Astronomical Science The Antares Emission Nebula and Mass Loss of A Sco A Dieter Reimers, Hans-Jürgen Hagen, Antares sible a 2-D mapping of the nebula with Robert Baade, Kilian Braun A B long-slit spectroscopy. We used UVES (all Hamburger Sternwarte, Universität with a 0.4 slit and a resolution of 80 000 Hamburg, Germany) 2.7 and covered the nebula with 23 long-slit positions (Figure 1). Exposure times were between 50 and 100 s (the Antares neb- The Antares nebula has been known as ula is bright!); the typical seeing was 0.6 a peculiar [Fe II] emission nebula, ap- which means that the angular resolution parently without normal H II region lines. corresponds roughly to the slit steps as Long-slit VLT/UVES mapping shows shown in Figure 1. that it is an H II region 3 in size around the B type star A Sco B, with a Balmer Quite unexpected was that all spectra – line recombination spectrum and [N II] even 10 from the M supergiant – are lines, but no [O II] and [O III]. The reason completely corrupted by scattered light for the non-visibility of [O II] is the low from the M supergiant. This was a sur- N electron temperature of 4 900 K, while prise, also to the UVES team at ESO. ESO [NII] is seen because the nitrogen abun- finally made a test placing the UVES slit dance is enhanced by a factor of three 10 from a bright star – with the same W due to the CNO cycle. We derive a result as in our spectra. Apparently, light –6 mass-loss rate of 1.05 ± 0.3 × 10 M/yr from the bright star is scattered by the slit for the M supergiant A Sco A. The [Fe II] Figure 1. Geometry of the Antares system together jaws back to the slit viewing camera and lines seem to come mainly from the with the UVES spectrograph slit sizes and positions from there again into the spectrograph shown to scale. edges of the H II region. slit. This should be a warning for anybody using UVES on faint targets close to bright UVES/VLT spectra, the Antares nebula sources, e.g. a disc, around a bright star. An iron rich nebula? puzzle could be solved. Did this mean that our spectra were use- The story of the Antares nebula began in The most extensive previous study was less? It meant that, due to the ubiquitous 1937 with the finding of O. C. Wilson and that by Swings & Preston (1978) based on scattered light, no general background R. F. Sanford at the Mount Wilson Ob- high-resolution long-slit photographic reduction was possible and no absolute servatory that the spectrum of the B-type spectra taken with the Mt. Wilson 100 line fluxes could be deduced. Data reduc- companion of Antares, itself an M super- and Palomar 200 telescopes. They found tion required an enormous load of extra giant, showed forbidden emission lines that the ‘[Fe II] rich nebula’ is strong roughly work. In brief, the scaled M star spectrum of Fe II. Later Otto Struve showed that the 3.5 around the B star and is surrounded was used as a background template. [Fe II] lines extend ~ 2 beyond the see- by a zone of weaker lines which may ex- The template was fitted to the spectrum – ing disc of the B star. Struve & Zebergs tend in a NW-SE direction up to 15. How- allowing for an offset and a scaling – (1962) describe extensively the Antares ever, they neither observed Balmer re- outside a region of ± 30 km/s of each de- nebula and conclude: “It is strange that combination lines nor any classical H II tected emission line and was finally the nebulosity around the B type com- region lines from ions like O II, O III, S II, and subtracted from the data. ponent shows only emission lines of iron N II. Therefore the nature of the Antares and silicon, but not those of hydrogen. nebula could not be understood. In a first The final results after this elaborate data The enormous abundance of hydrogen in quantitative study of the Antares system reduction – each of the roughly 90 de- all known gaseous nebulae and the con- it could be shown that the B star creates tected emission lines had to be treated ditions of excitation and ionization result- an H II region within the wind of A Sco A individually in each of the spectra – are ing from the radiation of a B type star and that the [Fe II] lines must be formed in 2-D brightness distributions (location on would render it almost inexplicable for a this H II region (Kudritzki & Reimers, 1978). the slit versus velocity) as a function of gas of normal composition to show only Later the H II region was detected as an the location relative to the M supergiant, the iron and silicon lines in emission. optically thin radio emitter with a diameter as shown for the HA line in Figure 2 and This gave rise to the hypothesis that the of ~ 3(Hjellming & Newell, 1983). for [Fe II] 4 814.55 Å in Figure 3. material in the vicinity of Antares is metal rich, in turn leading to the supposition that the envelope may be composed of UVES 2-D spectroscopy A peculiar H II region? solid substances such as meteors that have become vaporized and are excited With the construction and operation of The main puzzle of the Antares Nebula by the radiation of the hot B star. This UVES at the VLT it was obvious that pro- has been resolved by our UVES data. The hypothesis is admittedly improbable and gress in our understanding of the kine- Balmer recombination lines expected may have to be abandoned in the light matics of the Antares nebula should be for an H II region: HA, HB ... up to H10 are of future work.” We may add here that the possible with its high spectral resolution present, and their geometrical extent mystery has lasted 45 years until, with and high pointing accuracy, making pos- on the sky is virtually identical to that of The Messenger 132 – June 2008 33 Astronomical Science Reimers D. et al., The Antares Emission Nebula and Mass Loss of A Sco A 0.9 1.4 1.9 2.4 2.9 0.9 1.4 1.9 2.4 2.9 3.4 3.9 4.4 4.9 5.4 3.4 3.9 4.4 4.9 5.4 10 8 8 ) ) ( ( 6 . 6 . s s o o P 4 P 4 2 2 0 30 0 0 30 –20 0 –40 –40 –20 Velocity (km/s) Velocity (km/s) 5.9 6.4 6.9 7.4 7.9 5.9 6.4 6.9 7.4 7.9 Figure 2. HA 2-D brightness distributions (location Figure 3. [Fe II] 4814.55 Å 2-D brightness distribu- on the slit versus velocity) as a function of spectro- tions as a function of spectrograph slit position (see graph slit position relative to A Sco A. Offsets from caption of Figure 2). Notice the double structure the position of A Sco A are shown in arcsec. Notice is probably formed by the front (blueshifted) and rear that due to the M star wind expansion, the velocity (redshifted) edges of the H II region carried by the coordinate corresponds to the spatial depth (per- wind to large distances from the system. pendicular to the plane of the sky). The density in the plots has been rescaled in order to provide maxi- mum visibility. The true line fluxes vary by a factor of 150 between 1.9 and 7.9. the radio emission seen with the VLA normally prominent in H II region spectra regions, are not? There are two reasons: III II (Figure 4). So why have previous observa- like [O ] 4959/5 007 Å, [O ] 3726/3729 a low electron temperature Te and an en- tions not shown the hydrogen lines? Å, [N II] 6 583/6548 Å and [S II] 6716/6731 hanced N/O ratio. At first, the electron The reason is probably that the strongest Å had not been detected. Our UVES temperature Te must be so low that the lines HAand HBsuffered, in the Mt. Wil- spectra allow a more critical assessment [O II] lines are not excited and below the son and Mt. Palomar photographic spec- than earlier photographic spectra. The an- detection limit. We have modelled the tra, from seeing-dominated contamination swer is that we do see faint [N II] 6 583 Å H II region created by A Sco B in the wind by the M supergiant and the signal-to- and 6 548 Å lines, and that none of the of the M supergiant using the Cloudy noise of the photographic spectra was other lines is present. [N II] 6 583 Å follows code (version 07.02.00, last described by never sufficient to enable an efficient closely the HA distribution on the sky Ferland et al., 1998). Since the B star, with background subtraction. shown in Figure 2. Since it is close to HA, an effective temperature of 18 200 K, is we can measure the HA/[NII] 6 583.4 Å relatively cool, the resulting electron tem- The strongest lines are of [Fe II] such as ratio as a function of the location in the perature is below 5 000 K which does not 4287.4 Å, 4 359.34 Å and 4814.55 Å.