C Martayan
email: cmartaya@eso.org

T Rivinius

D Baade

A.-M Hubert

J Zorec

Review about populations of Be stars: stellar evolution of extreme stars

Keywords: stars: emission-line, Be, stars: formation, stars: evolution, stars: rotation, stars: abundances, stars: magnetic fields, stars: oscillations (including pulsations), binaries: general, surveys, gamma rays: bursts

Among the emission-line stars, the classical Be stars known for their extreme properties are remarkable. The Be stars are B-type main sequence stars that have displayed at least once in their life emission lines in their spectrum. Beyond this phenomenological approach some progresses were made on the understanding of this class of stars. With high-technology techniques (interferometry, adaptive optics, multi-objects spectroscopy, spectropolarimetry, highresolution photometry, etc) from different instruments and space mission such as the VLTI, CHARA, FLAMES, ESPADONS-NARVAL, COROT, MOST, SPITZER, etc, some discoveries were performed allowing to constrain the modeling of the Be stars stellar evolution but also their circumstellar decretion disks. In particular, the confrontation between theory and observations about the effects of the stellar formation and evolution on the main sequence, the metallicity, the magnetic fields, the stellar pulsations, the rotational velocity, and the binarity (including the X-rays binaries) on the Be phenomenon appearance is discussed. The disks observations and the efforts made on their modeling is mentioned. As the life of a star does not finish at the end of the main sequence, we also mention their stellar evolution post main sequence including the gamma-ray bursts. Finally, the different new results and remaining questions about the main physical properties of the Be stars are summarized and possible ways of investigations proposed. The recent and future facilities (XSHOOTER, ALMA, E-ELT, TMT, GMT, JWST, GAIA, etc) and their instruments that may help to improve the knowledge of Be stars are also briefly introduced.

Introduction

The emission-line stars are spread over the entire HR digram. They concern young, not evolved stars or at the opposite evolved stars and of course Main Sequence stars. They are or not massive stars. For instance, among others, one can found TTauri stars, UV Ceti stars, flare stars, Mira stars, HBe/Ae stars, B[e] stars, Be/Oe stars, Of stars, Supergiant stars, LBV, WR stars. This document concerns mainly OeBeAe stars. The OeBeAe stars have properties similar to OBA type star (see Evans contribution, this volume) and specific properties due to their extreme nature. The first Be star, γ Cas was discovered by the father Secchi in 1867 (Secchi 1967). Collins (1987) wrote the first definition of Be stars: "they are non-supergiant OBAtype stars that have displayed at least once in their spectrum emission lines (Hα)." As example, see Fig. 1. Actually the emission-lines come from the circumstellar decretion 2 C. Martayan, T. Rivinius, D. Baade, A.-M. Hubert, & J. Zorec disk (Struve 1931) formed by episodic matter ejections from the central star. Are the matter ejections related to the rotation? Be stars are very fast rotators, rotating close to the breakup velocity but often an additional mechanism is needed for the star being able to eject the matter. In the next sections, the rotation and other mechanisms (binarity, stellar evolution, magnetic fields, pulsations) are examined as well as the effects of the metallicity, the stellar evolution, etc. Porter & Rivinius (2003) presented different kind of emission-line stars with their known properties and how they differ to the "classical Be stars". Since 2003, with the improvements of the technology (interferometry, MOS, etc) and the efforts made on the modeling and theory, the knowledge of Be stars was improved. The next sections will provide some very summarized informations (since 2004, more than 240 refeered articles were published dealing with Be stars) about these recent developments (see also contribution by Baade, this volume).

-Techniques for detecting a Be star Before studying Be stars, it is necessary to find/recognize them. There are several ways to detect them. The first possibility is to use the photometric techniques. Combining different colour/colour or colour/magnitude diagrams would help to pre-classify the stars and detect potential Be star candidates. Keller et al. (1999) provide examples of CMD with given thresholds above which the Be stars could fall. They also show that Be stars tend to form a redder sequence than normal B stars. Dachs et al. (1988) have shown that the infrared excess is related to the circumstellar disk of Be stars. It is also linked to the Hα equivalent width. More recent studies found similar results in other wavelength domains such as in the infrared with the AKARI survey (Ita et al. 2010) or with SPITZER (Bonanos et al. 2010). However, this kind of study has some limits due to the intrinsic properties of Be stars (variability, change of phases Be-B). McSwain & Gies (2005) Another way to find the emission-line stars, and the Be stars, is to do slitless spectroscopy and combine the spectroscopic diagnostic with the photometry as presented in Martayan et al. (2010a). The slitless spectroscopy gives the advantage to do not be sensitive to the diffuse nebulosity and the spectra are not contaminated by other nebular lines.

Another possibility for finding Be stars is to study the behaviour and the variation of lightcurves. The Be phenomenon is transient, as a consequence the star could appear as a Be star or as a B star when the disk was blown up. Moreover, Be stars show often outbursts. These characteristics help to find this kind of stars using the lightcurves variations as published by Mennickent et al. (2002) and Sabogal et al. (2005).

-The Surveys Be stars with the techniques mentionned above can be found mainly in open clusters and in fields. In the Galaxy, there are different surveys, let us mention, the photometric survey by McSwain & Gies (2005) and the slitless spectroscopic survey by Mathew et al. (2008).

In other galaxies such as the Magellanic Clouds but also in the Milky Way, there are the following surveys among others: a slitless spectroscopic survey of the SMC by Meyssonnier & Azzopardi (1993), another in the SMC/LMC/MW by Martayan et al. (2010a), Martayan et al. (2008) (see Fig. 2). (2008). The stars appear as spectra.

And recently, with the VLT-FLAMES, some spectroscopic surveys were and are performed in different galaxies (Evans et al. 2005, Martayan et al. 2006a, Martayan et al. 2007a). With the polarimetry, some Be stars were found in the SMC/LMC (Wisniewski et al. 2007, Wisniewski & Bjorkman 2006), and with FORS2 in IC1613 (Bresolin et al. 2007), which correspond to the farthest Be stars detected.

2. Mid/long term monitoring of Be stars, V/R variations -Monitoring of Be stars Why long-term monitorings are needed? Due to the transience of the Be phenomenon for having a chance to detect a Be star, a mid or long term monitoring of B/Be stars is needed. Indeed, for the follow-up of the variability of the star related to outbursts and/or pulsations, or binarity, or modifications of the CS disk, etc, a monitoring is needed.

Surveys such as the photometric ones of MACHO or OGLE are useful in that sense but also the spectroscopic surveys. Here the amateur astronomers have a role to play. With the new efficient, good quality, low cost spectrograph that can be mounted in amateur telescopes, see contribution by Blanchard et al. (this volume), some useful good quality observations can be performed.

Moreover, without long-term monitoring (more than 10 years of spectroscopy) the results concerning the Be star µ Cen and the predictions of the outbursts dates could not have been obtained. For all details about this pioneer study, see Rivinius et al. (1998a), Rivinius et al. (1998b).

This kind of study could thanks to the databases be performed more easily now, let us cite the Be Star Spectra database http://basebe.obspm.fr/basebe/ (Neiner et al. 2007, contribution by De Batz et al., this volume).

-The V/R variation of the emission lines A long-term spectroscopic monitoring of Be stars can show in certain cases variations of the V and R peaks of the emission-lines. This variation is now understood and is related to global one-armed oscillations in the circumstellar disk. The case of the Be star ζ Tau (see Stefl et al. 2009) is representative of that phenomenon with a V/R cycle of about 1400 days without correlation with the period of 133 days of the companion. Fig. 3 from Stefl et al. (2009) illustrates this phenomenon. The (polarimetric, spectroscopic, interferometric) observations are properly reproduced by a modeling of a one-armed spiral viscous disk as described in Carciofi et al. (2009), see also contributions by Carciofi and Stee, this volume.

Stellar winds

The diagnostic lines are mostly present in the UV and optical but also several of them are in the infrared. Let us mention the lines of HeI, HeII, Si IV, C IV, N V, Brγ, etc. For more details, see Martins et al. (2010), Mokiem et al. (2006), and Henrichs et al., this volume. The structure of the wind differs among B stars. For the normal B stars, the wind is spherical, while for Be stars, we can expect a radiative bi-polar wind mainly due to the Von Zeipel effect. The poles in case of fast rotation becomes brighter than the equator, it explains this asymetry. In case of Be stars, one can also expect an additional mechanical wind at the equator that creates the CS disk. In such case the needed value of the Ω/Ω c ratio (angular velocity to critical angular velocity ratio) for ejecting matter has to be determined (see Sect. 4).

Fortunately, recent interferometric facilities (such as the VLTI-AMBER), allowed to measure the shape of the CS environment of Be stars and both winds were found. It is for example the case of the Be stars α Arae, Achernar (Meilland et al. 2007, Kervella et al. 2006).

Rotational velocities

Be stars are known to be very fast rotator but how fast are they rotating? Fig. 4 shows effect of the rotation on the shape of the stars for different Ω/Ω c in the Roche model, rigid rotation approach. Next sections provide some clues about this point. -From the interferometry The interferometric observations show that all observed Be stars are found flattened revealing high Ω/Ω c ratios. The flattening of Achernar was found equal to 1.56 ±0.05 by Domiciano de Souza et al. (2003). However, this value could also be due to remaining small disk indicating that Achernar is not rotating at the breakup velocity but close to it.

-Line saturation effect at high Ω/Ω c Townsend et al. (2004) showed that the HeI 4471 Å line has a saturation effect at high Ω/Ω c , typically above 80%. This implies that the Vsini measurements based on the use of 1 He line are limited in case of fast rotating stars, typically Be stars and possibly Bn stars. Using several He lines as well as metallic lines, the saturation effect occurs again but at higher Ω/Ω c (∼90-95%), see Frémat et al. (2005). This explains the discrepancy between the measurements of Be stars Vsini between Hunter et al. (2008) and Martayan et al. (2007a). Fig. 5 shows the measurements differences from Hunter et al. (2008) and the saturation curves at Ω/Ω c =70 and 80 % from McSwain et al. (2008). It clearly shows that the Be star Vsini measurements done by Hunter et al. (2008) are saturated and the measurements by Martayan et al. (2007a) saturate at ∼90%. -Fast rotation effects correction Accordingly to the previous section, in case of fast rotators, one has to take into account the fast rotation effects (flattening, gravitational darkening). These effects introduce underestimates of the Vsini, and it implies too that the spectral type/luminosity classes are modified from the late to earlier spectral type and from evolved to less evolved luminosity classes. The codes BRUCE by Townsend et al. (2004) and FASTROT by Frémat et al. (2005) are able to properly correct the fundamental parameters of the fast rotating stars.

-The veiling effect Another important effect in case of Be stars is the veiling effect due to the circumstellar disk that affects the continuum level and thus the lines depth. Not taking this effect into account will imply bad measurements of the fundamental parameters but also of the chemical abundances. For taking this effect into account, Ballereau et al. (1995) proposed a method based on the He4471 emission-line EW measurements. Other possibility is to prefer the bluest lines when spectrum fitting is performed but also to determine the disk influence by fitting the Spectral Energy Distribution. The new spectrocopic facility, the VLT-XSHOOTER that covers the near UV to the K-band simultaneously should help in that approach. Alternatively, there is the method consisting to fit the Balmer discontinuity as described by Zorec et al. (2009) -Metallicity effects Keller (2004), Martayan et al. (2006a), Martayan et al. (2007a), andHunter et al. (2008) found that SMC OB stars rotate faster than LMC OB stars, which rotate faster than their Galactic counterparts. In the SMC, the mass-loss of OB stars by radiatively driven winds was found lower than in the Galaxy by Bouret et al. (2003) and Vink (2007). Consequently, the stars should loose less angular momentum and could rotate faster (Maeder & Meynet 2001). The comparisons between SMC Be stars Vsini fairly agree with the theoretical models by Ekström et al. (2008). The SMC Be stars are found to be rotating very close to the critical velocity.

-

ZAMS rotational velocities

The ZAMS rotational velocities of SMC, LMC, and Milky Way (MW) intermediatemass Be stars were determined by Martayan et al. (2007a). They found a metallicity effect on the ZAMS rotational velocities. At lower metallicity Be stars rotate faster since their birth. This could be related to an opacity effect, for an identical Ω/Ω c , at lower metallicitty, the radii are smaller thus the stars can rotate faster. The comparison of the SMC ZAMS rotational velocities of intermediate-mass and massive Be stars with the theoretical tracks by Ekström et al. (2008) are also in a fair agreement.

-Number of Be stars vs. the metallicity and redshift As a consequence of faster rotational velocities at lower metallicity, one can expect to find more Be-type stars in low-metallicity environments. Using the ESO-WFI in its slitless mode (Baade et al. 1999) Another consequence of this increase of the Be stars number with decreasing metallicity is when a survey of OBA-type populations is performed, one should find an increased number/ratio of OeBeAe stars in the sample. This was exactly the case of the survey by Bresolin et al. (2007) who found 6 Be stars in the 6 MS B stars observed in the low metallicity galaxy IC1613. Actually at lower metallicity than the SMC, and probably at higher redshift, the stars could rotate again faster, then one can expect to have more Be stars but also that the Be phenomenon is extended to other categories of stars.

-Chemical abundances The measurement of chemical abundances is very important for testing the stellar evolution models with fast rotation, including the rotational mixing due to the rotation. Moreover, it is also important for distinguishing the stars following a chemical homogeneous evolution than the usual evolution (Maeder et al. 1987, Yoon et al. 2006). It is also of interest to know them for the pulsating stars of different metallicity environments. Lennon et al. (2005) found 2 Be stars without N enhancement, while the theoretical models using the rotational mixing expect a N enhancement and a C depletion. Dunstall et al. (this volume) and Peters (this volume) also found no N enrichment in Be stars. Does it mean that the rotational mixing theory is not able to explain the measurements? Not necessarily, Porter (1999) found that due to the temperature gradient in Be stars, the ions are displaced from the poles to the equator. This could occur to the N, implying that the corresponding lines become weaker and the measurements biased.

Stellar pulsations

Be stars lie in β Cep and SPB regions, thus p and g modes are expected to be found. The κ mechanism is at the origin of the observed pulsations due to the iron bump. A beating of non radial pulsations combined to the fast rotation of Be stars could be at the origin of the matter ejection. According to Townsend et al. (2004) an Ω/Ω c =95% is needed for launching the matter in orbit, while Cranmer (2009) shows that depending on the Ω/Ω c value the ejected matter will be part of the wind or will form a disk. Rivinius et al. (1998a) have shown that the combination fast rotation + beating of non radial pulsations is able to reproduce and predict the outbursts in the Be star µ Cen. More recently, with the new photometric space missions MOST, COROT, KEPLER, a large number of pulsational frequencies (g and p modes) were found in several Be stars. Huat et al. (2009), in the Be star HD49330, have found that the pulsational frequencies but also their amplitude change between the quiescence phase, the pre burst phase, and during the burst phase. Is it related to the subsurface convection zones in OB stars? For more details, see Cantiello contribution (this volume).

Spectropolarimetry: magnetic fields and disks

The spectropolarimetry is a useful tool for finding the Be stars due to their disks but also for detecting the magnetic fields. Recent spectropolarimetric facilities such as ESPADONS, NARVAL, allow the search of magnetic fields in Be stars but up to now, there is only 1 Be star in which a magnetic field was found: ω Ori (Neiner et al. 2003). The presence of magnetic field (Ω transport from the core to the surface or magnetic reconnection) combined to the fast rotation could also be an additional mechanism for ejecting matter and create the Be star CS disk.

Circumstellar disks

Huge progresses were performed thanks to the interferometry (VLTI, CHARA, NPOI) and the multi-wavelengths studies. The main result is that the disk is rotating in a Keplerian way around its central star, see for example Meilland et al. (2007). In addition, the thermal structure of the CS disk is also better known (see Tycner contribution, this volume).

Stellar formation and evolution, Gamma Ray Bursts

What kind of objects is at the origin of Be stars and their fast rotation? Could they be related to the Herbig Ae/Be objects? Are the fast B rotator Bn stars at the Be star origin, while their number is roughly equal to the number of Be stars in the MW? What is the IMF of the Be stars? Does it differ from the normal B stars? It seems that the open cluster density has no effect on the appearance of Be stars (McSwain & Gies 2005, Martayan et al. 2010a).

Taking into account the fast rotation effects, Zorec et al. (2005) derived the evolutionary status of Be stars in the MW (see Fig. 8). Their location/life in the MS strongly depends on their mass and on the evolution of the Ω/Ω c ratio. In the LMC, the diagram of Be stars appearance is similar than in the MW but in the SMC, at the opposite of the MW, massive Be stars and Oe stars appear or are still existing in the second part of the MS (Martayan et al. 2007a). These differences are also consistent with the Ω/Ω c evolution at different metallicities.

However, the life of a Be star is not finished at the end of the MS. It seems that the more massive SMC Be and Oe stars could play a role in the explanation of the type 2 GRBs (also called Long Gamma Ray Bursts). Martayan et al. (2010b) were able to reproduce the number and ratio of LGRBs using the populations of SMC massive Be and Oe stars accross their stellar evolution. The type 3 GRBs (without SN counterpart) could be explained in a binary scenario involving Be-binary stars (Tutukov & Fedorova 2007).

Binarity

While 60 to 70% of O stars are known to be binaries (see Sana et al. 2009 and his contribution in this volume), the question remains opened for the B and Be stars. The last study by Oudmaijer et al. (2010) using the VLT-NACO AO facility found that: 29 to 35 ±8% of the nearby B stars are binaries and that 30 to 33 ±8% of the nearby Be stars are binaries. They scanned the separations larger than 20 AUs with a flux contrast ratio of 10 corresponding to M companions. Scanning smaller scales needs the interferometry. About 40 % of all Be stars observed with the VLTI and CHARA were found to have a companion. However, certain type of Be stars are known to be binaries such as the Be-X rays (see for example Coe et al. 2008). The γ Cas-like stars have possible companions and have magnetic field activity in their CS environment (Smith et al. 2006).

About the Be-binary disruptions scenario when the companion has exploded in SN, one should find that Be stars are runaway stars or see SN remnants. In both cases, this is not observed for the large majority of Be stars. The chemical abundances could help to find a potential previous interaction with a companion. In that case, one can expect abundances anomalies.

Conclusion

To better understand the properties of these stars, it is necessary to do multiwavelength studies and multi-techniques studies with the use for example of: the VLT-XSHOOTER, IR and X-rays satellites, MOS such as the VLT-FLAMES, the interferometry, the AO facilities, and if possible simultaneous observations.

In the future, it will be possible to do similar studies than today but in very far galaxies, to scan other metallicity ranges, etc, thanks to the ELTs and their instruments but also to ALMA, SKA, etc.

As shown, the Be phenomenon could be not restricted to the late O, B, and early A stars and could also not be restricted to the MS, specially in lower metallicity environments. Some of the questions about them and their CS environment are now solved but some others are still opened and the definition of Be stars should be revised.

Finally, one can propose this new definition of the Be phenomenon: This is a star with innate or acquired very fast rotation, which combined to other mechanism such as non-radial pulsations beating leads to episodic matter ejections creating a CS decretion disk or envelope.

This definition implies that the stars are not restricted to the MS, not restricted to B-type, and from the CS material the emission lines come.

Figure 1 .

 1 Figure 1. Examples of Hα emission-line in different LMC Be stars. The shape depends among other parameter of the inclination angle.

  in their photometric survey of Galactic open Review about populations of Be stars: stellar evolution of extreme stars 3 clusters found 63 % of the known Be stars due to the transience of the Be phenomenon. Thus with the photometry the coupling of indices is necessary. It is also necessary to disentangle the local reddening and the infrared excess due to the circumstellar disk of Be stars.

Figure 2 .

 2 Figure 2. Example of slitless spectroscopy in the NGC6611 open cluster by Martayan et al.(2008). The stars appear as spectra.

Figure 3 .

 3 Figure 3. Example of V/R variations in the Be star ζ Tau. Figure adapted from Stefl et al. (2009).

Figure 4 .

 4 Figure 4. Example of star shape depending on the Ω/Ωc ratio in the Roche model and rigid rotation approach. Due to the Von Zeipel effect there is a temperature gradient from the poles to the equator. Courtesy by Y. Frémat, see also Frémat et al. (2005).

Figure 5 .

 5 Figure 5. Vsini measuremnts of B and Be stars by Hunter et al. (2008) compared to those of Martayan et al. 2007a. The dotted curve corresponds to Ω/Ωc=70% and the solid curve to Ω/Ωc=80% (line saturation effect with 1 He line measurement). The curves come from McSwain et al. (2008).

  , Martayan et al. (2010a) found 3 to 5 times more Be stars in SMC open clusters than in the Galactic ones (McSwain & Gies 2005), see Fig.6. This quantified the results byMaeder et al. (1999) andWisniewski et al. (2008).

Figure 6 .

 6 Figure 6. Ratio of Be to B stars by spectral categories in the SMC and MW. Figure adapted from Martayan et al. (2010a).Red dashed bars are for the SMC, white bars for the MW.

Figure 7 .

 7 Figure 7. Evolution of pulsational frequencies in the Be star HD49330 from COROT data during the different pre, during, and after an outburst phases. The figure comes from Huat et al. (2009).

Figure 8 .

 8 Figure 8. Evolutionary diagram in the MS of the Be stars in the MW by Zorec et al. (2005).

Discussion

T. Rivinius: Comment on the remark by A. Miroshnichenko about the Be binary fraction. The binary fraction could well be 50 %. What is really important is that in studies where B and Be stars are investigated, after having made sure that the same selection effects apply to both, the binary fraction is found the same for B and Be stars. At least for binary separations wide enough not to destroy the disk by interaction. R. Prinja: Are the fast, high-ionization polar winds of Be star clumps (and perhaps structures) in a manner akin to massive OB stars? Fundamental differences in wind characteristics of these classes of hot stars may constrain the pivotal role of for example Fe-bump driven convection and the role of disk-wind driving in Be stars.

C. Martayan:

The clumping and the wind structure of Be stars are not yet known. Maybe the current interferometric programs on Be stars could bring some preliminary informations.

A. Maeder: During the cycle Be-B, are luminosity variations observed? C. Martayan: Yes, during the life cycle of the CS disk of Be stars, some luminosity variations were reported and they follow a sequence that is shown by De Wit et al. (2006).


Christophe Martayan

Ronny Blomme

Jean-Baptiste Le Bouquin

Antoine Merand

Guillaume Montagnier

Fernando Selman

Julien Girard

Andrew Fox

Dietrich Baade

Yves Fremat

J.-B Le Bouquin

A Merand1

A Lobel

F Martins

F Patru

Th Rivinius

H Sana

S Štefl

J Zorec

T Semaan

High-angular resolution observations of the Pistol Star

Keywords: stars: binaries, stars: circumstellar matter, stars: early-type, stars: mass loss

à la diffusion de documents scientifiques de niveau recherche, publiés ou non, émanant des établissements d'enseignement et de recherche français ou étrangers, des laboratoires publics ou privés.

Introducing the Pistol Star

At the time of its formation ∼2 million years ago, the Pistol Star may have been one of the most massive stars in the Milky Way (Figer et al. 1998). Like most other LBV's, it is surrounded by nebulosities, which, therefore, may actually have been ejected by the Pistol Star (cf., e.g., η Car). The shape of the main gas cloud has earned the star its nick name. Owing to its own ejecta but mostly due to its far-away location (∼ 8 kpc) in the region of the Galactic Center, observations of the Pistol Star suffer from very high extinction (20-30 magnitudes in the optical).

Among massive stars, the fraction of binaries may approach 100% at the time of formation (cf. Mason et al. 2009, Sana et al. 2009 ). In fact, binarity may be one means to enable the formation of massive stars, which with single-star-plus-disk symmetry encounters challenges already at the 10-M ⊙ level while observations of LBVs imply original masses that are larger by more than an order of magnitude. One of the most prominent massive stars reported to be a binary is η Car (Damineli, Conti, & Lopez 1997).

Observations

In order to more closely examine the nature of the Pistol Star, near-IR observations were undertaken with the VLT AO camera NACO and the VLTI spectrograph AMBER. The angular resolutions of NACO, ranging from natural seeing to near the defraction limit at ∼15 mas, and AMBER, 20-3 mas, are perfectly complementary to one other. At a distance of 8 kpc, this permits spatial scales to be studied down to ∼25 AU. The new VLT spectrograph X-Shooter provided spectra with R ∼ 5000 -11000 over the range 0.3-2.5 µm. They show a vast number of nebular lines; the stellar continuum begins to emerge at ∼ 800 nm. At the time of the Symposium, some additional observations were pending, also of other massive stars.

First preliminary results

The NACO images resolve most of the stars within 30 arcsec of the Pistol Star into several point sources. A quick preliminary inspection of the AMBER interferometry suggests the presence of a further point source at a separation of about 10 mas or ∼ 80 AU. For comparison: If the 5.5-year period in η Car is orbital, its companion would be at one-fifth or less of this distance. In the NACO images, the inner nebulosity looks less structured than in the case of η Car. An effort will be made to continue the observations in order to establish whether the Pistol Star has a physical companion with effects on formation and evolution of the system.

  † based on ESO runs 085.D-0182(A), 085.D-0625(A, C) , R. Blomme, J.-B. Le Bouquin et al.

Figure 1 .

 1 Figure1. NACO image (2.3x2.2 arcsec) of the Pistol Star. Note the appearance of several nearby point sources, while the Pistol Star itself is surrounded by extended nebulosity, measuring roughly 700 AU across. The structure in the nebulosity is not due to discontinuities in the color coding but believed to be real, possibly evidencing multiple shells.

  Fig. 1 depicts the immediate vicinity of the Pistol Star itself as observed with NACO in K band. Apart from several point sources, it reveals an extended nebulosity. At a distance of 8 kpc, its diameter would correspond to more than 6000 AU. Figer et al. 1998 find gas velocities with a range of ∼50-100 km/s but over a larger area. Combining these numbers leads to a crude estimate of the expansion age of 290 years. If the inference from Fig. 1 of multiple shells is correct, they would differ in age by ∼30 years.


Christophe Martayan

Dietrich Baade

Jean Zorec

Yves Fremat

Juan Fabregat

Sylvia Ekstrom

Y Frémat

S Ekström

Massive Oe/Be stars at low metallicity: Candidate progenitors of long GRBs?

Keywords: gamma rays: general, stars: supernovae: general, stars: rotation, Magellanic Clouds

come

Long GRBs and the collapsar model

In the collapsar model for Ib/c supernovae (Woosley 1993), matter with high specific angular momentum is accreted by the already-formed black hole with a short delay. The intermediate disk structure gives rise to the formation of relativistic jets emitting the so-called long gamma-ray bursts (GRBs). The gamma-ray radiation is strongly relativistically beamed with opening angles of order 10 degrees.

In order for a rapidly rotating core to be present, quasi-homogeneous chemical evolution is essential. Stellar evolution models show that in stars with very high initial rotation rates mixing is faster than chemical gradients are built up by nuclear burning (for more details see Yoon & Langer 2005).

The angular momentum still available at this late phase of stellar evolution is the larger, the less mass and, therefore, angular momentum the progenitor star has lost during its lifetime. Since radiatively driven mass loss from massive stars decreases with metallicity, this has led to the notion that low-metallicity star-forming regions should be preferred hosts of GRBs. In fact, observations are beginning to support this picture (e.g., Modjaz et al. 2008). Moreover, the GRB rate per unit mass increases with redshift z. But it is not clear whether this requires any explanation other than the general increase in star formation with redshift (up to z∼2).

Oe/Be stars

In the Milky Way, the near-main sequence stars with the highest rotation velocities are Oe/Be stars. Moreover, many of them rotate at ∼90% of the critical rate. Two arguments 1 2 C. Martayan, D. Baade, J. Zorec, et al. predict a larger fraction of nearly critically rotating massive stars in low-metallicity environments: (i) The mass (and angular momentum) loss rates are lower and (ii), all else being equal, the radii are smaller. Various comparisons of open clusters in the Magellanic Clouds and the Galaxy have concluded (cf. Martayan, Baade & Fabregat 2010) that the abundance of Oe/Be stars increases with decreasing metallicity Z. If one keeps in mind that rapid rotation only is a necessary, but not also sufficient, condition for a rapidly rotating O/B star to become an Oe/Be star, this can pass as a confirmation of the expectation. In any event, because of their emission lines, Oe/Be stars probably are more reliable tracers of a population of rapidly rotating OB stars than broad photospheric lines would be. Therefore, Oe/Be stars more massive than ∼18 M ⊙ (metallicity dependent; Yoon, Langer & Norman 2006) qualify as candidates, also at low Z. Note that, perhaps, Oe/Be stars do not quite reach the theoretical upper mass limit for the progenitors of GRBs (Yoon, Langer & Norman 2006).

GRBs and Oe/Be stars in a synopsis

The strong relativistic beaming makes GRBs observable out to extreme redshifts. But it also implies that much fewer nearby GRBs are detected. Therefore, while it is still relatively straightforward to determine the number of Oe/Be stars per SMC mass and also the space density of such environments as a function of redshift can be estimated, statistics with just a few dozen entries (GRBs) pose serious problems. Nevertheless, using a distribution of beam angles determined from observations, 3-6 long GRBs are predicted at z 0.2 per 11-year period while the observed number is 8 (cf. Martayan et al. 2010). Since the agreement is better than 2 σ, there is no need to fine tune adjustment parameters such as the range in spectral type, from which GRB progenitors are drawn, or the number, type, and star-formation level of the assumed host galaxies. But the excellent agreement could still be coincidental.

Discussion

In addition to the challenges already mentioned, a number of other questions still warrant further examination: • Do Oe/Be stars really preserve their large intial angular momentum throughout their evolution? After all, they regularly eject rotationally supported disks. This behavior is not observed in On/Bn stars, i.e. rapidly rotating O/B stars without emission lines.

• Are there other differences in the evolution of Be and Bn stars?

• Is the SMC a good proxy for GRB-forming environments?

• How incomplete are the observations, especially of GRBs? • Can candidate progenitors other than Oe/Be stars reach the GRB state along different evolutionary paths?


Christophe Martayan

Ronny Blomme

Jean-Baptiste Le Bouquin

Antoine Merand

Guillaume Montagnier

Fernando Selman

Julien Girard

Andrew Fox

Dietrich Baade

Yves Fremat

J.-B Le Bouquin

Y Frémat

A Lobel

F Martins

F Patru

T Rivinius

H Sana

S Stefl

J Zorec

T Semaan

X-shooter, NACO, and AMBER observations of the LBV Pistol Star 1

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Introduction

LBV stars are exceptional transition objects in the stellar evolution of massive stars. Among the emblematic objects of this category, there is η Car. It is now postulated that this kind of star could be a possible aborted supernova. In all cases, they present a special interest for the stellar evolution theory and models because they have strong winds, and asymmetries in their structure, are very luminous at the Eddington limit or above. According to Groh et al. (2009), 2 kinds of "LBV-phases" could exist: the strong variable LBV with S-Dor variability that could be very fast rotators rotating close to the breakup velocity and the group of dormant LBV with less variability like P Cygni. Moreover, the star at the origin of different SN was a LBV like in SN2006jc (Foley et al. 2007) and SN2006gy (Smith et al. 2007). This direct evolution from LBV to SN without the step of WR is not yet understood (Smith et al. 2007) but could lead to the ultra-powerful pair-instability SN.

Observations

The observations are spread over several months and periods: 12/2009, 05-06-09/2010. Early O stars were observed with Xshooter: BRRG56, [ELS2006]N11-026, [ELS2006]N11-029 in the LMC. These stars were selected as they are expected to be among the earliest O stars in the LMC. BRRG56 is classified as an O2 star. They show few absorption lines and several emission lines. In the Galaxy, the brightest O star known [BSP2001]8 (Martins et al. 2008) and one of the most massive claimed with 250 M ⊙ the LBV Pistol Star were observed with Xshooter. In addition, Pistol Star was partially observed with the VLTI-AMBER with the UTs and with the VLT NACO-AO assisted imager.

Brief instruments description

The VLT-Xshooter (D'Odorico et al. 2006): It is the first 2nd generation instrument and is a very efficient echelle spectrograph. One of its interests is the large simultaneous wavelength coverage from the near UV to the K-band with three different arms. The UVB arm allows the observation of the range λλ 290 to 600nm, the VIS arm from λλ 535 to 1050nm, and the NIR arm from λλ 980 to 2500nm. The resolving power of each arm is defined by the slits used, in the UVB arm from 3300 to 9100, in the VIS arm from 5400 to 19000, in the NIR arm from 3500 to 11500.

The VLT-NACO (Lenzen et al. 2003. Rousset et al. 2003) is an adaptive optics assisted imaging in the 1-5 microns. It allows the scan of spatial scales from several arcseconds down to 20-15mas. The observations were performed with the K, Lp, and NB4 filters.

The VLTI-AMBER (Petrov et al. 2007) is a near-infrared, multi-beam interferometric instrument, combining simultaneously up to 3 telescopes. It allows to scan the spatial scales from 20 to 3 mas. The observations of Pistol Star were challenging because this star is the faintest object observed at the VLTI and at the limit of AMBER capabilities and the fringe tracker FINITO cannot be used.

These 3 instruments were used in order to scan different spatial scales for probing the environment of Pistol Star, the nature of Pistol Star itself, and to obtain some parameters/properties of the surrounding nebula.

VLT-Xshooter spectroscopy for the Pistol Star

The observations in the UVB arm (R used=4000, see Fig. 1-left) show a very faint object of magnitude ∼27 in the last 2 red orders (540-600nm). The exposure time was 3000s. this object is at about 3" of the Pistol Star. It could correspond to a faint star in the NACO images at 2.8"-W. In the VIS arm (R used=6700, see Fig. 1-middle), Pistol Star appears in the last 5 red orders (770-1050nm) and disappear in the blue orders due to the huge nebulosity in its vicinity, the blue light is absorbed. The exposure time was the same than in the UVB arm but here the faint UVB object is not seen. The reason is that the UVB arm is more sensitive than the VIS one by more than 1 magnitude. There is also another object that can be seen at the edge of the slit (at the right edge of the orders). One can also note the large number of skylines present. Up to now no bluer spectrum than 1000nm was obtained of Pistol Star. Finally, in the NIR arm (R=11500), the spectrum displayed in Fig. 1-right was corrected of the sky lines, the exposure time is 50s. About 100 NIR exposures were taken. The Pistol Star spectrum corresponds to the central spectrum, there are some emission lines present due to the nebulosity surrounding the central star. A small extract of the NIR spectrum close to the Brγ line is shown in Fig. 2. Some emission-lines around Brγ can be seen, among them there are [FeII], FeII, MgII, NaI, HeI lines. These lines will be compared to the observations by Figer et al. (1999). As a preliminary result, the -EW-of the Brγ line is higher than the measures provided by Figer et al. (1999). NLTE modelling with CMFGEN (Hillier & Miller 1998) and other codes for determining the parameters of the star(s), the wind will be done. The reduction of the data and the identification of the lines are in progress.

Pistol Star and the surrounding stars with the adaptive optics NACO.

The NACO images in Fig. 3 clearly show that the 2 B and C 2MASS stars are 2 groups of stars within few arcseconds. The E star corresponds to the Mira V4644 Sgr and seems to be surrounded by a shell. The D star corresponds to the WR102e. The A star corresponds to Pistol star. The star in the right top of the bottom part of Fig. 3 shows the NACO PSF. One can see several stars close to the central system of Pistol Star at less than 0.4. Moreover it seems that Pistol star is surrounded by different concentric shells of matter. The largest one roughly measures 0.07pc (Pistol Star is at 7.7kpc). This confirms the spectroscopic result by Figer et al. (1999). They found that the spectroscopic lines show the ringlike signature of spherical expansion.

Possible binarity of Pistol Star (VLTI-AMBER)

The figure 4 shows a rough estimates of the UV plane of the VLTI-AMBER observations. There is a departure from single-star morphology. A companion to Pistol Star has been detected with a rough separation of 10-50 mas ( 80 AU). The reduction and the analyse is in progress but this result indicates that like in the Eta Car case (Weigelt et al. 2007), the evolution and the understanding of the Pistol Star structure must go through a binary channel. Long term monitoring of the system must be done in order to confirm the presence of the companion and determine its orbit (check the gravitational link) and the masses of the stars.

Conclusion

Combining multi-wavelength and multi-technique observations, one can better probe the nature of the objects and better understand their properties. In the case of Pistol Star, it was found that it has a spectroscopic variability, that there are small inner shells of matter, and that the star is a binary. These properties combined to others make of Pistol Star a twin of the LBV η Car. More details will be published in a forthcoming article.

Figure 1 :

 1 Figure 1: Snapshots of the 3 XSHOOTER arms of Pistol Star observations. Left: UVB arm (290-600nm), middle: VIS arm (525-1050nm), right: NIR arm (980-2500nm). One can see the spectra of Pistol Star mainly in the NIR arm as well as in the reddest orders of the VIS arm. another very faint object can be seen in the UVB arm.

Figure 2 :

 2 Figure 2: Small extract of the XSHOOTER NIR spectrum of Pistol Star in the region close to the Brγ line.

Figure 3 :

 3 Figure 3: NACO observations of Pistol Star and the stars in its vicinity. The top panel shows the 2MASS image and extracts of the NACO K band image. Bottom panel shows Pistol Star with inner small shells of matter in K-band. The right column shows different stars including a PSF reference star.

Figure 4 :

 4 Figure 4: UV plane of the Pistol Star observations with the triplet of UTs UT234. There is a companion in the NE at 10-50 mas.


Christophe Martayan

Jean Zorec

Dietrich Baade

Yves Fremat

Juan Fabregat

Sylvia Ekstrom

Y Frémat

S Ekström

Evolution of massive Be and Oe stars at low metallicity towards the Long Gamma Ray bursts 1

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Introduction

Let us recall that a Be star is a non-supergiant B-type star which spectrum has shown at least once emission-lines. Actually a Be stars rotate very fast, in the Galaxy at ∼ 85% of the critical rotational velocity. With episodic matter ejection from the central star, a circumstellar decrection disk is formed. This phenomenon is not restricted to B type stars but can also occur in late O and early A stars in the Galaxy. In the first part of this document the metallicity effects on the rotational velocities and on the ratios of Be to B stars are examined.

The second part of that document deals with the long soft gamma ray bursts (here type 2 bursts) and their possible relationship with the massive Be and Oe stars at low metallicity.

2 Ratios of Be stars with respect to the metallicity Maeder, Grebel, & Mermilliod (1999) found that the ratio of Be stars to B stars in open clusters seems to increase with the the metallicity decrease. However, only 1 open cluster in the SMC was used. Wisniewski & Bjorkman (2006) found a similar result but the number of open clusters in the SMC/LMC remained small. The strong variation of the Be/B ratio from an open cluster to another prevents definite conclusions from these studies. It was thus needed to increase the samples (much more SMC open clusters) for improving the statistics, for better constraining the freedom degrees such as the age, the metallicity, the density, etc, and for quantifying the trend found with respect to the metallicity.

• The WFI Hα spectroscopic survey and Be stars ratios Observations with the ESO-WFI in its slitless spectroscopic mode (Baade et al. 1999) were carried out on September 25, 2002 with the aim to map the LMC/SMC central parts for detecting the Hα emission-line stars and the Be stars. The WFI was used with a grism and a Hα filter insuring a bandpass of 700nm, a resolving power of 100 adapted to find the emission in the Hα line of Be stars. The exposure time was 600s. In the SMC 14 images and in the LMC 20 images were obtained (see Martayan, Baade & Fabregat 2010). Let us recall that this kind of instrumentation is not sensitive to the diffuse ambient nebulosity and is not sensitive enough to the weak emission, thus only lower estimates of the emission-line stars content can be provided. Finally 3 million spectra were obtained in the SMC (covering 3 square degrees), and 5 million in the LMC. In this first part of their survey, Martayan, Baade & Fabregat (2010) focused in 84 SMC open clusters. Once the emission-line stars and normal stars detected, the stars were classified.

-1) the astrometry (∼0.5") was performed with the ASTROM code [START_REF] Wallace | Users guide of ASTROM Watson[END_REF].

-2) the WFI stars were cross-matched with Ogle-II photometric catalogues (Udalski et al. 1998). For the stars successfully correlated, the B, V, and I colours were obtained.

-3) for each open cluster, the E[B-V] and the age were obtained from Pietrzynski & Udalski (1999), then the photometry of the stars was corrected of the reddening.

-4) the absolute magnitudes of the stars were obtained with the SMC distance modulus (Udalski 2000).

-5) ∼4400 stars in SMC open clusters were classified using the calibration of Lang ( 2001).

-6) The SMC data were compared to the results from McSwain & Gies (2005) in the Galaxy. At that step it is possible to compute the ratios of Be stars to B stars per spectral-type categories, i.e. for example B0e/(B0+B0e), etc, in the SMC and MW. Finally Martayan, Baade & Fabregat (2010) found that the Be stars are 3 to 5 times more abundant in the SMC (at low metallicity) than in the MW.

• Metallicity (Z) and rotational velocities Keller (2004), Martayan et al. (2006Martayan et al. ( , 2007)), Hunter et al. (2008) found that OBBe stars rotate faster at lower Z. The SMC OBBe stars rotate faster than LMC OBBe stars, which rotate faster than their Galactic counterparts. This is due to a lower mass-loss (the stellar winds are radiatively driven winds, see Bouret et al. 2003, Vink 2007) and lower angular momentum (Ω) loss at lower Z resulting in faster rotation (Maeder & Meynet 2001).

The ZAMS rotational velocities for SMC, LMC, and MW Be stars samples were determined by Martayan et al. (2007). Be stars at their birth rotate faster in the SMC than in the LMC, which rotate faster than in the MW. This is an opacity effect, at lower Z, the radii should be smaller, thus for an identical intial angular momentum, the stars rotate faster.

Theoretical tracks for Be stars at Z SM C from Ekström et al. (2008) fairly agree with the mean Vsini and mean VZAMS of SMC Be stars (see Martayan et al. 2010).

Long soft Gamma Ray Bursts and relationship with Be/Oe stars

The gamma ray bursts (GRBs) are the most energetic events since the Big Bang. Among them, 3 classes are distinguished. The type 1 or short GRB (less than 2s) probably resulting of the collapse of 2 compact objects. The type 2 or long GRB (LGRB) is associated to the SNIc and is probably resulting of a massive fast rotating star collapsing in a black hole (Woosley 1993, Fryer 1997). And the rare type 3 pseudo LGRB is not associated to a SN. New theoretical models (Hirschi, Meynet & Maeder 2005, Yoon, Langer & Norman 2006, Georgy et al. 2009) provide some informations about the LGRBs progenitors masses and rotational velocities from the ZAMS at different Z: they must be massive fast rotator at low Z (to avoid too much mass-loss and Ω loss), the H, He envelopes should be lost by fast rotation mixing and/or by a quasi chemically homogeneous evolution [START_REF] Maeder | [END_REF], Yoon, Langer & Norman 2006;Woosley & Heger 2006), and the stars should have a weak magnetic field.

From the observations point of view, Iwamoto et al. (1998Iwamoto et al. ( , 2000) ) found that the massive fast rotating stars are at the origin of the LGRBs. Thöne et al. (2008) found that the LGRB060505 occured in a low-Z galaxy, with high star-formation rate, in a young environment (6 Myr) and from an object about 32 M ⊙ . Campana et al. (2008): the LGRB060218 progenitor had an initial mass of 20 M ⊙ (∼B0 star), with Z = 0.004 (∼ Z (LM C,SM C) ). Martins et al. (2009) observed several SMC WR stars whose evolutionary status and chemical properties can be understood if they are fast rotators with a quasi chemically homogeneous evolution. One can compare the theoretical models at the ZAMS of Yoon, Langer & Norman (2006), that show areas in rotational velocities vs. the masses for different Z of LGRBs progenitors, with masscategories of Be/Oe stars. In Fig. 1, such comparison is done using the determined ZAMS rotational velocities for Be stars (ABCD) and for massive Be and Oe stars (E) in the SMC by Martayan et al. (2010).

• Comparison theory/observations of LGRBs progenitors areas

From observational and theoretical clues it seems that massive fast rotators at low Z, i.e, massive Be and Oe stars could be progenitors of LGRBs. The next sections provide additional tests for supporting this finding, including the predictions of the LGRBs rates and numbers in the local universe.

• Prediction of the LGRBs rates and numbers from Oe/Be stars populations To predict the rates of LGRBs from a well defined population, here the massive Be and Oe stars, different assumptions and computations were performed:

-the SMC is a representative galaxy of the Im.

-The number of OB stars in the SMC was determined using the SMC OGLE-III catalogue that is complete down to B9.

-The corresponding number of Oe/Be stars is determined with the rates from Martayan, Baade & Fabregat (2010).

-The binaries were removed of the sample and the transience of the Be phenomenon was taken into account -The lifetime of different spectral type stars was taken into account.

This provides the base rate of LGRBs, then the LGRBs beaming angles distribution from Watson et al. ( 2006) is used. In such case, the average predicted rate is: R pred LGRBs= (2-5) x 10 -7 LGRBs/year/galaxy.

These rates should be compared to the observed value: R obs LGRB ∼ (0.2-3) x 10 -7 LGRBs/year/galaxy (Zhang & Mészáros 2004;Podsiadlowski et al. 2004;Fryer et al. 2007).

One can then predict the number of LGRBs in the local universe, here for a redshift z ≤ 0.2 by taking into account that:

-there are 17% of Im for z<0.5 (Rocca-Volmerange et al. 2007), -the number of galaxies with z ≤ 0.2 comes from Skrutskie et al. (2006), -the previous rates are used.

-11 years are considered from 1998 to 2008, years for which the follow-up of the GRBs is the best (in term of redshift and classification) in the GRBOX survey1 .

In such case the predicted number is N pred = 3-6 LGRBs in 11 years at z≤0.2. The predicted numbers could be compared to the observed number of LGRBs in 11 years at z≤0.2 from the GRBOX survey: N obs = 8 LGRBs.

From these different tests and comparisons it seems that the populations of low Z massive Be and Oe stars could play a major role in the explanation of the LGRBs.

Remaining questions and conclusion

However, among other there are remaining important questions:

-What is actually the stellar/chemical evolution at low-Z with fast rotation? Some discrepancies between the theory and the observations were found by Hunter et al. (2009).

-The LBV stars could explode in SN without WR phase (Smith et al. 2010), that is not understood by the theory.

-Which observational clue could be found for the chemical evolution? Maybe from the GRB environment according to Woosley & Heger (2006) and Georgy et al. (2009). However, the fast rotating stars such as some SMC massive Be/Oe stars have Vsini ≥ 500 km/s, which according to the models could imply a quasi chemically homogeneous evolution.

-Is the SMC really representative of all the Im? Are the proportions Be/Oe similar in all SMC-like galaxies? Bresolin et al. (2007) observed OB stars in IC1613 (Z IC1613 ≤ Z SM C ). They found 6 Be stars among the 6 main sequence B stars observed.

-How many LGRBs/SGRBs are not detected by the Gamma Ray observatories? -There are more LGRBs at high-z. The first stars (very very metal-poor stars) were also probably very fast rotators (like Be/Oe stars?) and thus were probably LGRB progenitors.

Figure 1 :

 1 Figure 1: Comparison of LGRBs progenitors area (dashed) at Z SM C at the ZAMS with categories of SMC Be/Oe stars (grey-plain ABCDE areas). This Figure was adapted from Martayan et al. (2010).

http://lyra.berkeley.edu/grbox/grbox.php


Luis ; Olmos

Didier Bouvard

Tetsu Takahashi

Christophe L Martin

Luc Salvo

Daniel Bellet

Marco Di Michiel

Didier Bouvard

Analysing the sintering of heterogeneous powder structures by in situ microtomography

Keywords: composite materials, powder metallurgy, sintering user supplied): X-ray microtomography, Porosity Sintering, X ray microtomography, composites, porosity

come

In the last decades, progresses in understanding the physics of metal and ceramic powder sintering have been largely supported by the development of modelling, from prominent two sphere models [1][2][3] to more recent discrete element simulations [4,5]. Experimental investigations have not been as successful in capturing operative mechanisms because of the difficulty of directly observing the sintering process. Generally, sintered powders are observed post mortem by scanning electronic microscopy (SEM). With such a procedure, statistical data on the advancement of sintering can be obtained through the comparison of microstructural features at various points in the sintering cycle. However, following the alteration of an individual particle or of an individual pore during sintering is not possible.

Additionally, the information directly provided by this technique is two-dimensional. Hence it is difficult to quantitatively investigate issues that are related to the collective behaviour of particles such as rearrangement and coordination evolution. The importance of particle rearrangement has been demonstrated early on by Petzow and Exner, who have measured particle rotations in 2D packing. [6].

X-ray microtomography now provides a non-destructive method to observe the microstructure of materials in 3D. This technique consists of recording radiographs of a specimen at various angular positions. The spatial distribution of the linear attenuation coefficient within the specimen is deduced from a mathematical algorithm. A 3D image of the microstructure is thus obtained when the phases composing the material have distinctive attenuation coefficients [START_REF] Kak | Principles of Computerized Tomographic Imaging[END_REF]. This is particularly convenient for powder materials where the pore phase is transparent to the X-rays. The use of the synchrotron radiation from third generation sources allows highly absorbing materials such as copper and iron to be characterized with a spatial resolution in the range from one to two micrometers [START_REF] Bilderback | [END_REF]. Also, if the time of image recording is convenient, in situ observation of a material under mechanical or thermal loading is possible.

The first study of a sintered powder material by synchrotron microtomography has been performed by Bernard et al. [9] with glass beads sintering by viscous flow. These authors obtained 3D images of a specimen which sintering had been interrupted after various times.

More recently, Lame et al. [10] and Vagnon et al. [11] achieved in situ characterization of the sintering of copper and steel metal powders with particle size of approximately 50 µm. After a quantitative analysis of the data obtained on copper powder [12], they found that sintering was quite homogeneous and that classical models based on mean field assumption agreed well with the experimental data. Nöthe et al. [13] confirmed the uniformity of sintering when he found limited amount of rotations of copper particles in 3D microtomography images taken at various stages of sintering. It should be pointed out that all of the aforementioned studies are concerned with quasi monosize, closely packed powders.

The present paper presents results obtained by in situ synchrotron microtomography with more complex copper-based powder systems. These systems comprise particles with a broad size distribution and some of them include defects such as artificially created macropores or ceramic inclusions. One may anticipate a large effect of particle rearrangement for such systems with local heterogeneities. Novel tools have been developed for a more comprehensive analysis of 3D images, which involves in particular the evolution of interparticle contact indentation and the displacement of particle centres throughout sintering.

The data drawn from this analysis allow discussion of the sintering behaviour of heterogeneous powder systems with particular emphasis on collective phenomena.

Experimental procedure and image processing

Microtomography experiments were carried out at the European Synchrotron Radiation Facilities (ESRF) in Grenoble, France. Because copper is a material with a high absorption coefficient, it has been necessary to use high energy X-ray, which has been provided by ID15 beamline of ESRF. At ID15, the incoming X ray energy ranges between 35 and 90 keV after the filter. In the present experiments, a 50 keV energy was used. More details on the set-up can be found elsewhere [10,11,14]. The experimental set up allows the sample to be translated and rotated by 180°. A specialized furnace has been built at ESRF for sintering. This furnace is composed of a horizontal ceramic cylinder surrounded by a resistive wire. Three holes, one at the bottom of the furnace to allow for the introduction of the sample and two on both sides for the beam passage are located on the ceramic cylinder. Images were obtained with a 1 micron thick europium-doped Lutetium Aluminium Garnet (LAG) single crystal scintillator screen located in front of the detector system. The images provided by the scintillator were magnified 9 times with a microscope objective and recorded by a 1024 x 1024 CCD Frelon camera. The CCD camera has the following characteristics: a pixel size of Three samples have been investigated: a close packed copper powder, a copper-based powder composite and a copper powder compact with artificial pores. The atomised spherical copper powder was sifted to retain only particles smaller than 63 µm (Figure 1a) and then poured into a quartz capillary with 1 mm internal diameter and approximately 1 cm height. This capillary was pasted to an alumina rod attached to the positioning system. The second sample was composed of slightly sintered copper particles with artificially created pores. The procedure to create these large pores consisted of mixing the copper powder with 15% volume fraction of NaCl salt particles of 40 -63 µm diameter. This mixture was then pre-sintered at 500°C for 30 minutes in a 1 mm diameter capillary, which was removed after cooling. The salt was finally dissolved by keeping the samples in distilled water during 5 hours. The resulting sample is displayed in Figure 1b. The sample was introduced in a quartz capillary for subsequent sintering. The third type of sample, the composite sample, was made of copper particles mixed with 16% volume fraction of spherical porous alumina particles (Figure 1c) with a diameter between 40 and 60 µm. These particles were mixed for one hour in a turbula and then poured in the same type of capillary as the one used for copper powder. The volume fraction has been selected based on the results of a previous study on the densification of various copper-alumina powder mixtures by dilatometry and scanning electron microscopy [15]. This prior study showed that the reduction of the densification rate due to the presence of alumina inclusions was already significant with 5 vol.% inclusions and was particularly drastic when the inclusion content was above 20%. Indeed, significant clustering of inclusions was observed for alumina content above 15%.

Sintering for all three types of sample was carried out at 1050°C after heating at 40 K/min under a reducing atmosphere (helium -5 at.% hydrogen) to prevent oxidation of copper particles. The duration of the sintering plateau was different for the three samples. Because of the thermal expansion of the powder and of its swelling at intermediate temperature [16], the sample translated with respect to the X-ray beam during the observation. Thus, a slice of only 600x600x300 voxels could be tracked all along the thermal cycle. 3D raw images were treated in order to separate the phases by thresholding. The separation of pore and copper phases was straightforward and provided binary images of regular and porous copper samples. Concerning the composite material, alumina particles and porosity can hardly be distinguished by eye and not at all by automatic processing because the grey levels of alumina particles and porosity are very close. Hence, for this sample, we have prepared binary images where alumina inclusions and porosity appear as the same phase. This is unfortunate since it prevents the investigation of interesting phenomena such as inclusion coordination and clustering. Still, useful information concerning evolution of copper particles, which constitute the "active" phase of the composite material, has been drawn from binary images, as demonstrated in the following.

Copper particle segmentation was achieved by a watershed procedure applied on binary images (for more details refer to Vagnon et al. [12]). The volume and the centre of gravity of each particle were then calculated from the segmented images. Contact indentation was defined as the relative displacement of particle centres. We evaluated the diameter of each particle as the diameter of an equivalent sphere with the same volume. Because of the presence of defects on the images (ring artefacts in particular) and of computer memory capacity, only ∼500 particles have been followed during the whole sintering cycle. We present in Figure 2 an example of 2D images (virtual slices) extracted from 3D images showing the same microstructure after reconstruction, binarisation and segmentation. The segmentation procedure has proved more to be more problematic for increasing densities since particle interfaces become less clear. The upper acceptable limit was found to be around 0.75 in relative density. Such a limit is certainly higher in a monosize packing since the size of every particle to be separated is known. has been presintered before being introduced in the capillary. The radial shrinkage of each sample after sintering is clearly observed. The shrinkage of the composite sample is much smaller than the shrinkage of the other samples. Two types of pores appear in the porous sample: regular pores resulting from particle packing, as in the regular samples, and larger pores created by salt particle elimination. The regular pores exhibit a large decrease in size during sintering whereas the changes of the large pores are less significant. As mentioned above, the alumina particles can hardly be distinguished from the pores in the composite sample. Consequently, the alumina particles are incorporated in the porous phase in the following calculations. Obviously, these "filled pores" cannot shrink as the large pores of the porous sample.

Results

General observations

Table I shows various data for every specimen: the initial relative density, D 0 , the relative density after the heating stage, the final relative density, D f , the final densification ratio, defined as (D f -D 0 )/ D 0 , the initial and final average particle sizes (averaged by number), the initial and final coordination numbers, Z 0 and Z f . The relative density is defined as the weight density of the sample divided by the weight density of pure copper. The average particle size is initially between 33 and 42 µm, depending on the sample. It increases slightly during sintering due to the disappearance of the smallest particles. The initial densities of the regular and of the composite samples are equal because copper and alumina particles have approximately the same size. The initial density of the porous sample is 14% lower than the initial density of the regular sample. The final densification ratio for the regular sample is 0.24, sligthly larger for the porous sample (0.31), and much lower for the composite sample (0.1).

Coordination number

As expected, the initial coordination number is lower for the porous material than for the regular one (5.9 vs. 6.8). It is even lower in the composite material (5.6) because only coppercopper contacts are accounted for. It is worth noting that multiplying the initial coordination number of the regular sample by the volume fraction of copper particles in the composite sample leads to a coordination number of 5.7, which is very close to the one calculated for the composite material. To compare the coordination number evolution in the three materials, we have plotted in Figure 4 by their initial values. Interestingly, all points lie close to a straight line of slope 1 which means that the coordination number is well approximated by the very simple relation:

0 0 D D Z Z = (1)
It is shown in Figure 4 that this expression does not agree with the classical relation obtained by Arzt [17] from the assumption that the coordination number varies in accordance with the Radial Distribution Function (RDF) characterizing particle packing :

        -         + = 1 5 . 15 3 / 1 0 0 D D Z Z (2)
Arzt's relation significantly underestimates the increase in coordination found in microtomography specimens.

Pore size distribution

The pore size distribution is an important marker of the progress of sintering, especially so for heterogeneous systems such as those studied here. It should be pointed out that the porosity in such materials is completely open. It means that the pores cannot be considered as distinct elements and that the definition of a pore size distribution is somewhat ambiguous. Anyway, following Matheron [START_REF] Matheron | Eléments pour une théorie des milieux poreux[END_REF], such a distribution can be estimated by measuring the porous volume accessible to a structural element of increasing size. For this purpose, we chose an octahedra of edge 3.2 µm (twice the voxel size) and multiples of this. Figure 5 shows the pore size distribution obtained in this way at various times during the sintering cycle for each of the three samples. The volume fraction corresponding to a given size interval is calculated as the volume inaccessible to the octahedron of the larger size of the interval minus the volume fraction accessible to the lower size. Each volume fraction is normalized by the total volume of the particles, which is constant throughout sintering. Using this measure, the global pore volume decrease related with densification is clearly detected. In the regular sample (Figure 5a), pore size is initially between 0 and 40 µm and most pores have a size around 11 µm.

During sintering, the volume of the largest pores decreases and eventually pores between 30

Deleted: of In the porous sample (Figure 5b), the pore size is initially between 0 and 70 µm. The distribution exhibits a peak at 11 µm, as in the previous case, but also shows a clear plateau between 30 µm and 50 µm, which corresponds to the artificially created pores. These large pores get slightly smaller during the cycle. Meanwhile, the volume fraction of regular pores decreases with a maximum again at 11 µm. The pore size distribution in the presintered sample before sintering is very close to the one at the beginning of the 1050°C plateau.

Concerning the composite sample (Figure 5c), whatever the sintering time, the pore size distribution exhibits two clear peaks around 14 and 60 µm, respectively. The value of 14 µm is larger than the value of 11 µm found in the previous samples. This is because the average particle size was fortuitously larger in the composite sample. The second peak features single alumina particles and pairs of particles. A third peak, although of lesser extent, can also be detected close to 100 µm. It indicates the presence of a few triplets of alumina particles, which form volumes approximately twice as large as the volume of a single particle.

Interparticle indentation

In classical analysis derived from two sphere models, densification occurs during sintering because contacting particles get closer to each other through various diffusion mechanisms such as grain boundary diffusion and volume diffusion. The relative displacement of the centres of the two particles in contact with each other is defined here as the interparticle indentation and is denoted as h. According to the two sphere models, for two contacting particles with the same radius, R, the normalized indentation, i.e. the local deformation, h/2R, is a power function of R, with a negative exponent depending of the predominant diffusion mechanism. For particles with different sizes, R p and R q , Parhami et al. [START_REF] Parhami | [END_REF] have suggested that the relevant parameter to account for size effects is the equivalent radius defined as: This definition has also the advantage of being in line with the general contact theory for elasticity and plasticity [START_REF] Johnson | Contact mechanics[END_REF]. The pertinence of this parameter has been checked by using the following procedure. For each pair of contacting particles, we have determined the local deformation as the indentation divided by the distance between particle centres, d 0 . We have found large scattering in values of this parameter at any given time, likely due to local heterogeneities in particle shape and arrangement. We have thus calculated the average of the local deformation for pairs of particles with radii R p and R q for which the equivalent radius has a given value. We have chosen three values of equivalent radius: 14, 20 and 28 µm. For each value, at least 30 pairs of particles were tracked. Note that, for example, the 14 µm equivalent radius set includes a pair with particles of 9 and 27 µm radii as well as pairs with particles of the same radius, 14 µm.

q p q p R R R R R + = 2 ˆ (4)
The average local deformation for the three samples is plotted in Figure 6 as function of the actual densification ratio defined as (D-D 0 )/D 0 . The use of the densification ratio allows for a fair comparison of the different materials. In each case, the deformation increases with increasing densification. Also, as expected from sintering theory, the deformation is larger for lower equivalent radius whatever the sample is considered. However, the curves associated with each sample show clear differences. For regular and composite samples, the deformation is very low until a densification ratio of approximately 0.025 is reached. This corresponds to the heating stage. The deformation then regularly increases throughout densification. For the same equivalent radius, larger local deformation is observed in the composite sample as compared to the regular sample at a given densification ratio. For the porous presintered sample, the deformation first exhibits a steep increase followed by a more gradual increase. It is smaller than the deformation in regular and composite samples, except at low densification ratio. Standard deviation bars plotted in Figure 6 indicate that the data corresponding to each equivalent radius are clearly separated. This observation suggests that the equivalent radius is able to reveal the size effect in sintering deformation.

Particle tracking

We take advantage of the ability of X-ray tomography to track individual particles to gain knowledge on particle rearrangement during the whole thermal cycle. We can draw, for example, the trajectories of particle centres in the analysed volume. Since 3D lines are hardly discernible in a 2D plot, we show in Figure 7 the horizontal projections of the trajectories. According to the mean field theory, which assumes that the shrinkage is uniform throughout the sample, particle trajectories should be radial and particle displacement should be proportional to the distance between the particle and the volume centre. Most experimental trajectories seem to follow this rule but it is interesting to note that some of them significantly deviate from the radial direction. This is particularly obvious in the porous sample where a few trajectories are perpendicular to the radial direction.

For a quantitative analysis of particle displacement, we have calculated two parameters for each particle:

-the "deviation angle", θ, defined as the angle between the displacement vector -going from the initial position of the particle centre to its position at a given sintering time -and the radial direction, between 0 and 180°;

-the "closing parameter", δ = 1 r a u u -, where u r is the radial component of the displacement vector and u a is the length of the displacement vector predicted by the mean field theory.

When the particle gets closer to the volume centre, exactly as predicted by the mean field, δ is equal to 0.

As an example, Figure 8 shows the values of both parameters for 350 particles analyzed in the porous material at a density of 0.62 (densification ratio : 0.1). It is observed that 90% of the particles exhibit a deviation angle between 10° and 70° and a closing parameter smaller than 0.7. There is no clear correlation between both parameters except for their largest values.

Indeed, any particle with a deviation angle larger than 90° shows a closing parameter larger than 1, since it moves backward (u r < 0) instead of following the forward mean field displacement (u a > 0). There are only five particles with a deviation angle above 90°.

The average values of δ and θ have been calculated at various sintering times to quantify the relative amount of rearrangement for each sample. These values are plotted vs. the densification ratio in Figure 9. We observe that δ and θ averages show first a steep increase followed by a continuous decrease. The curves corresponding to the regular and composite samples are very close to each other whereas those corresponding to the porous specimen are always above the other two, except in the low densification ratio domain. However, it is not relevant to compare the beginning of the curves since the porous material experienced a presintering step before its introduction in the microtomography experimental set up. averages tend towards constant values, respectively, 0.14 and 18° for the regular material; 0.2 and also 18° for the porous material. For the composite material the final densification ratio is rather small and there is no clear asymptotic value.

Discussion

Quantitative data drawn from microtomography experiments provide useful information for a better understanding of microstructure evolution of metal powders during sintering. The relevance of the most interesting results is discussed in the following.

First, we have found a simple linear relation that reasonably describes the increase of average coordination number in the three investigated samples as function of relative density. It is worth pointing out that the same relation is relevant for very different materials such as a close packing of particles, a packing with large pores and a composite packing. This linear relation may be justified by using Arzt's Eq. ( 2) with an undefined factor α replacing 15.5.

Linearization of this equation for D close to D 0 leads to :

0 0 0 3 D D Z Z D α - = + (4)
where α is a constant. Eq. ( 4) yields Eq. ( 1) for α = 3Z 0 , i.e. , α = 20.4 for our regular material. Eq. ( 1) is thus formally close to Arzt's model when the slope characterizing the Radial Distribution Function is increased. Also, Eq. ( 1) gives some credence to the linear empirical equation of Helle et al. [START_REF] Helle | [END_REF].

D Z 12 = (3)
With the values of Z 0 and D 0 obtained for the regular material, we obtain Z = 10.5 D.

Next, we have measured the pore size distribution according to Matheron definition [START_REF] Matheron | Eléments pour une théorie des milieux poreux[END_REF]. For the regular sample with close packed particles, this distribution shows a clear peak around 11 µm throughout densification. This indicates that, although the powder densifies and the total amount of porosity decreases, the majority of pores consistently exhibit an equivalent size, which is approximately 27% of the mean particle size (40 µm). It does not mean however that 11 µm size pores are stable. All pores in the packing shrink but, as the largest pores shrink more slowly, the peak of the distribution remains the same. This trend has been observed for relative densities below 0.8. It is expected that the pore size decreases when the density approaches unity. In the other samples, a peak also appears in the distribution at approximately the same size value but a second population of much larger pores is also detected. In the porous sample, the size of the artificially created pores is in between 30 and 70 µm. These pores merely show a slight reduction in size during sintering. The analysis of composite sample data is not so pertinent since alumina inclusions have been lumped together with the porosity in the image analysis. We find once again that real pores exhibit a stable size, which is slightly larger than in the previous cases. Inclusion-filled pores also show evidence of a stable size of approximately 60 µm.

To analyse the indentation between particles of different sizes, an equivalent radius defined as the harmonic mean of both radii has been introduced. We have found that in average the normalized indentation -i.e. the local deformation -increases when the equivalent radius decreases. This means that the equivalent radius allows size effect in sintering to be taken into account in a simple manner. For the same densification ratio, in the porous sample, the local deformation is lower, whatever the value of the equivalent radius. This result demonstrates that the contribution of interparticle indentation to the densification is less important in the porous sample than in the regular one. In other words, a more significant part of the densification seems to come from particle rearrangement in the porous sample. In contrast, a larger indentation is required in the composite sample to reach a given densification ratio.

This is clearly linked to the presence of non-deformable inclusions that impose copper particle to sinter more for the sample to attain a given macroscopic relative density.

The evaluation of particle rearrangement has been carried out by analysing the deviation of particle displacement with regard to the mean field theory. .For this purpose, we have defined two parameters, one to characterize the deviation from the radial direction and another one to characterize the radial component of the displacement. These parameters have been calculated for all particles in the investigated volumes and for each specimen. Average values of these parameters have been calculated for quantification of particle rearrangement. The analysis of the variations of both parameter averages prove that rearrangement mainly occurs at the beginning of the densification. However, the three materials show dissimilar behaviours. For the regular and the composite samples, rearrangement is most important during the heating stage. This is because contact size, at this early stage of sintering, is very small and thus does not impede particle motions. Also, the packing density is still small enough at this stage to allow for particle rearrangement. For the porous material, which has been presintered, the rearrangement peak is observed later, after a few minutes at the temperature plateau. A reasonable interpretation is that this peak results from the conjunction of low packing density and high temperature. In any case, the values of average parameters at the end of the sintering plateau prove that rearrangement has been globally more important in the porous specimen in comparison with the regular and composite specimens. This last result is consistent with the indentation data described above, which showed that particle indentation was smaller in the porous material for a given densification ratio.

Conclusion

It has been shown that in situ X-ray microtomography was particularly appropriate for gathering data on collective particle behaviour during sintering. The wealth of information that this technique provides in 3D allows a better understanding of phenomena such as the evolution of multiple contact deformation or particle rearrangement. These phenomena play a significant role in sintering, and more particularly in heterogeneous powder systems as demonstrated here. The most relevant outcomes are :

-The estimate of a prominent pore size equal to 27% of the average particle size, which is

1. 4

 4 µm, and a full image readout time of 200 ms per image. 650 radiographs were recorded at covering an interval of 180° for each data set, which led to the capture of one microtomography data set in about one minute. With such an experimental set up, images with 1.6 µm sized voxels have been obtained.

Figure 3 shows

 3 Figure3shows 2D images (virtual slices) of the microstructure inside every sample before and after sintering. Copper particle are easily discernible in the initial state but not so easily in the final stage. The quartz cylindrical capillary can be seen in the images. Regular and composite samples are initially in contact with the capillary because the powders have been directly poured into it. Conversely, the porous sample is smaller than the capillary because it

  the coordination number versus the relative density, both normalized

  disappear. However the 11 µm pore size remains the most represented size in the distribution. Note that most of the pore changes occur during the sintering plateau at 1050°C.

  constant at least up to 0,8 of relative density, -the minor shrinkage of artificially introduced large pores, -the linear relation between the coordination number and the relative density, for all samples, including the composite packing, -the relevance of the defined equivalent radius to describe the sintering of pairs of particles with different sizes, -the definition of two parameters related with particle displacement that allowed for a quantification of rearrangement, -the conjunction of a slightly stronger rearrangement and a lower interparticle indentation in the porous material in comparison with the regular one at the same densification ratio.

Figure 1 :Figure 2 :Figure 3 :Figure 4 :

 1234 Figure 1 : SEM observation of copper powder (a), presintered porous sample (b), alumina powder (c) Figure 2 : Virtual slice of a regular copper sample (a) after reconstruction and (b) after binarisation and segmentation Figure 3 : Virtual slices of samples in the initial state (left side) and after sintering (right side).(a) et (b) : regular sample; (c) and (d) : porous sample; (e) and (f) : composite sample Figure 4 : Normalized coordination number vs. normalized relative density.

Figure 5 :Figure 6 :

 56 Figure 5 : Pore size distribution at various times of the sintering cycle for regular (a), porous (b) and composite (c) samples Figure 6 : Average normalized particle indentation as function of the densification ratio for three values of equivalent radius, 14 , 20 and 28 µm, (a) regular, (b) porous and (c) composite.

Figure 7 :Figure 8 :

 78 Figure 7 : Planar projection of particle centre trajectories in the horizontal plane during sintering in regular (a), porous (b), composite (c) samples and schematic definition of displacement parameters (d) Figure 8 : Distribution of the parameters characterizing particle rearrangement, the deviation angle θ and the closing parameter δ, in the porous specimen at a densification ratio of 0.1.

Figure 9 :

 9 Figure 9 : Evolution of the average values of δ (a) and θ (b) during densification for the three samples.

  Figure 1

Figure 5 Page

 5 Figure 5

Table caption :

 caption The investigation methods presented in this paper are also useful for guiding the development of modelling techniques such as discrete element modelling[22]. They can provide both It should be remembered that this work has been restricted to the initial stage in sintering for two reasons. First, individual particles cannot be separated from 3D images when the relative density is larger than 0.75. Secondly, due to microtomography resolution, we have been compelled to sinter relatively coarse powders, which can hardly been densified above 85% relative density in a reasonable time. With a higher resolution in the future, we should be able to work with finer powders and then investigate denser structures, thus allowing the transition from open to close porosity together with the final stages of pore evolution to be analysed. Table I : Data characterizing the three samples.

	information for building realistic numerical microstructures and data for validating the results
	of the models.
	Deleted: commented
	F o r F o r F o r
	P P P
	e e e
	e r e r e r
	R R R
	e e e
	v i e v i e v i e
	w w w
	O n l O n l O n l
	y y y

Table I :

 I Data characterizing the three samples.

		0.645	0.665	0.800	0.24
	Porous	0.563	0.57	0.735	0.31
	Composite	0.650	0.673	0.713	0.10
	Material	Initial average	Final average	Initial	Final
		particle size	particle size	coordination	coordination
		(µm)	(µm)	number, Z 0	number, Z f
	Regular	37	43	6.8	8.6
	Porous	33	40	5.9	7.7
	Composite	42	44	5.6	6

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