---
abstract: |
    This paper is the first one in a series investigating the properties of the S stars belonging to the Henize sample (205 S stars with $\delta<-25^\circ$ and $R<10.5$) in order to derive the respective properties (like galactic distribution and relative frequencies) of intrinsic (i.e. genuine asymptotic giant branch) S stars and extrinsic (i.e. post mass-transfer binary) S stars. High-resolution (R=30000 to 60000) spectra covering the range $\lambda\lambda4230-4270$Å  have been obtained for 76 S stars, 8 M stars and 2 symbiotic stars.

    The $\lambda4262$Å  and $\lambda4238$Å  blends involving a line were analysed separately and yield consistent conclusions regarding the presence or absence of technetium. Only one ‘transition’ case (Hen 140 = HD 120179, a star where only weak lines of technetium are detectable) is found in our sample. A resolution greater than R $=30\,000$ is clearly required in order to derive unambiguous conclusions concerning the presence or absence of technetium. The Tc/no Tc dichotomy will be correlated with radial velocity and photometric data in a forthcoming paper.
author:
- 'S. Van Eck'
- 'A. Jorissen [^1]'
date: 'Received date; accepted date'
subtitle: 'I. The technetium dichotomy[^2] '
title: The Henize sample of S stars
---

\#1\#2[BD[$#1$]{}$^\circ$[\#2]{}]{} \#1\#2[[$#1$]{}$^\circ$[\#2]{}]{}

Introduction
============

S stars have been identified as a class of peculiar red giants by Merrill ([@Merrill22]). Basically, the S stars emerge as a sequence parallel to the M stars as far as temperature is concerned, but with enhanced abundances of s-process elements. The chief observational difference between M and S spectra is the presence of ZrO bands in the latter. The s-process overabundances in S stars are explained in the framework of the Thermally Pulsing Asymptotic Giant Branch (TPAGB) evolution. Thermal instabilities (called [*thermal pulses*]{}) that affect the helium-burning shell of these stars have two important consequences: they provide the proper environment for the nucleosynthesis of s-process elements, and they trigger an envelope response (the [*third dredge-up*]{}) that allows the s-process elements and carbon to be brought to the surface \[see e.g. Mowlavi ([@Mowlavi]) for a review\].

Technetium is an s-process element with no stable isotope that was first identified in the spectra of some M and S stars by Merrill ([@Merrill52]). $^{99}$Tc, with a laboratory half-life of $t_{1/2} = 2.13\times 10^5$ yr, is the only technetium isotope produced by the s-process. The high temperatures encountered during thermal pulses strongly shorten the effective half-life of $^{99}$Tc ($t_{1/2}\sim 1$ yr at $\sim 3\times 10^8$ K, Cosner et al. [@Cosner]) but the large neutron densities at these high temperatures more than compensate the reduction of $^{99}$Tc life-time (Mathews et al. [@Mathews]) and enable a substantial technetium production. Third dredge-up episodes then carry technetium to the envelope, where it decays steadily at its terrestrial rate of $t_{1/2}=2.13 \times 10^5$ yr. Starting from an abundance corresponding to the maximum observed one, technetium should remain detectable during $1.0-1.5\times 10^{6}$ yr (Smith & Lambert [@Smith]). If the dredge-up of heavy elements occurs after each thermal pulse (occurring every $\sim 1-3\times 10^{5}$ yr), virtually all s-process enriched TPAGB stars should exhibit technetium lines.

However, Straniero et al. ([@Straniero]) advocated that the s-process nucleosynthesis mainly occurs during the interpulse. When technetium is engulfed in the subsequent thermal pulse, it will decay at a fast rate because of the high temperature, and will not be replenished if there is no neutron source operating within the pulse itself. The conclusion that s-process enriched TPAGB stars should necessarily exhibit technetium would then be challenged. Nevertheless, all the S stars identified as TPAGB stars by Van Eck et al. ([@Van; @Eck98]) thanks to the HIPPARCOS parallaxes turned out to be Tc-rich. A survey of technetium in a large sample of S stars like the Henize sample may be expected to provide further constraints on the s-process environment in AGB stars (e.g. interpulse s-process versus thermal-pulse s-process, thermal-pulse duration and temperature versus $^{99}$Tc half-life).

Not all S stars exhibit Tc lines though (Little-Marenin & Little [@Little-Marenin], Little et al. [@Little]), but technetium-poor stars (also called [*extrinsic*]{}, as opposed to technetium-rich, [*intrinsic*]{} S stars) are currently believed to emerge from a totally different evolutionary history: because they are members of binary systems (Brown et al. [@Brown], Jorissen et al. [@Jorissen93], Johnson et al. [@Johnson], Jorissen et al. [@Jorissen]), they rather owe their chemical peculiarities to the accretion of s-process-rich matter from their companion (formerly a TPAGB star, now an undetected white dwarf). They are technetium-poor, because enough time has elapsed for the technetium to decay since the mass transfer event. The $^{99}$Tc half-life is indeed much shorter than any stellar evolutionary timescale (but the TPAGB). Such a polluted giant star will be classified either as a G or K giant with enhanced heavy elements (i.e. as a barium star) or, if it has cooled enough for the ZrO molecular bands to appear, as a technetium-poor S star.

Besides technetium detection, several spectroscopic criteria (of various efficiencies) aiming at distinguishing extrinsic from intrinsic S stars have been mentioned in the literature \[e.g. oxygen isotopic ratio (Smith & Lambert [@Smith90]), presence of the $\lambda 10830$ line (Brown et al. [@Brown]), zirconium isotopic ratio (Busso et al. [@Busso])\], but technetium detection appears to be, by far, the most secure and tractable way to unmask extrinsic S stars.

This unmasking operation is crucial when deriving fundamental stellar quantities such as the third dredge-up luminosity threshold (commonly measured as the minimum luminosity of S stars). Evolutionary timescales of TPAGB stars can be strongly in error if the considered star samples are polluted by non-AGB, mass-transfer S stars.

We therefore decided to study a large and properly defined sample of S stars in order to disentangle the two sub-families and to study their respective characteristics. The Henize sample of S stars (Henize [@Henize], as listed in Stephenson [@Stephenson]) consists of 205 S stars south of $\delta=-25^\circ$ and brighter than $R=10.5$. Radial velocity data, low- and high-resolution spectroscopy, as well as Geneva photometry have been collected over several years. The present paper deals with high-resolution technetium spectra for 72 Henize stars. Some additional K, M and symbiotic stars data are also presented. Results concerning binarity and photometry, as well as the global analysis of the different data sets, are postponed to a forthcoming paper.

Observations and reduction
==========================

Instrumental set-up {#Sect:setting}
-------------------

The high-resolution spectra used in the present study were obtained during several runs (1991-1998) at the European Southern Observatory, with the Coudé Echelle Spectrometer (CES) fed by the 1.4m Coudé Auxiliary Telescope (CAT). The 1991-1993 runs were performed with the short camera (f/1.8) and CCD \#9 (RCA SID 503 thinned, backside illuminated, $1024 \times 640$ pixels of $15~\mu$m), whereas the long camera (f/4.7) and CCD \#38 (Loral/Lesser thinned, backside illuminated, UV flooded, $2688 \times 512$ pixels of $15~\mu$m) were used during the 1997-1998 runs. Details on these configurations can be found in Lindgren & Gilliotte ([@Lindgren]) and Kaper & Pasquini ([@Kaper]). The resolution ranges from 0.14Å  (R=30000) to 0.07Å  (R=60000) for a central wavelength of 4250Å. The spectra approximately cover the wavelength range $\lambda\lambda$ 4230-4270Å.

Stellar samples
---------------

The observed stars are a subset of the sample of 205 S stars collected by Henize ([@Henize]) from his objective-prism survey (with a dispersion of 450Å/mm at H$\alpha$) of ZrO stars south of $\delta = -25^{\circ}$ and brighter than $R=10.5$. Given the limitations on the CAT pointing and on the detectors sensitivity, only stars with $\delta > -75^{\circ}$, $V<11$ and $B-V<2$ (translating into 70 objects) could be observed in a reasonable amount of time, i.e. less than 1h30 per star. A few bright redder stars could also be observed (but see the discussion on SC stars in Sect. \[Sect:SC\]). A sample of bright M stars with an excess at $60\mu$m \[indicative of a possibly detached dust shell; see Zijlstra et al. ([@Zijlstra])\], as well as the two symbiotic stars RW Hya and SY Mus, and some radial-velocity standards have also been observed. Three non-Henize S stars from our radial-velocity monitoring (Udry et al. [@Udry]) have been included as well. The log of the observations, including the instrumental setting, is given in Table \[Tab1\] and, for Henize stars, in Table \[Tab2\].

Data reduction and S/N ratio
----------------------------

The CCD frames were corrected for the electronic offset (bias) and for the relative pixel-to-pixel response variations (flat-field). Wavelength calibration was performed from thorium lamp spectra taken several times per night. An optimal extraction of the spectra was performed according to the method of Horne ([@Horne]). The whole reduction sequence was performed within the ‘long’ context of the MIDAS software package.

The signal-to-noise (S/N) ratio was estimated for each spectrum in the following way: three S/N values were computed for the three best exposed CCD lines (along the dispersion axis), in the neighborhood of the spectral region of interest (either 4262Å  or 4238Å). These three S/N values were then combined according to Eq. 17 of Newberry ([@Newberry]). When the exposure time on a given star has been split in two (in order to reduce cosmics detrimental effect), the final S/N ratio was computed using Eq. 18 of Newberry ([@Newberry]). The degradation of the S/N ratio due to flat-field correction has not been taken into account, since flat-fields have little degrading effect for the low S/N values under consideration. The S/N ratio values are listed in Tables \[Tab1\] and  \[Tab2\] for each target star. Because of the CCD spectral response, the S/N ratio near 4238Å  is systematically lower than the one near 4262Å.

Analysis
========

Fit of the technetium blends {#Sect:Analysis}
----------------------------

The three strong resonance lines of are located at 4238Å, 4262Å  and 4297Å, with intensity ratios of 3:4:5. All three lines are severely blended (Little-Marenin & Little 1979, their Table III). With the adopted instrumental configurations, a single exposure spans 35 to 50Å; it is thus possible to observe simultaneously the 4238Å  and 4262Å  lines. In this analysis we follow the guidelines provided in the landmark paper of Smith & Lambert ([@Smith]) and therefore concentrate on the most useful 4262Å  line, while the 4238Å  line is used as an independent confirmation.

Fig. \[Fig:figTc62\] shows examples of spectra in the 4262Å  region for an M3-4 giant (HD 73341) and for seven S stars (four being technetium-poor: Hen 3, 187, 31, 7 and three technetium-rich: Hen 140, 39 and 202 = $\pi^1$ Gru). It can be seen that the Tc $\lambda 4262.270$Å  line is blended with two features; the bluest includes primarily (4262.050Å) and (4262.087Å), and the reddest (4262.373Å) (see Fig. \[Fig:figTc62\]). A weaker contribution of at 4262.228Å, almost on the top of the line, may also be present. These composite features are much weaker than the resonance line at its maximum strength; moreover, the - blend and the line are 0.18Å  apart. Therefore the shape and location of the - (-) blend (hereafter called $X_{4262}$ feature) clearly depends on whether it contains the technetium line or not. Quantitatively, the minimum of a gaussian fitted to a Tc-containing $X_{4262}$ feature is shifted redward by $\sim0.14$Å  with respect to the minimum of a gaussian fitted to a no-Tc $X_{4262}$ feature; such a shift is easily detectable on our spectra (compare Hen 3 or Hen 7 with Hen 39 on Fig. \[Fig:figTc62\]).

In practice, each spectrum has been rebinned to zero-redshift in the following way: 10 nearby ($\le 5$Å  on either side) apparently unblended stellar features with unambiguous identification, are adopted as wavelength standards. Gaussian profiles are fitted to these lines and provide a mean redshift. The wavelength of the $X_{4262}$ feature is then computed as the minimum of a gaussian centered on the $X_{4262}$ feature of the redshift-corrected spectrum. Typical uncertainties on the $X_{4262}$ wavelength amount to 0.013Å  for technetium-poor stars and 0.020Å  for technetium-rich stars (as derived from the standard deviation on the mean redshift).

The same method is applied to the 4238Å  technetium line, where the CH- blend has been taken as the $X_{4238}$ feature (Fig. \[Fig:figTc38\]). Typical uncertainties on the $X_{4238}$ wavelength are slightly larger (0.023Å  for Tc-rich S stars and 0.025Å  for Tc-poor S stars) because of the lower S/N ratio and the stronger blending at 4238Å. Results are listed in Tables \[Tab1\] and  \[Tab2\]; the $X_{4238}$ and $X_{4262}$ features always yield consistent results regarding the absence or presence of technetium, except for Hen 140 (=HD 120179).

This star is indeed unique in having very weak technetium features (see Figs. \[Fig:figTc62\] and \[Fig:figTc38\]). A second spectrum, taken 3.5 years later, is almost identical to the one displayed in Fig. \[Fig:figTc62\] and \[Fig:figTc38\]. The blind application of gaussian fitting to the $X_{4238}$ feature of Hen 140 yields a central wavelength that would qualify it as Tc-poor; however, the extreme weakness of the pseudo-emission separating the LaII-CH blend from the SmII line as seen in Hen 140 (Fig. \[Fig:figTc38\]) is unusual for Tc-poor stars, and suggests the presence of a weak technetium line, as confirmed from the appearance of the $X_{4262}$ feature. We therefore believe that Hen 140 is the unique example in our sample of an S star with very weak Tc lines.

In all the other cases, technetium (non-) detection relies on the location of the minimum of a gaussian fitted to the $X_{4262}$ (or $X_{4238}$) blend.

Are there, with this method, risks (i) to misclassify as Tc-rich a truly Tc-poor star, and (ii) to misclassify as Tc-poor a truly Tc-rich star? We show in the remaining of this section that both risks are most probably non-existent in the present study.

Error (i) could, in principle, affect very luminous Tc-poor stars, because their large macroturbulence would broaden their $X_{4262}$ feature, which could then possibly mimick a Tc-rich feature.

In order to test this hypothesis, gaussian filters of different widths have been applied to Tc-poor spectra, so as to make their line widths comparable to those of the stars classified as Tc-rich. This simulation clearly shows that even the largest macroturbulence value observed in our sample (T Ceti) is not large enough to make truly Tc-poor stars appear as Tc-rich from the broadening of their $X_{4262}$ feature. That conclusion is even more stringent when considering the $X_{4238}$ feature. However, this risk cannot be excluded for very luminous stars (class I or II) if observed at lower resolution (R $< 30\,000$).

Error (ii) could, in principle, occur for stars displaying a ‘weak technetium line’ (weaker than the Tc lines of Hen 140 discussed above) with an intensity not large enough to shift the $X_{4262}$ blend redward from the Tc-poor wavelength. In fact, some stars in our sample exhibit an ‘ambiguous’ $X_{4262}$ blend, in the sense that the pseudo-emission located between the - lines and the line becomes very weak or even disappears, mimicking a ‘weak technetium line’ (a typical example is Hen 7 on Fig. \[Fig:figTc62\]). Such a star is classified by our method as technetium-poor, for the minimum of the $X_{4262}$ blend remains unchanged with respect to the no-Tc cases.

In fact, all intermediates exist between the ‘unambiguous’ $X_{4262}$ Tc-poor blends (with a clear central pseudo-emission, see HD 73341 and Hen 3 on Fig. \[Fig:figTc62\]) and the ‘ambiguous’ $X_{4262}$ blends (where this pseudo-emission is absent, as in Hen 7); two typical transition cases are plotted on Fig. \[Fig:figTc62\] (Hen 187 and Hen 31). These ‘ambiguous’ spectra were taken during different observing runs; the shape of the $X_{4262}$ blend is independent of the resolution and of the S/N ratio of the spectra.

Do these ‘ambiguous’ spectra correspond to stars with a weak technetium line, intermediate between the clear Tc-poor and Tc-rich cases? In fact these ‘ambiguous’ spectra are clearly different from the spectum of the weakly Tc-rich star Hen 140, for [*their $X_{4238}$ feature is identical to the $X_{4238}$ feature of the unambiguous Tc-poor stars*]{} (Fig. \[Fig:figTc38\]), which clearly indicates that technetium is absent in these stars.

It may therefore be concluded that our method of gaussian fit to the $X_{4238}$ and $X_{4262}$ features is able to properly separate technetium-rich from technetium-poor S stars.

What then causes the variety of $X_{4262}$ features observed in Fig. \[Fig:figTc62\] for technetium-poor S stars? The spectral sequence going from Hen 3 to Hen 7 on Fig. \[Fig:figTc62\] is not a temperature sequence. The temperature of the stars of our sample have been derived from the $V-K$ color index using the Ridgway et al. ([@Ridgway]) calibration, the $K$ magnitudes from Catchpole et al. ([@Catchpole79]) and our Geneva photometry. Although the bulk of technetium-rich S stars are clearly cooler than technetium-poor S stars (see also Van Eck et al. [@Van; @Eck98]), there is no sign whatsoever of a possible correlation between the shape of the $X_{4262}$ blend of technetium-poor stars and their temperature. MOOG (Sneden [@Sneden]) synthetic spectra (for stars with T$_{eff}\sim 3400-3800$K as derived from their $V-K$ index) indicate that neither gravity nor metallicity can significantly modify the $X_{4262}$ blend.

A closer inspection of the spectral sequence of Fig. \[Fig:figTc62\] (from HD 73341 to Hen 7) reveals that several lines become stronger as the central pseudo-emission of the $X_{4262}$ blend weakens. The major contributors to these features, identified with the help of synthetic spectra, are indicated on the top of Fig. \[Fig:figTc62\]. It is noteworthy that [*all these elements are s-process elements*]{}. The sequence of spectra (drawn with a thin line) in Fig. \[Fig:figTc62\] is thus, from top to bottom, a sequence of increasing s-process line strengths (s-process lines being weak, as expected, in the M star HD 73341). The line which progressively blends the $X_{4262}$ feature of technetium-poor stars is thus probably an s-process line as well. Since it cannot be technetium (see above), a good candidate is the 4262.228Å  line of , or perhaps the wing of the line at 4262.087Å.

It is not surprising to find a wide range of s-process enhancements among technetium-poor S stars, since these stars have accreted their s-process-enriched matter from a companion star. Hence the level of chemical peculiarities is not linked to the evolutionary status of the star, but rather depends upon the amount of s-process accreted matter (see Jorissen et al. [@Jorissen] for a detailed discussion).

These s-process lines are more difficult to see in the technetium-rich S stars, probably because in these cooler and more luminous stars, lines are broader (because of a larger macroturbulence) and the molecular blanketing is more severe.

Misclassified and SC stars {#Sect:SC}
--------------------------

The method outlined in Sect. \[Sect:Analysis\] cannot be applied to four stars of our sample which exhibit peculiar spectra (Hen 22, 135, 154 and 198). In order to check the assignment of the Henize stars to spectral type S, low-resolution spectra ($\Delta \lambda \sim 0.3$ nm, $4400$Å$<\lambda< 8200$Å) have been obtained for all stars from Hen 3 to Hen 165 at ESO on the 1.52m telescope equipped with the Boller & Chivens spectrograph and grating \#23 (Van Eck et al. [@Van; @Eck]).

Two misclassified stars have been uncovered: Hen 22 and Hen 154 show no sign of ZrO bands whatsoever in their spectra. Besides, Hen 22 is classified as ‘S:’ by Henize. Both stars cannot be dwarfs because their NaD and MgH $\lambda$4780Å  features are too weak. Their prominent $\lambda$ 4455Å  line and their weak CN $\lambda$7895Å  band point towards them being giant stars rather than supergiants. Type Ia supergiants can certainly be ruled out because their absolute magnitudes ($M_{\rm v}=-7.8$ and -7.5 for G8Ia and K3-5I respectively, Landolt-Börnstein [@Landolt]) would result in much too large heights above the galactic plane (7.5 kpc for Hen 154 and 13.2 kpc for Hen 22). Hen 154 is probably a late G giant ($\sim$G8), and Hen 22 a mid-K giant ($\sim$K3-5). These assignments are compatible with the Geneva photometry available for these two stars.

As far as Hen 135 and Hen 198 are concerned, Fig. \[Fig:figTc62except\] shows that the spectra of these stars are very different from those of other S stars of the Henize sample. Many spectral features adopted as wavelength standards, as well as the technetium blend, are difficult or even impossible to identify in the spectra of Hen 135 and Hen 198. In fact, we show below that these two stars are the only two SC stars in the subsample of Henize S stars observed with the CAT[^3]. Hen 135 ($V\sim 7$) and Hen 198 ($V\sim 7-10$) were indeed the only very red stars ($B-V>2$) which were bright enough to allow spectra to be taken in the violet.

SC stars are known to have very peculiar spectra. Their spectrum is filled with strong atomic lines and almost no molecular bands in the optical, a consequence of their C/O ratio being very close to unity (Scalo [@Scalo]). Catchpole & Feast ([@Catchpole71]) define SC stars from the following three criteria: (i) extremely strong Na D lines, (ii) drop in the continuum intensity shortward of 4500Å, and (iii) bands of ZrO and CN simultaneously present (though quite weak), as well as general resemblance of the spectrum (i.e. regarding ‘the absolute and relative strength of metal lines’) with that of UY Cen.

Hen 135 (=UY Cen) is thus the prototype SC star. Our two spectra of that star (taken in March 1993 and January 1997, see Fig. \[Fig:figTc62except\]) are quite different; in particular the shape of the $X_{4262}$ feature has changed noticeably. Therefore it is hazardous to infer the technetium content of UY Cen from these data alone without the help of appropriate model atmospheres and synthetic spectra, which is beyond the scope of this paper.

Hen 198 (=RZ Sgr) has an Se-type spectrum; Stephenson ([@Stephenson]) quotes the HD catalogue noting that ‘the spectrum is similar to class N, but does not belong to that class’. It is probably associated with a reflection nebula (Whitelock [@Whitelock94]). RZ Sgr is a large-amplitude ($\sim2.5$ mag) SRb-type variable ($P=203.6$ d). Its H$\alpha$ emission, as well as the TiO and ZrO band strengths, are variable. Catchpole & Feast ([@Catchpole76]) also note that the Zr:Ti ratio of RZ Sgr is unusually high for an S star, and rather close to the one of N-type carbon stars.

Although RZ Sgr has not been classified as an SC star, it shares many common features with that family. Indeed, it reasonably meets the three criteria mentioned above for SC stars:\
(i) Reid and Mould ([@Reid]) measured the strength of the Na D lines for several S, SC and C stars, including RZ Sgr. A spectrophotometric index of 1.07 is found for RZ Sgr, much larger than typical values for S stars (0.22 for and 0.28 for NQ Pup), but comparable to values obtained for SC stars (0.55 for LMC 441, 1.73 for R CMi, 2.57 for VX Aql). Thus RZ Sgr has abnormally strong Na D lines with respect to other S stars.

\(ii) The ultraviolet flux deficiency of SC stars is clearly apparent from photometric data in the Geneva system. Indeed, the mean wavelengths of the $B$ and $V$ filters are $\lambda_0(B)$ = 4227Å  and $\lambda_0(V)$ = 5488Å  (Rufener & Nicolet [@Rufener]); therefore the $B-V$ index is highly sensitive to the ultraviolet flux deficiency of SC stars occurring for $\lambda < 4500$Å. SC stars have $B-V>2$, whereas the bulk of S stars have $B-V<2$. In that respect again, RZ Sgr ($2.0<B-V<3.0$) is typical of SC stars.

\(iii) ZrO is present (although weak) in RZ Sgr; we found no information about the possible presence of CN bands. Infrared CO bands are stronger in RZ Sgr than in many other S and SC stars (Whitelock et al. [@Whitelock85]), probably locking a great quantity of carbon.

The IRAS colours of RZ Sgr also share many similarities with SC stars: it is located in a region of the ($K-[12]$,$[25]-[60]$) color-color diagram (‘region E’ as defined by Jorissen & Knapp [@JorissenKnapp]) containing mainly SC stars with large 60$\mu$m excess and often resolved shells (see also Young et al. [@Young]).

All these arguments therefore indicate that RZ Sgr is closely related to the SC family. As pointed out for UY Cen, the 4262Å  and 4238Å  lines of are very difficult to analyse in SC stars. An assignment of these two stars to either the Tc-rich or Tc-poor group has therefore not been attempted here. Abia & Wallerstein ([@Abia]) nevertheless suggest that SC stars are Tc-rich, based on a quantitative analysis.

Discussion {#Sect:RESULTS}
==========

The technetium dichotomy
------------------------

The lower part of Fig. \[Fig:Tc\] shows the frequency histogram of the wavelength of the $X_{4262}$ spectral feature for stars of Tables \[Tab1\] and  \[Tab2\]. The stars of our sample clearly segregate in two groups. The average wavelength of the bluer group is 4262.093Å; this group thus corresponds to Tc-poor S stars. The average wavelength of the redder group is 4262.235Å, thus revealing the contribution of the 4262.270Å  line to the - blend. The standard deviation on the $X_{4262}$ wavelengths is 0.012Å  for Tc-poor stars and 0.017Å  for Tc-rich stars. These values are in good agreement with the estimated errors on the $X_{4262}$ wavelength (0.013Å  for Tc-poor and 0.020Å  for Tc-rich stars, Sect. \[Sect:Analysis\]).

The two groups are clearly separated by a 0.08Å  gap, [*with no intermediate cases*]{}. Therefore, in order to distinguish Tc-poor from Tc-rich stars (on our spectra of resolution in the range 30000-60000), a delimiting wavelength of 4262.16Å  may be safely adopted.

A similar conclusion holds for the $X_{4238}$ feature (Fig. \[Fig:Tc3862\]), where a boundary wavelength of 4238.29Å  unambiguously separates the two kinds of S stars. Fig. \[Fig:Tc3862\] further shows that the diagnostics provided by the $X_{4238}$ and $X_{4262}$ features are consistent with each other.

For comparison purpose the frequency histogram of the wavelength of the $X_{4262}$ spectral feature as obtained by Smith & Lambert ([@Smith]) is plotted in the upper part of Fig. \[Fig:Tc\], for their sample of MS and S stars (their Table 2). The segregation into Tc-poor and Tc-rich S stars (with 4 stars falling on their boundary wavelength at 4262.14Å) is not as clean as with our higher resolution spectra.

The small number of ‘transition stars’ in our sample (i.e. stars with weak Tc lines, the only case being Hen 140) is noteworthy. This result may provide constraints on the evolution with time of the technetium abundance along the TPAGB (Smith & Lambert [@Smith]; Busso et al. [@Busso]) and clearly deserves further studies. For example, it would be of interest to investigate whether the small number of S stars with weak technetium lines found in our sample implies that the very first objects to dredge-up heavy elements on the TPAGB are not S stars but rather M stars. Indeed, the Tc detection threshold might not coincide with the ZrO detection threshold, but be slightly lower (i.e. [*stars would appear Tc-rich before being ZrO-rich*]{}); this would explain the puzzling Tc-rich M stars discovered by Little-Marenin & Little ([@Little-Marenin]) and Little et al. ([@Little]). It is moreover necessary to disentangle abundance effects from atmospheric effects on the technetium line strength.

Although a more detailed study is deferred to a forthcoming paper, it may already be mentioned at this point that the Tc/no-Tc dichotomy reported in this paper is not due to technetium being entirely ionized in the warmer S stars. Indeed, the /( + ) ratio is still $\sim40\%$ in the warmest S stars (T$_{eff}$=3800 K) while it amounts to $\sim70\%$ at T$_{eff}$=3500 K and to $\sim95\%$ at T$_{eff}$=3000 K, according to the Saha ionization equilibrium formula (with representative electron densities taken from model atmospheres).

M stars with 60$\mu$m excess and symbiotic stars
------------------------------------------------

None of the four M stars with 60$\mu$m excess taken from the sample of Zijlstra et al. ([@Zijlstra]) show technetium. This observation clearly indicates that these stars, which are surrounded by cool dust dating back to a former episode of strong mass loss, do not currently experience heavy elements synthesis followed by third dredge-ups. The same conclusion holds true for the two observed symbiotic stars (SY Mus and RW Hya).

Conclusion
==========

High-resolution spectra have been obtained and analysed to infer the technetium content of 76 S, 8 M and 2 symbiotic stars. The presence or absence of technetium was deduced from the shape of two blends involving technetium at 4238Å  and 4262Å  (more precisely: from the wavelength of their minimum). However this method does not apply to SC stars. Two misclassified S stars (Hen 22 and Hen 154) have emerged. The technetium (non-)detection at 4238Å  is consistent with the result at 4262Å. Only one ‘transition’ case (Hen 140 = HD 120179, a star where only weak lines of technetium are detectable) is found in our sample.

A resolution in excess of $30\,000$ is definitely required to provide unambiguous conclusions regarding presence or absence of technetium. For example, at 4262Å, an s-process line (possibly ) is suspected to sometimes mimick a weak technetium line (although the 4238Å  feature clearly shows that technetium is absent). The shape of the $\lambda$4262Å  feature varies from one Tc-poor star to another, depending on the s-process overabundance level, which is in turn a function of the amount of accreted matter by these binary S stars.

Among the 70 analysed Henize S stars, 41 turn out to be technetium-poor and 29 technetium-rich. That fraction may not be used, however, to infer the relative frequencies of intrinsic and extrinsic S stars, since the subsample of Henize S stars observed with the CAT is biased towards the brightest and bluest stars. The frequency of extrinsic/intrinsic S stars will be derived from the whole data set in a forthcoming paper.

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[^1]: Research Associate, National Fund for Scientific Research (FNRS), Belgium

[^2]: Based on observations carried out at the European Southern Observatory (ESO, La Silla, Chile)

[^3]: Although the Henize sample contains several SC stars, they are usually too red, hence too faint at 4250Å, to be observed with the CAT.
