1994aj at the same phase (Paper I), but is very similar to that of SN 1996L (Paper II).
The velocity of the wind, as derived from both the narrow emission and absorption features, shows a mean value of ∼ 590 km s −1 .
In this period, the Hα/Hβ flux ratio begins to increase, reaching a value of ∼ 6.5 at a phase of +100 days. This could be a clue that Hα becomes more collisionally excited (Branch et al. 1981), in line with the proposed scenario (see Sect. 3.5).
Days +142 to ∼ +800 From day +142, the Hα profile undergoes another profound change. The Hα profile is now better fitted by three Lorentzian components (see Figure 8c), up to phase +523 days, and later on with three Gaussian components. Since their earliest appearance, these components show a very similar FWHM, the Red component being the broadest (FWHM∼ 970 km s −1 ), followed by the Core component (FWHM∼ 950 km s −1 ), while the Blue component has a FWHM∼ 850 km s −1 (see Table A4 and Figure 9).
Between 142 and 800 days, the flux ratios of the different components stay constant, with the Core and Blue components carrying the largest fraction of the Hα flux (30-40% each). In this lapse of time, the Hα luminosity shows a relatively flat and linear (in a log-log diagram) decrease up to phase ∼ +500 days, followed by a sharper drop (by almost 1 dex) up to about +800 days. This behaviour resembles that of the bolometric curve which, however, stops at about 700 days after explosion (see Figure 4).
We interpret these Hα components as the result of the interaction between a mostly spherical ejecta with an highly asymmetric CSM. The geometric configuration of the system will be discussed with more detail in Sect. 3.5. The Hα/Hβ flux ratio continues to steadily rise (see Figure  11) to a value of about 10.
Later than +800 days There is an observational gap between +800 and +1850 days. From day +1853 the Hα profile is again fitted by the three Gaussians mentioned before, and both the Hα flux and the bolometric curve show steeper evolutions.
While the overall flux carried out by the line dim by 0.28 dex during this time interval, the relative intensities of the three components change: the Blue component is now much stronger with respect to the other two, while the Red one remains the dimmest (see Figure 13). At phase +2177 days, the flux ratio of the Core component decreases to the level of the Red component one, while the Blue component is carrying almost 80% of the line flux. After this phase, the line ratios do not change very much. The sharp change in the relative flux ratios may be a signature of either a sudden dust formation in the SN ejecta, or a progressive quenching of the interaction between the ejecta and the receding part of the CSM. Most probably, as the central wavelengths of the Core and Red emission peaks do not show any evolution (see Figure 12), both scenarios are at work. Also the Blue component does not show any shift from the central position. For this reason, if dust is forming, it condensates in the inner part of the ejecta or in the clumps where the Core and Red Hα components originate.
Starting from ∼ 2900 days after the explosion, the Hα luminosity (see Figure 10) decreases below the luminosity emitted by the precursor star. This, along with its distorted profile, proves that most of the Hα flux comes from the SN itself and that the ambient contamination is by now negligible. Moreover, this is a clue that the Hα source detected before the explosion was indeed associated to the precursor star and that it has been swept up by the SN ejecta.

The He I lines
In our first spectra, the best evidence for the presence of helium is the broad, asymmetric bump (FWHM∼ 11000 km s −1 ) centered at about 7065Å (rest frame) that can be identified as He I 7065Å (see Figure 6). Figures 14 and  15 show the complete evolution of He I 7065Å, the only helium transition which is not contaminated by close-by strong lines. In fact, this is not the case for He I 5876Å, which could be heavily contaminated by the nearby Na ID doublet. An estimate of the relative contribution of two lines throughout the complete spectroscopic evolution of SN 1996al is a very difficult task. The last phase in which the He I lines are unequivocally detected is +523 day (see Figure 5).
In Figure 16, we compare the profiles of He I 5876Å -Na ID with the He I 7065Å line, during the first 100 days of the SN evolution. In the first two spectra (+29 and +32 days) they are very similar, both in the overall profile and velocity limits. The lines show a double peaked profile with the blue one peaking at ∼ −2000 km s −1 and the red one at ∼ 3500 km s −1 , with the red wing extending far to higher velocities (∼ 9000 km s −1 ) than the blue one. Thanks to the similarity of the profiles, we may safely conclude that at phases +29, and +32 days the He I 5876Å -Na ID feature is almost entirely due to He. Figure 16 shows that, at least during the first month, the broad wings of the Hα profiles are similar to those of He I lines, suggesting that the three lines form in the same, asymmetric region. Figure 14 also shows that the profile of the He I 7065 A emission remains unchanged up to +53 days after the V maximum, with the blue bump dominating the line profile. By this time, the profile of He I 5876Å -Na ID feature starts to deviate from that of He I 7065Å (see Figure 16) with a broad bump emerging slightly redder than the Na ID rest wavelength position. Apart from this difference, the two He I lines and the broad Hα still share the same basic characteristics (see third panel from top in Figure 16).
Starting from phase +67 days, the He I 7065Å line shows a new component centered at the rest position, with the blue and red bumps still being visible. Also the He I 5876 A -Na ID feature has an overall similar profile, with an increasingly contribution from the Na ID. Afterwards, (on day +93 and +99) the He I 7065 and 7281 A features become progressively more symmetric, are centered at the rest position and reach terminal velocities of about 2500-3000 km s −1 (see bottom spectra in Figure 14). The He I 5876Å -Na ID profile is different, indicating an increased contribution by Na ID (see below). On the contrary, the terminal velocities of the Hα broad component are higher, implying different zones of formation. The broad Hα component is still coming from the outer, fast-moving region of the ejecta whose emissivity strongly depends on CSM interaction, while the He-Na ID lines mostly form in the slow expanding, spherically symmetric inner ejecta. In order to explain the broad He I line profiles at early times, we may conceive two possible alternatives: either some 56 Ni Figure 10. Hα luminosity evolution of SN 1996al compared with those of SNe 1980K (Uomoto & Kirshner 1986), 1990K (Cappellaro et al. 1995, and the twin SNe 1994aj (Paper I) and 1996L (Paper II). The mean error for the log (LHα) of ±0.10 dex is reported. The horizontal long-dashed (red) line marks the Hα flux emitted by the precursor star. The Hα luminosities here reported are the total Hα flux measurements, which may be slightly different from the sum of the fitting components reported in Tables A3 and A4 and resulting from the deconvolution of the entire line profile. The Hα luminosity on day +4178 is the average between the flux measured from the FORS2 spectrum and that derived from the calibrated Hα image. Figure 11. Evolution of the Hα/Hβ ratio for SN 1996al. The changes of the main properties of Hα are labelled. The latest points, after 1000 days, are essentially lower limits since (at these late phases) Hβ barely detectable. The mean error bar for the Hα/Hβ ratio (±0.5) is shown.
was ejected at high velocity (∼ 10000 km s −1 ), or the He I lines were supported by the interaction of the fast expanding symmetric ejecta with an asymmetric CSM. Since starting from about 2 months past explosion the He I lines become symmetric (the inner geometry of the innermost layers ap-  pears to be symmetric), we believe that the interaction scenario is the most likely.
The disappearance of the He I 5876Å emission line after about day +100 is prima facie at odds with the presence of the pronounced He I 7065Å emission. The clue can be found in the coeval emergence of the Na ID emission. In fact, in a co-moving frame the He I 5876Å photon emitted by the SN ejecta can be redshifted and scattered into the Na ID. As a result we see the Na ID feature instead of the He I 5876Å emission. The mechanism can be explored through a model based on a Monte-Carlo simulation (Figure 17). The underlying model assumes freely expanding spherical ejecta with the He I 5876Å emissivity distribution recovered from the He I 7065Å profile (Figure 17, inset). The adopted He I 5876Å flux is 1.2 times larger than that of He I 7065Å, corre- sponding to the case of collisional excitation dominating over recombination (Almog & Netzer 1989). The optical depth of the Na ID is assumed to linearly fall with increasing velocity from τ (5890) = 20 at the SN center (v = 0), to zero at v = 1800 km s −1 . The relative fractions of the Rayleigh and spherical scattering are determined by the Hamilton parameter E1 which, in the representation of Chandrasekhar (1960), is 1/2 for Na I 5890Å and zero for Na I 5896Å. In the optically thin limit, therefore, the scattering phase function for Na I 5896Å line is spherical, while for Na I 5890Å line the phase function is the superposition of the spherical (1/3) and Rayleigh scattering (2/3). The latter is taken into account for optical depth τ (5890) 1. For τ (5890) > 1, the scattering is assumed to be spherical. This approximately accounts for the scattering spherization during the multiple local scattering. It is reassuring that the simulated spectrum ( Figure 17) from this simple model provides an excellent fit to the observed emission line.
The evidence of the non-local scattering of He I 5876Å photons into Na ID implies that the late time He I line originates in the inner layers of the SN ejecta. Figure 15 shows the evolution of the He 7065Å emission up to almost two years after explosion. While the flux evolution of the line shows a bump between +300 and +500 days (see Table A4 and Figure 18), it becomes narrower with a terminal velocity of about 1200 km s −1 (see Figure 9).
The flux evolution of the He I 6678Å and 7065Å lines reported in Figure 18 shows that the evolution of the line fluxes is consistent with that of the 56 Co decay in the time interval +100 ÷ +300 days, while at earlier and later phases some extra energy from the interaction between the ejecta and CSM is required. Despite the CSM/ejecta interaction is likely the primary source for the SN luminosity, the mechanism responsible for the evolution of the He I emission is far from being clear. On day +449, the He I 7065Å luminosity is ∼ 4 × 10 38 erg s −1 . The He I 5876Å and He I 10830Å lines are typically stronger than the He I 7065Å line by 1.2 and 4−5 times, respectively (Almog & Netzer 1989). The total He I luminosity around day +450 is estimated to be ∼2.8×10 39 erg s −1 , which is larger than the total Hα luminosity of 2 × 10 39 erg s −1 and significantly larger than the luminosity of the Hα Core component. This fact and the low velocity of He-emitting gas suggest, again, that the helium emission originates from the He-rich central zone of the SN ejecta.
Interestingly, the Hα line does not show a significant luminosity enhancement during +300 ÷ +450 days (see Figure 10), when the He I 7065Å flux increases by a factor of 4. This may be understood if the hydrogen ionisation fraction  ( is close to unity, so that the hydrogen emission does not react to the enhanced ionisation that leads to the emergence of strong helium lines. If this is the case, then the total mass of hydrogen in the SN ejecta should be close to the the mass of the ionized hydrogen, i.e., of the order of 0.1 − 0.3 M (see Sect. 3.5.1). Alternatively, if the helium is ionised by the central source, the X-UV radiation could be strongly attenuated already in the He layer, so that the external hydrogen layer would be only marginally affected.
Two possible sources for powering the transient He I emission are conceivable: X-ray emission from the enhanced CSM interaction or energy release due to accretion into the black hole. The latter requires the accretion of ∼10 −7 M during about +200 days and a super-Eddington X-UV luminosity. This could originate in the thermalisation of the Figure 16. Comparison in the velocity space of the profile evolution of the He I 5876Å -Na ID feature (reference wavelength is 5876Å, dotted, red lines) with those of the He I 7065Å feature (short-dashed, black lines) and Hα solid, green lines. In the last phase (+93 days), we also show the profile of He I 7281Å feature (long-dashed, blue line).
accretion disk wind in a scenario of disk outflow interacting with the SN ejecta (Dexter & Kasen 2013).

Photospheric radius
In order to derive the evolution of the photospheric radius and temperature, we perform a blackbody fit to our early time spectra, after correcting for reddening and redshift values reported in Table 3. The fits were done selecting the portions of the spectra free of evident lines (examples of the fits are reported in Figure 6). The fit errors, which depend on a number of parameters, mainly the wavelength extension of the spectra, lines contamination, the uncertainty of the flux calibration and the assumed extinction, have been estimated to be of the order of 250 o K. Adopting the bolometric curve of Figure 4, spherical symmetry and a filling factor of 1, we estimate the evolution of the blackbody radius and the temperature (Figure 19). The computation is extended up to 100 days post explosion. This does not imply that the SN has a clear spheric symmetric photosphere up to that epoch (cfr. next Section).
The photosphere several days after the explosion is relatively hot, with a TBB ∼ 10000 o K and a radius close   Figure 4 and has been tuned to fit the +100 ÷ +300 days flux points. As in Figure 4, a rise time to maximum of 10 days has been assumed. A mean error-bar for the line fluxes is reported.
to 7.4 × 10 14 cm. While the temperature shows a smooth decrease to about 6500 o K at a phase of +100 days after explosion, the photospheric radius remains almost at a constant value of about 6.9 × 10 14 cm up to 80 days after explosion, with a sudden decline to 3.0 × 10 14 cm at phase about +100 ÷ +110 days (see Figure 19). Figure 19. Evolution of the blackbody radius and the temperature in days from the explosion, assumed to occur 10 days before the maximum light. The average temperature error bar of ±250 o K is shown in the upper-right corner, while the mean error bar for radius (±0.3 × 10 14 cm) is shown in the lower-left corner.

Parameters of the progenitor star and environment geometry
In the following, we will sketch a scenario for the SN explosion and its environment that will attempt to account for all available observations. Despite a high expansion velocity of the ejecta (∼10 4 km s −1 , see Table A3) as indicated by the broad wing of Hα and the boxy He I 5876Å emission, the spectra have never shown pronounced broad absorption lines of either hydrogen or other species. This implies that the photosphere is located in the external layer of the SN ejecta. The absence of pronounced broad absorptions is typical of SNe IIn, in which the SN ejecta interact with a dense circumstellar matter (CSM), and is attributed to the high optical depth in the cool dense shell (CDS), that is located at the ejecta/CSM interface (Chugai 2001). It should be emphasised that -in the case of SN 1996al -the CDS is probably patchy because we can detect the broad Hα emission arising in the inner SN ejecta.
The higher resolution spectrum on day +43 (see Figure 7) showing the pronounced broad and narrow Hα components can be used to probe both the SN ejecta and the CSM. The broad component is emitted by the freely expanding ejecta with maximum velocity of 10000 km s −1 , while the narrow line with P-Cygni profile presumably forms in the CSM, which expands with v < 2000 km s −1 . With these numbers, it is easy to calculate the radius of the CDS during the first 1-2 months after the explosion (assuming it occurred 10 days before the maximum light), which spans between 3 − 4 × 10 15 cm. Given the radii calculated in the previous section, and supposing that the photosphere is placed close to the CDS, an upper limit of ∼ 10% for the filling factor of the patchy CDS is derived.

Modelling the early Hα profile
The broad Hα line is modelled with the Monte Carlo technique adopting the recombination emission in the ionised ejecta, with the H + and ne distributions described by a broken power law (ne ∝ v k ) with power index k = k1 in the inner zone (v < v0) and k = k2 in the outer zone. The Thomson scattering with thermal broadening in the ejecta is taken into account, assuming a kinetic temperature of 10 4 o K. Note that Thomson scattering turns out to be a crucial diagnostic probe for estimating the total amount of ionised hydrogen. Since our model does not compute the hydrogen level population, we consider two extreme cases for the Hα optical depth: optically thin Hα line (τ23 1) and optically thick one (τ23 1). The optically thin case is plausible, because the flux in the Balmer continuum seems to be large enough to strongly depopulate the second level. Two sources of continuum are included: the SN ejecta with the emissivity ∝ n 2 e , and the CDS at v = 10000 km s −1 . We do not compute the absolute flux neither in the line nor in the continuum. We neglect the contribution of the line emission from the CDS compared to the continuum radiation bearing in mind that, in the case of the opaque CDS, the contrast of the CDS line emission with respect to the CDS continuum is small. This is indicated by observations of interacting supernovae (e.g. SN 1998S at the age of 10-40 days, cf. Fassia et al. 2001) and supported by arguments based on the model of a thin opaque shell with the estimated line-to-continuum contrast of the order of the thermal-to-expansion velocity ratio (Chugai 2001). Instead, we let line and continuum emissivities of SN ejecta to fit the observations. The contribution of the CDS in the line emission is then neglected whereas the continuum radiation of the CDS is taken into account. The angular dependence of the CDS continuum brightness (Ic) affects the strength of the Hα broad absorption. We assume Ic ∝ 1/ cos θ, where θ is the angle with the normal to the CDS surface. This approximation holds when the CDS is either optically thin or fragmented in such a way that the average number of clouds on the line of sight is less than unity. Strong fragmentation is expected given the fact that CDS is Rayleigh-Taylor unstable (Chevalier 1982). The contribution of the CDS continuum is a free parameter derived from the best fit of the Hα profile. The latter, is determined by the normalised distribution of ionised hydrogen along the radius and by the Thomson optical depth.
We show two cases for the broad component: τ23 1 (no Hα resonance scattering) and τ23 1 (Figure 20). A reasonable fit of the broad component to the observed Hα profile on day +43 is found for the parameters: k1 = 1, k2 = −1.7, v0 = 1800 km s −1 , and Thomson optical depth τT = 2.1. The case without Thomson scattering (dotted line in Figure 20a) demonstrates the crucial role of electron scattering for the formation of the line profile. In the case of the optically thick Hα line, we find that the relative contribution of the ejecta continuum is 10%; for a larger contribution of the ejecta continuum, the broad Hα absorption component becomes prominent. For the optically thin Hα, the calculated recombination continuum in the Case B, can contribute only to 3.3% of the observed continuum. The recovered Thomson optical depth, together with the estimated CDS radius on day +43 and the normalised electron distribution recovered from the Hα profile, suggest a mass of ionised hydrogen of 0.07 M in the ejecta and a Hα luminosity of 2 × 10 40 erg s −1 on day +43. The latter is in excellent agreement with the observational estimate.
We now turn to the narrow line. The peak of the narrow emission on days +23 and +43 suggests the presence of the Hα-emitting circumstellar gas with velocities of ∼ 300 km s −1 . On the other hand, the wings of the narrow emission profile and the blue edge of the narrow absorption indicate CSM velocities as high as ∼ 2000 km s −1 . This is consistent with the expansion law v ∝ r for the CSM with a maximum velocity of ∼ 2000 km s −1 which is used for the Hα modeling. We adopt the external radius for the fast CSM outflow of 5 × 10 15 cm on day +43, keeping in mind that the spectra suggest that this kinematics is maintained at least until day +63. We find that the resonance scattering is a good approximation for the narrow P-Cygni line, although we need to add a net emission which amounts to 8% of the resonant scattering component. We checked also a model where the CSM outflow has a constant velocity. However, such a model reproduced neither the absorption nor the emission components of the narrow line. In the best-fit case shown in Figure  20, the Hα optical depth in the CSM decreases outward from 0.9 at the CDS radius (Rs) down to zero at the boundary radius R b according to τ = 0.9[(r − Rs)/(R b − Rs)] 1.6 .
The kinematics v ∝ r followed by the observed narrow Hα poses a question about its origin. Generally, this sort of the kinematics combined with large value of the terminal velocity (∼ 2000 km s −1 ) suggests a mass ejection via shock wave. At present, we are not able to propose a reliable mechanism for this vigorous mass loss. Beside the stellar shock, the common envelope phase in a binary system could also play a role.

Modelling the light curve
The fast initial (t < +100 days) drop of the bolometric light curve (see Figure 4) strongly indicates that the diffusion time in the SN envelope is small, which implies low-mass (∼1 M ) ejecta. At first glance, this is at odds with the high mass of the SN progenitor suggested by the pre-explosion Hα luminosity (∼10 37 erg s −1 , see Table 4). In fact, in the case B recombination, the pre-explosion Hα requires a rate of ionising radiation of ∼10 49 photons s −1 . To sustain this ionisation rate, we need a hot progenitor star with bolometric luminosity log L/L 5.4 and radius R ∼ 10R (see discussion in Sect. 3.1). According to atmosphere models (Simón-Díaz & Stasińska 2008), the above precursor may have been a 25 M ZAMS star at the final evolutionary stage (Woosley, Heger, & Weaver 2002). The apparent controversy between the SN light curve and pre-SN Hα constraints can be solved if one assumes that the 25 M ZAMS progenitor had lost almost the whole hydrogen envelope, with most of the helium core falling back onto the collapsing core making a black hole (see models in the 25 − 30 M mass range of Heger et al. 2003).
For our fiducial model we assume that the 25 M progenitor (Woosley, Heger, & Weaver 2002) ends up as a stripped pre-SN star with a 8.3 M helium core, and a 0.3 M envelope, with a hydrogen mass fraction of 0.5. The adopted pre-SN radius is R = 10 R constrained by the preexplosion Hα luminosity. To construct the stripped pre-SN configuration, we scaled the mass and the radius of the hydrogen envelope of the 25 M progenitor to the required values. Using a radiation hydrodynamics code (Utrobin 2007;Utrobin et al. 2015), we analysed the outcome of the explosion which is initiated by a supersonic piston applied to the bottom of the stellar envelope at the boundary of the 2.1 M central core, which is then removed from the computational mass domain and assumed to collapse into a neutron star. We computed a set of SN models with explosion energies of (0.1 − 1.0) × 10 51 erg, ejecta masses of 0.45 − 2.1 M , and kinetic energies of (0.2 − 4.5) × 10 50 erg. Note that the explosion energies released at the boundary of the collapsing core are in the plausible range for the neutrino-driven explosion mechanism (Janka 2012). The light curves of these models have a very narrow shock breakout peak (∼10 −4 days), a short overall duration (∼10 days), and a luminosity of <10 42 erg s −1 (after the shock breakout peak), while the expansion velocity at the boundary is large (>10 4 km s −1 ). These properties are in agreement with the absence of broad absorption lines in the spectra which, in turn, suggest a small energy contribution of the intrinsic SN luminosity to the bolometric light curve, and instead support the circumstellar interaction as the dominant source powering the SN luminosity.
The bolometric luminosity powered by the CS interaction is calculated in a thin shell approximation (Chevalier 1982;Chugai 2001). The model suggests that freely expanding SN ejecta collide with dense CSM. The interaction proceeds via forward and reverse shocks with the CDS at the contact boundary between SN and CSM. The CDS radius Rs and velocity vs are recovered from the numerical solution of the equations of motion and mass conservation. The forward shock velocity is vs − vw (vw is the wind velocity) and the reverse shock velocity is R/t − vs. The post-shock temperature is set by the shock velocity assuming an isothermal (Te = Ti) strong shock, while the post-shock density is set to be four times the pre-shock density. The X-ray luminosity of the j-th shock (j = 1 for the forward shock and j = 2 for the reverse shock) is Lj = 2πR 2 s v 3 j ηj, where vj is the shock velocity, ηj = t/(t + tc,j) is the cooling efficiency and tc,j is the post-shock cooling time. The optical bolometric luminosity is calculated assuming that the full X-ray luminosity of both shocks is absorbed by the cool material of the CDS, the un-shocked SN ejecta and the CSM.
We consider two cases: (i) spherically symmetric CSM with velocities v = v0(r/r0) for r < r0 and v = v0 = 2000 km s −1 for r > r0 ≈ 5 × 10 15 cm and (ii) anisotropic CSM with a dense equatorial disk and a rarefied, more symmetric flow. This latter outflow presumably does not contribute significantly to the interaction luminosity, although it is responsible for the narrow Hα absorption. The equatorial CSM is assumed to have a low expansion velocity: we set it to be 100 km s −1 (the result does not depend on the exact value of vw). The CSM anisotropy is set as the fraction of a sphere occupied by the equatorial CSM, G = Ω(r)/4π. This factor multiplied by the luminosity computed for the spherical CSM gives the luminosity in the anisotropic case. Indeed, the luminosity and its temporal evolution constrain the CS density and its radial distribution. We find that the requirement of a rapid deceleration suggests a low mass and low energy supernova ejecta in order to minimise the ejecta momentum. A search for the best fit model in an extended parameter space leads us to conclude that one cannot describe simultaneously the light curve and the low expansion velocity in the framework of a spherically symmetric model. This is the primary reason to invoke an asymmetric CSM. We do not pursue the goal to find the best fit model for the CSM using a minimisation procedure. Our best-fit model should be considered rather as a reasonable possibility.
The model light curves and the expansion velocity of the CDS for the isotropic and anisotropic CSM cases are shown in Figure 21. As parameters for the hydrodynamical model we adopt an ejecta mass M = 1.15 M and a kinetic energy E = 1.6×10 50 erg. The ejecta consist of 0.15 M of hydrogen and nearly 1.0 M of helium. The rest of the He core is assumed to fall back, thus possibly making a black hole with a mass of 7−8 M . The SN explosion produces the brief shock breakout peak followed by a short quasi-plateau (∼12 days) with a luminosity of ∼10 41 erg s −1 . This contributes negligibly to the SN luminosity, which is instead dominated by CS interaction (Figures 21, panels a and c). The modelled interaction luminosity for both the spherical and the anisotropic CSM well fits the observed light curve (Figures 21, panels a  and c), whereas the expansion velocities of the CDS at late times are quite different: the spherical model predicts deceleration down to 3000 km s −1 , while the CDS for the equatorial CSM is decelerated down to 1300 km s −1 (Figures 21,  panels b and d). In that respect, the non-spherical model is preferred, given the similarities of the model with the velocity evolution of Hα and the He I lines shown in Figure  9. The total mass of the non-spherical CSM in the range of r < 3×10 16 cm is 0.13 M . The geometric factor is G = 0.02 at the large radii (Figure 21d, inset) which corresponds to a polar angle width of the equatorial disk of 23 • .
The scenario with an equatorial CSM predicts a doublehorn line emitted in the interaction zone. Indeed, at t > +142 days the Hα shows Blue and Red bumps with radial velocities of about −1100 km s −1 and 1550 km s −1 , respectively (see Table A4). The velocity and intensity differences between those peaks indicate that the structure of the CSM deviates from axial symmetry.

Evidence from narrow absorption lines
The narrow absorptions of the Hα,Hβ,and Fe II (mult. 42) lines are shallow in the higher resolution spectra of days +23 to +26 (see Sect. 2.4). Moreover, Hα and Hβ have similar relative depths, A ≈ 0.1. Given a theoretical ratio of optical depths Hβ/Hα = 1/7, the facts that the Hβ and Hα absorptions have similar depths and are both shallow imply that these lines are saturated, but only about 10% of the continuum source is covered by line absorbers. The spherical component of the CSM responsible for the narrow absorptions is therefore clumpy and consists of clouds optically thick in the Hα and Hβ lines embedded in the optically thin inter-cloud gas. The interaction with the clumpy spherical CSM might contribute to the central (Core) Hα emission component (t > 142 d). The fact that this emission component is redshifted by 300−400 km s −1 could be attributed to the asymmetry of the CSM in a sense that the far side of the quasi-spherical CSM is more dense and mostly contributes to the Core Hα emission.

The progenitor scenario
The dominant role of the continuum originated by the CSMejecta interaction and the significant deceleration of the SN ejecta indicate lower ejecta mass and explosion energy than in normal core-collapse SNe. We have shown in the previous Sect. that ejecta with Mej ≈ 1.15 M and E ≈ 1.6×10 50 erg are consistent with the observations. The low hydrogen mass (∼ 0.15 M ) and the high He I luminosity around day +450 imply that the ejecta are made up primarily of helium (∼ 0.8 M ).
The CSM close to the exploding star has an overall spherical distribution (see Figure 22), but with some density enhancement in an oblate structure to explain the asymmetric, early He I line profiles. The inner radius of the equatorial disk is r ≈ 3 × 10 15 cm, while its external limit is > 5 × 10 17 cm (> 0.15 pc), since the ejecta are still interacting with the disk at the time of our last observation (+5542 day). The disk is embedded in a more spherically symmetric, clumpy CSM. The mass of the CSM in the anisotropic model, including the inner mostly spherical zone is 0.13 M , within a distance r < 3 × 10 16 cm. This value is very similar to the amount of CS mass found in the twin SN 1994aj (Paper I) and about ten times more than in SN 1996L (Paper II). According to this scenario, and given the progressive quenching of the Hα Core and Red components, we may suppose that both the receding part of the disk and the quasi-symmetric, clumpy CSM get shallower and shallower in density with time, or that some obscuring dust is forming mostly inside the shocked ejecta (or a combination of the two).
The presence of dense CSM around the progenitor star of SN 1996al is also supported by the recovery of the progenitor in an archive Hα image, from which a Hα luminosity of logLHα ∼ 37.28 is derived. This Hα luminosity implies a massive ∼ 25 M ZAMS progenitor star (see Sect. 3.1), which has lost almost the entire hydrogen envelope during its evolution.
It is interesting to note that such a highly asymmetric geometric configuration derived for the SN 1996al CSM has been observed around very massive evolved stars such as MN18 of Large Magellanic Cloud (LMC), and in other blue The inset shows the CS density distribution, where ρ and r units are g cm −3 and cm, respectively. Panel b: the CDS velocity (thick line) and the boundary velocity of the un-shocked SN ejecta (thin line). For comparison, the observed velocities of the He I 7065Å are also plotted (solid, red squares connected with a dotted, red line) The inset shows the CDS radius. Panel c: the same as in panel a but for the non-spherical case. Panel d: the CDS velocity (thick line) and the boundary velocity of the un-shocked SN ejecta in the equatorial plane (thin line). For comparison, the observed velocities of the He I 7065Å are also plotted (solid, red squares connected with a dotted, red line). The inset shows the geometric factor describing the angular width of the equatorial disk (see text).
In this scenario, the absence of [O I] 6300-6364Å emission lines at any phase (see Figure 23) is naturally explained since most of the oxygen layer has fallen onto the black hole 5 . 5 It is interesting to note that a similar scenario (low energy, 4 × 10 50 erg; massive, 26 M , progenitor; substantial fallback of material on to the collapsed core) has been proposed to explain the physical characteristics of SN 1997D (Turatto et al. 1998;Benetti et al. 2001). However, this high mass scenario has been questioned by Chugai & Utrobin (2000) who, instead, proposed a low mass progenitor for SN 1997D. In the case of SN 1996al the key parameter that point towards a high mass progenitor scenario is the pre-explosion Hα luminosity.

CONCLUSIONS
Extensive photometric and spectroscopic observations of SN 1996al along 15 years of its evolution are presented. The light curve indicates that the SN is of Type IIL, reaching peak absolute magnitudes of MB ∼ −18.6 and MV ∼ −18.2. However, a closer inspection shows that the decline rate changes with time, and this is correlated with significant changes in the spectral appearance. The complex SN evolution is very likely a consequence of the primary role played by the ejecta-CSM interaction during all phases of the SN life.
The spectra obtained soon after the maximum light show narrow P-Cygni Balmer lines superimposed on broader components with asymmetric profiles. In addition, broad and asymmetric He I lines in emission are detected. At these early phases, the broad H component and the He lines form in the same outer region of the ejecta. By day +142, the Hα profile dramatically changes: the narrow emission component disappears, and Hα splits in three components of comparable widths, that remain at constant wavelengths afterward. This composite profile remains visible over 15 years, although the relative strengths of the three components change with time.
The He lines behave differently. After day +100, they become significantly narrower ( 1000 km s −1 ) and symmetric, a clear indication that they form in a different region of the ejecta. Later on, between day +300 and +600, they experience a sudden increase in flux. Then, the flux of the He I line decreases again, and by day +700, these lines fade below the detection threshold.
We have also recovered information on the progenitor star by inspecting a pre-explosion archive image in the Hα narrow band. A luminous source is detected at the location of SN 1996al, showing a total Hα luminosity of 1.9 × 10 37 erg s −1 , which favours a massive (25 M ZAMS) progenitor star for SN 1996al. This, along with the SN parameters derived by the modelling, suggests that the progenitor was stripped of a large fraction of the hydrogen envelope in the latest stages of its life which ended with the collapse of a 7 − 8 M helium core.
The peculiar evolution of the H and He I lines is likely a consequence of the geometry, the non-uniform density profile of the CSM surrounding SN 1996al and the fallback onto the compact remnant. Observational evidences favour an highly structured CS scenario characterised by inner, high-density material with oblate structure plus an equatorial ring extended over 0.15 pc, both embedded in a more spherically symmetric, lower-density but clumpy CSM. The total mass of the CSM is modest, being 0.13 M within a distance 3 × 10 16 cm. Models show that the luminosity of SN 1996al is best sustained by SN ejecta-CSM interaction. The inferred ejected mass is relatively small (∼ 1.15 M , ∼ 0.15 M of which is hydrogen, the residual fraction is helium) and is expelled in a low kinetic energy explosion (∼ 1.6 × 10 50 erg). The amount of ejected 56 Ni (if any) expelled in the explosion is constrained to be 0.018M . The modest ejecta and  Fesen et al. (1999), while that of SN 1998S is from Mauerhan & Smith (2012). 56 Ni mass are explained with a massive fallback of material into the compact remnant, producing a 7 − 8 M black hole.

ACKNOWLEDGMENTS
We thank the anonymous referee for very useful comments and suggestions. We thank T. Augusteijn, M. Birlan, J. F. Le Borgne, S. Covino, P. Rosati and G. C. Van de Steene for giving us part of their telescope time; P. Fouqué for supplying the DENIS photometry; and M. Riello for performing some of the observations. We are in dept with N. Richardson for providing unpublished spectra of LBVs. We thank R.    Table A3. Spectral lines parameters as derived from spectra of SN 1996al up to 100 days. Table A4. Spectral line parameters as derived from spectra of SN 1996al observed after 100 days.