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Université de Montréal
Étude des environnements circumstellaires d' étoiles Ae/Be de Herbig et d'autres étoiles jeunes
Par Roger Hajjar Département de physique Faculté des arts et des sciences
Thèse présentée à la Faculté des études supérieures en vue de l'obtention du grade de Philosophie Doctor (Ph.D.) en physique
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Cette thèse intitulée:
Étude des environnements circumstelIaires d' étoiles Ae/Be de Herbig et d'autres étoiles jeunes
présentée par:
Roger Haijar
a été évaluée par un jury composé des personnes suivantes:
Gilles Fontaine, président-rapporteur Pierre Bastien, directeur de recherche Nicole St-Louis, membre du jury Jerôme Bouvier, examinateur externe
Thèse acceptée le: . . .. 30. .. . r. .. / . .? .-.. .(. 4 .? f- SOMMAIRE
Cet ouvrage présente une étude des environnements circumstellaires d'étoiles pré-séquence principale de masse intermédiaire dites he/Be de Herbig ainsi que d'objets stellaires jeunes de différents types, des T Tauris. des FU Ori et d'autres de types encore indéterminés. Nous avons recouru pour cela à deux techniques d'observation: la photométrie et la cartographie dans le continu millimétrique et la cartographie en polarisation linéaire.
Nous avons comparé. pour la première fois. des cartes de polarisation obtenues a différentes longueurs d'onde. allant du visible à l'infrarouge dans le but d'observer et d'exploiter une prédiction faite par Bastien S: Menard (BSI) voulant que la dimension du disque polarimétrique, représenté par un profil de vecteurs de pola- risation alignés et délimités par deux points nuls. varie en fonction de la longueur d'onde. Les points nuls indiqueraient la zone de transition entre le milieu optique- ment épais et celui optiquement mince. L'observation de cet effet à été constatée pour tous les objets de notre échantillon, indépendamment de Leur àge ou de leur type. Le seul critère d'appartenance a l'échantillon étant un angle d'inclinaison du disque de zs 90°. Nous avons également obtenu, en se basant sur l'interprétation du modèle de BM, le profil de densité dans le disque et, pour un objet, HL Tau, les densités effectives du disque. La masse trouvée à partir de ces calculs, MD O. 1 Mo. concorde avec les estimations obtenues par d'autres méthodes. Le profil de densité en loi de puissance obtenu, p(r) = po(ro/r)7où 7 = k&1.9, est comparable à ceux utilisés pour la nébuleuse protosolsire ainsi que pour les modèles de disques d'accrétion autour des étoiles jeunes. L'étude des cartes de polarisation nous a permis également le développement d'une méthode de détermination de la source dïiiumination de la nébuleuse. Elle se base sur le caicd des coordonnées de tous les points d'intersection des perpendiculaires aux vecteurs de polarisation. Ceux-ci formant majoritairement un patron centrosymétrique, l'intersection des perpendiculaires devrait indiquer la position de la source. Nos résultats semblent indiquer que les histogrammes en deux dimensions des points d'intersection portent de l'information sur la géométrie ainsi que les propriétés optiques de la source.
La cartographie millimétrique et submillimétrique a permis la découverte et la caractérisation d'une nouvelle source dans la région de V633 Cas. Elle a également montré la complexité de la distribution de poussière autour de V628 Cas. Ceci montre l'importance de la cartographie et l'identification de toutes les sources contribuant au flux électromagnétique dans une région. L'étude des étoiles jeunes se faisant premièrement à partir de leurs distributions d'énergie specr rale. une telle étude est donc fondamentale pour la compréhension de l'é~olutionde ces deux objets au moins.
Un autre résultat important vient de l'étude détaillée de HL Tau. Les observa- tions polarimétriques ont montré que les densités du disque proche de l'étoile sont suffisamment élevées pour rendre le milieu opaque à la radiation submillimétrique. Jusqu'à ce jour, on a togjours considéré, aux fins d'analyse des données millimétri- ques et submillimétriques, que le milieu était transparent à ces longueurs d'onde.
Mots clés: étoiles jeunes - polarisation: cartes - submillimétriques: cartes - environnement circumstellaire: disque. *. . SOMMAIRE ...... 111
LISTE DES TABLES ...... xi ... LISTE DES FIGURES ...... xiii -.. LISTE DES ABRÉVIATIONS ...... XVIII
-A\*-AYT-PROPOS ...... xx
ENTRODC~CTIOX ...... L
0.1 La formation stellaire dans ses grandes lignes ...... 3
0.2 Les structures du milieu circumstellaire ...... 9
0.3 Le millimétrique et submillimétrique ...... 12
0.4 L'imagerie polarimét tique ...... L.5
0.4.1 Uneintroduction ...... 15
0.4.2 L'interprétation et les modèles ...... 15
0.4.3 L'utilité des cartes de polarisation ...... 18
0.5 Les objectifs ...... 20
RÉFÉRENCES ...... 23 CHAPITRE 1: Observations millimétriques de V633 Cas et V628 Cas . 27
1.1 Introduction ...... 30
1.2 Observations ...... 32
1.2.1 Cdibration ...... 32
1-22 Sensitivities ...... 33
1.3 V633 Cas ...... 34
1.3.1 Morphology ...... 35
1.3.1.1 At thecenterofadisk:' ...... :36
1.3.1 The V633 Cas C Peak ...... 38
1.3.2.1 Dust properties ...... 38
1 .3.2.2 Mass calculations ...... 139
1.3.3 Discussion ...... 40
1.3.3.1 A Class O Object or an extreme Class I? ..... 41
1.1 Cr628 Cas ...... 44
1.1 . 1 Morp hology ...... 45
1.1.2 Dust Properties and Distribution ...... 16
1.5 Conclusions ...... 47
RÉFÉRENCES ...... 54
CHAPITRE 2: NGC 7129 ...... 68 vii
2.1 Introduction ...... 11
2.2 Observations ...... 73
-* 2.3 Results for the Optical and Near IR Observations ...... r a
2.3.1 PJoor ...... TCda
-* 2.3.1.1 . Direct Imaging ...... ia
LI 2.3.1.2 ?ohrimetrie Imaging ...... 4,
2.3.1.3 Aperture Polarimet ry ...... 79
2-32 LkHa234 ...... S1
2.3.2.1 Direct Imaging ...... 31
2.3.2.2 Polarimetric Imaging ...... S4
2.3.3 BD +6.j0 1631 ...... $3
2.3.4 BD +6s0 1635 ...... 36
2.1 %fillimeterand Submillimeter P hotomet ry ...... Si
2.5 Spectroscopy of 'Joor ...... Y8
2.6 Discussion ...... 89
2.6.1 Interpreting Polarization Maps ...... 89
2.6.2 Individual Objects ...... 91
2.6.2.1 Noor ...... 91
2.6.2.2 LkHa234 ...... 91 2.6.3 A Star Formation Scenario ...... 92
7 Summary and Conclusion ...... 93
2.8 MueUer Matrix Calculations ...... 9.5
29 Polarimetric Map Cent roid Determination ...... 97
2.9.1 Cdculating the centroid ...... 98
2.9.2 MainContributors ...... 99
RÉFÉRENCES ...... 109
CHAPITRE 3: HL Tau . une interprétation des cartes de polarisation . . 129
.3.1 Introduction ...... 132
.3.2 Observations ...... 13.5
32.1 Optical ...... 135
3.2.2 Subrnillimeter ...... 136
3.3 Polarization ...... 136
3.3.1 Pohrization maps ...... 136
3.3.1.1 Polarimetric Centroid ...... 1:38
3.3.2 An aperture poiarirnetry review ...... 139
3.4 Fax infrared dust distribution ...... 130
3.5 Discussion ...... 141
3.5.1 Nul1 points reviewed ...... 111 3.5.1.1 HL& ...... 115
3.6 Conclusion ...... 118
RÉFÉRENCES ...... 157
CHAPITRE 4: Cartes de Polarisation infrarouge ...... 173
4.1 introduction ...... Li6
4.1.1 Disks and Pseud-Disks ...... 177
4.1.2 Interpreting Polarization Maps ...... 178
4.2 Observations ...... 179
4.2.1 The Sample ...... 179
4.22 Observations ...... 1S0
4.2.3 Reduction and Calibration ...... 1S1
4.3 Results for Individual Objects ...... 1SL
4.3.1 Parsamian 21 ...... 182
3.3.2 LkHa 233 ...... 153
4.3.3 V376Cas ...... 185
4.3.5 RN0 138 ...... 189
4.4 Disk Density Structure ...... 191
4.5 Cdculating the Centroid Position ...... 192 1.6 Discussion ...... 194
4.6.1 The Disk ...... 194
4.6.2 Thestar's Position ...... 197
1.7 Conclusion ...... 198
REFERENCES ...: ...... 206
CONCLUSION ...... 23.5
RÉFÉRENCES ...... 110 ... RE>IERCIEMENTS ...... xxirt LISTE DES TABLES
1.1 Log of Observations ...... 19
1.2 Selected Gain and scaling factors for different runs and wavelengths . 50
1.3 V633 Cas C photometry and related parameters ...... 5 1
1.4 Mass calculation results ...... 52
1 .. V628 Cas photometry aad related parameters ...... -53
2.1 Log of Observations ...... 100
-1 3 LI- Nooros companion stars...... 101
2 ..3 Polarizat ion map centroid ...... 102
2 ..3 - continued ...... 103
2.4 Aperture polarimetry of NGC 7129 stars...... 104
2.5 LkHa 231 cornpanion stars...... 105
2.6 Coordinates of sources tracing the R boundary line...... 106
2.7 Millimeter photometry of sources in NGC 7129 ...... 107
2.8 Polarization maps and related parameters ...... 108
3.1 Log of Observations ...... 151 xii
3.2 Polarimatric disk size at different X ...... 153
3.3 Centroid calculations...... 1.53
3.4 Linear polarization measurements on HL Tau ...... 154
3.5 Density measurements in the disk for diflerent grain models ... 1.55
3.6 Disk masses for the three diflerent grain models...... 156
4.1 Program stars and related parameters ...... 'LOO
4.2 Disk sizes and density calculations ...... 701
4.3 Calculation results for centroid determination...... 202
1.3 -continued ...... 203
4.3 -continued ...... 203
4.3 -continued ...... 20.5 LISTE DES FIGURES
O. 1 Courbe d'extinction interstellaire ...... 1
O .2 Courbe de polarisation interstellaire ...... >
0.3 Comparaison des DES des OSJs et des EAEBEH ...... 5
0.4 Vue schématique des structures majeures du MCS ...... 11
0.5 Carte de Polarisation dans la bande 1 de HL Tau ...... 16
0.6 Cartes de polarisation théoriques, Fisher et a1 (1996) ...... 19
1.1 Calibration fits for each cun ......
1.2 Error distribution for the V633 Cas rnap ......
1.3 Error distribution for the V628 Cas rnap ......
1.3 Continuum map at 1.1 mm of V633 Cas ......
1.5 Continuum map at 1.1 mm of V628 Cas ......
1.6 Molecular condensations around V633 Cas C (adapted from Nakano et al 1990) ......
1.6 (b) ......
1.7 First SED estimate of the new source V633 Cas C ...... xiv
1.8 Emissivity index of V633 Cas C ...... 66
1.9 Emissivity indexes for peaks 1 and 2 on the V628 Cas map ..... 67
2.1 RbandimageofNGC7129 ...... 111
2.2 ùitensity contour plots of Noor ...... 119
2.3 Greyscale J and K images of Noor ...... 113
2.4 1 band polarization map of Noor ...... 111
2.5 J band polarization map of Noor ...... 11.5
2.5 (b) ...... 116
2.6 P(X) and @(A) curves for Noor ...... 117
2.7 Same as Fig . 2.2 but for LkHa '231 ...... 11s
'>.Y Same as Fig 2.3 but for LkHa 231 ...... 119
2.9 Sameas Fig . 2.4 for LkHa 231 ...... 1-0
2.10 Sarne as Fig . 2.3 for BD +6j0 1637 ...... 121
2.1 1 Sarne as Fig . 2.4 for BD +65" 1637 ...... 1-2
2.12 Sameas Fig . 2.2 for BD +65O 1638 ...... 123
2.13 Same as Fig . 2.4 for BD +65* 1638 ...... 124
2.14 Calibration fits for 3 mi!limeter and submiiiimeter wavelengths . . 125
2.15 Fit through the photometric measurements of Noor ...... 126
2.16 Noor's Spectra ...... 127 XV
2.17 IRS objects and cavity walls near LkHa 234 ...... 128
3.1 Histogram of errors for the 800pm map of HL Tau ...... 160
3.2 The 1 band polarization map of HL Tau ...... 161
3.2 (b) ...... 162
3.3 Polarkation vectors of HL Tau overlayed on diflerent types of isocontours ...... 163
3.3 (b) ...... 164
3.4 J . H and K polarization map of HL Tau (a. b and c repectirely) . iveintraub et al ( 1995) ...... 16.3
3 .4 (b) ...... 166
3.4 (c) ...... 167
3.5 Two dimensional histogram of the intersection points ...... 168
.3.6 The YOOprn map of HLTau ...... 169
3.7 .A zimuthal plot of al1 the pixels in the YODprn map ...... Li0
3.8 Schematics of the geometrical relation of the three cylindrical coordinates on the light path ...... 171
3.9 Logarithmic plots of the column density versus the radius ..... 172
4.1 J band polarimetric map of Parsamian 21 ...... 209
4.1 (b) ...... 210
1.1 (c) ...... 211 xvi
1.2 K band polarimetric rnap of Parsamian 21 ...... 212
4.2 (b) ...... 213
4.2 (c) ...... 211
1.3 Sameas figure 4.1 but for the K band mapof LkHa 233...... 215
4.3 (b) ...... : ...... 216
1.3 (c) ...... -.. . . . 217
1.4 Same as figure 4.1, for the K band map of V376 Cas ...... 2 18
4.4 (b) ...... 219
4.5 (c) ...... "'O
4.6 Contour plot of the J band of KI76 Cas ...... -,.>-- 1
7 Same as figure 4.1. for the h: band map of V633 Cas ...... -739---
4.7 (b) ...... -P>:3,,
4.7 (c) ...... , . 9-34,,
4.3 Same as figure 4.1, for the J band rnap of RSO 138 ...... -195,-
4.8 (b) ...... ->36,,
1.8 (c) ...... 227
4.9 Contour plots of the H and K images of RN0 138 ...... 998--
4.10 Peak displacement with wavelength of RN0 138 ...... 229
4.11 Logarithrnic plot of the column density versus the scaled disk diameter230 xvii
4.12 2D histogram of the density of intersection points of perpendiculars to a pure ellipse ...... 231
4.12 (b) ...... 232
4.13 Plot of the intersection points of perpendiculars to an ellipse with an axis ratio of 0.7...... '333
4.14 Density distribution of intersection points for al1 our polarization maps...... 23.1 OSJ Objet Stellaire Jeune YS0 Young Stellar Object IR Infrarouge PIR Proche Infrarouge ( NIR en anglais) TRL Infrarouge Lointain ( FIE€ en anglais) EAEBEH Etoiles Ae/Be de Herbig (HAEBES en anglais) MCS Milieu Circurnstellaire hl 1s hlilieu Interstellaire -... Heureur parmi les hommes D 'arorr iti chois1
Et par l'une et par IButre Je n'ai jamais connu
D'amour comme le uotre." Ires Duterl
A Roula et Sour AVANT-PROPOS
*...se tenant avec prudence dans le commerce d'un vieil arbre, appuyé du menton à la dernière étoile.
II uoit au fond du ciel de grandes . " choses pures qui tournent au plaistr. Saint-John Perse
Depuis la nuit des temps. l'univers avait été reconnu statique. immuable et inaltérable. Aucun savoir n'était possible quant à la composition des astres et leurs propriétés disait encore au dix-huitième siècle Auguste Comte. La mécanique de Yewton avait permis de déchiffrer le mouvement des astres qui. finalement. dans ce ballet, restait encore inaltérable. La physique et I'astronomie de La fin du dix- neuvième et du vingtième siècle mirent l'univers en branle et le parèrent d'une dimension et d'une richesse jusque là inimaginables. Les étoiles et galaxies y foisonnaient et les distances mesurées allaient en s'accroissant. Hubble montrait l'expansion de l'univers et confirmait ainsi la cosmologie, rendue enfin possible par la relativité générale. Des pionniers tels Eddington et Chandrasekhar calculaient les structures des étoiles et montraient qu'elles avaient une fin. Du spectre de Fraunhofer auexpériences de Bunsen et Kirchhoff, la voie était désormais ouverte à la co~aissancede la composition des étoiles. Celles-ci prenaient enfin vie, les étoiles naissaient, vivaient et mouraient. L'astrophysique était née et avec elle la possibilité de comprendre l'univers auquel nous participons. Les premières xxi mesures de polarisation interstellaire en 1949 par Hall et Hiltner pavèrent la voie. enfin, à l'utilisation de toutes les propriétés de la lumière dans cette quête scientifique.
11 fallut attendre longtemps avant de commencer à étudier les premiers stades de l'évolution stellaire. leur formation. En effet les premières identifications d'étoiles jeunes. communément appelées pré-séquence principale. remontent à 1958 avec Joy qui émet l'hypothèse qu'une certaine classe détoiles. appellées T Tauri, pourrait bien être les précurseurs des étoiles de masse comparable à la masse du soleil (0.5M- SM. 5 2Me). Cette découverte fut suivit en 1960 par I0identification d'une nouvelle classe d'étoiles pré-séquence principale par Herbig. Les étoiles de type Ae/Be de Herbig sont les précurseurs des étoiles de masses intermédiaires ('ME, SM. 5 10Mn).
Les efforts théoriques développés en parallèle ont mené à de premiers modèles de formation stellaire et à ce qui est communèment appelé les isochrones de Hayashi qui identifient le parcours d'évolution d'une étoile vers sa phase adulte. celle de la production d'énergie thermonucléaire.
Les études du milieu interstellaire furent également déterminantes pour la compréhension des processus de formation stellaire. La découverte de l'extinction interstellaire par Trurnpler en 1930 fut suivit par des études poussées sur la poussière qui peuple l'espace interstellaire. La polarisation fut et reste un instru- ment clé dans l'étude des propriétés de ces minuscules constituants de l'univers. L'ouverture de nouvelles fenêtres d'observation ainsi que le développement des techniques radio vint également montrer l'existence de nuages de matière, ultérieu- rement identifiés comme d'immenses pouponnières d'étoiles. Ce qui, d'évidence, amena les grains de poussières à mettre leur grain de sable dans la machine REMERCIEMENTS
Un jour d'août 1990. je suis entré dans un bureau au quatrièmeétage de l'aile D du pavillon principal. Ignorant tout de l'astronomie - tel que je l'ai découvert plus tard en suivant le cours d'astronomie galactique -. Pierre Bastien a bien voulu m'accepter comme étudiant dans son groupe de formation stellaire, je le fut pendant près de sept ans. Je lui dois d'avoir réalisé le rêve de devenir astronome. la chance d'avoir visité - et quatre fois s'il vous plaît - le site astronomique le plus important de la planète. le sommet du Mauna Kea à Hawaï. et le bonheur de travailler au CFH qui. pour l'adolescent que j'étais au début des années 80. symbolisait l'astronomie dans toute sa splendeur.
Un jour de mars 1992. déclarant devant la famille réunie que je ne m'engagerais dans un doctorat que si j 'obtenais une bourse, mon père me foudroya du regard. La bourse, je l'ai obtenue. mais je ne serais pas là sans le support moral et financier de mes parents. Après tout, les meubles de la chambre de Nour, les billets d'avion et notre séjour au Liban, toutes mes réalisations, de mes premiers gagas à ma thèse de doctorat, sont le fruit de leur éducation. Mes parents. c'est aussi mes beaux-parents. mon beau-père qui a accepté de passer trois mois aux frites et aux œufs rien que pour nous donner la chance d'avoir parmi nous au Canada. mon adorable Tante Adèle sans qui nous serions. Roula et moi, encore en train de se demander comment donner le bain à un bébé de trois jours!
Un jour de mai 1992. une femme bien assise dans son siège au bureau de Mkallès au Liban, accepta de me prendre comme époux. Ensemble, nous construisons la plus belle des réalisations. une famille dont le premier élément est Sour. notre petite étoile. et qui le restera même si la cornmunaut6 scientifique - <- -y--- n'adopte pas Noor comme nom de cette étoile jeune que Papa a étudiée. .\ toutes les deus. je dois mon bonheiir. un soutien et une présence de tous les jours et une raison de perséverer même aux jours où ma thèse. mon travail et tout mon être glissaient vers l'abime.
Un soir d'octobre 19S9. je juis arrivé au Canada. Je m'installais finalement au Québec. à Montréal plus précisément. le 31 décembre 1989. Et I'aventure de I'irnmigré que j'étais conimenqa. Je l'étais. oui. car je ne le suis plus. grâce à l'accueil estraordinaire quej'ai requ au sein du groupe d'astronomie de 1-L-niversité de Montréal et du Centre du hnt Slégantic. Je remercie tous ceux avec qui j'ai passé mes journées. étudiants et professeurs qui rn'ont accepté parmi eu\- comme
Louis d'un séminaire que présentait -notre Roger nationar ! Ca fait chaud au roeiir! Le Québec. hlontréal. le Groupe au sein duquel j'ai vécu. les Rencontres du Centre. les Bistros Brahé. lei .-\tdiers Étudiants. mon fameux billet d-infraction de $3 15. mon acquittement du dit billet et tous ceus que j 'ai plus ou moins connus re~terontcra\-és a jamais dans nion coeur.
Ln grand nierci a Luc Turbide qui a bien \*ouluque je branche rnori ordinateur au réseau quand je n'avais plus d'appartement. qui a mis au point le fichier LATES udm-t hese .sty pour nous faciliter la vie et la sienne aussi. .\krci pour son incroyable talent d'adminisrrareur de systême et parce qu'il a demandé qu'on le fasse dans ce petit manuel clii.il a écrit au sujet di1 fichier cité plus hait !
Je ni'inclinc dei-ant le talenr de Liniis Torvalds qui. étudiant de premier <->-rie i Iirlziiiki. a prwl~iitIr ~>.stZriivti'o~~Gratioii Li~,iis qui iii'a permis de çoriipléter ma thèse dans le cornfort douillet de mon domicile sur un Pentium 166 plus rapide que les stations de t rayai1 SL-S du groupe!
Khaled. Hanane. Alain. Sadine. Roland. Dan'. Sazha. Sami. Sleiman. Elie Sabil. Georges. Michelle. Jean-lfarc. Stefanir. Adel. Samer. -Aiméeet tant d'aut res sont de ces prénoms qui ont laissé leur trace sur ce travail. Ils ont contribué à la belle espérience que j'ai vécue durant ces années. .Jt. les reiiiercir et cipèrr pouvoir èire à mon tour. à Leur côté. ce qu'ils ont été pour moi.
Ainsi prend fin un tra~ailqui a mis piusieurs années à miirir. ptiisse Dieu me donner la chance de continuer à faire de la recherche le restant de mes jours (On oublie toiijours de le remercier Lui!). Disons qu'nu nombre de postes disponibles. une inter\-ention di\-ine est in&-i table pour en décrocher un! Dusse-je porii.oir un jour réalisé ie rève de redonner à nion pays iiiie partir de ce queilni 'a offert. inalgr4 1.5 ans de guerre! L'astronomie n'était-elle pas également cet appel de l'infini. ce
30th de l'élan vers les hauteun. que nous portons en nous! INTRODUCTION
Les études de la poussière ont montré qu'elle contrôle la température du SIIS en permettant la formation de molécules par l'absorption de la radiation CV qui les dissocierait. et en éjectant des électrons extraits par effet photoélectrique. Elle est également responsable de la polarisation linéaire et circulaire de la lumière qui nous provient des étoiles. Celle-ci étant due à de l'extinction dichroïque par des grains allongés et alignés par les champs magnétiques. Il y a deux observationels fondamentaux servant à ces études:
FIGURE 0.1. Courbe d'extinction interstellaire tel gubbservée pour différentes lignes de visée (Mathis 1989) 0 la courbe d'extinction intersteUaire(Figure0.1 ).
la courbe de polarisation interstellaire observée et modélisée par la loi
empirique de Serkowski (1971). & = exp [-h. ln2 (+)]*oz (Figure 0.2).
OU K = -0.10 + 1.86Xm, (A,, en pm),
FIGURE 0.2. Courbe de polarisation interstellaire. (tirée de Bastien 1991). calculée pour différentes valeurs de A,,, .
auxquelles il faut ajouter ce qui suit:
0 certaines caractéristiques spect roscopiques semblent aussi provenir des grains plutôt que du gaz interstellaire: la raie d'absorption à 10pm. ainsi que des raies infrarouges expliquées par l'existence des Hydrocarbures Polycycliques Aromatiques (PAH) (Puget et al l985)),
0 les raies interstellaires diffuses qui pourraient être expliquées aussi par les PAH (van der Zwet & Allamandola 1985).
Les étoiles, produit des condensations moléculaires des nuages interstellaires, accumulent autour d'eues des environnements provenant du materiau originel dont elles se sont formées. Ces environnements. du fait même de leur lien primaire à l'étoile naissante, vont subir des transformations et des influences reliés à la présence de cette source de rayonnement. Ils vont à leur tour influencer l'évolution de la proto-étoile. Déjà, au premier stade de l'effondrement gravitationnel. les gains de poussière vont jouer un rôle essentiel dans la détermination de la masse hale de la condensation stellaire primaire qui va se former. En effet, ils contribuent à la baiance des opacités mais dissipent également par rayonnement une partie de l'énergie gravitationnelle extirpée du nuage. Ultérieurement, en plus d'influencer considérablement l'évolution de 1'étoile, le milieu circurnsteilaire \a être essentiel à la formation possible de planètes.
0.1 La formation stellaire dans ses grandes lignes
La clé de voûte de l'étude de l'évolution des étoiles en formation est ce qui est communément appelé la distribution en énergie spectrale ( DES). Celle-ci est un spectre continu couvrant le rayonnement électromagnétique de I'W à \'IR lointain. La première classification des OSJ se base sur ce type d'observations pour construire un schéma évolutif des proteétoiles de faible masse. En effet Lada (1985) identifie 3 classes d'objets définis comme suit (voir Figure 0.3):
1. Objets dits de Classe 1, caractérisés par une DES croissante dans l'IR jusqu'aux plus grandes longueurs d'onde. Lada identifie ces objets comme des proteétoiles encore enfouies dans le nuage ou elles se sont formées.
2. Objets dits de Classe II, caractérisés par une DES montrant également un excès IR mais moins prononcé que ceux de Classe I et de pente décroissante. Les étoiles T Tauri classiques forment cette classe
3. Objets dits de Classe III, caractérisés par très peu ou pas d'excès infrarouge relativement à un spectre de corps noir. Ces étoiles sont très proches de la séquence principale et constituent la dernière phase de formation. L'étoile s'est débarassée de la majorité de la matière où elle était plongée.
Récemment, l'observation millimétrique ayant mis en évidence des objets ext ré- rnement froids, supposés en phase d'effondrement initial, André et al (1993) ont alors introduit une Classe O d'objets qui, précisement. contiendrait tous les objets représentant les tous premiers pas vers la formation d'une étoile. Un des succès de ce schéma est qu'il fut rapidement modélisé théoriquement. En effet. les travaux de Adams et al (1987) identifient cinq phases pour les étoiles de faible masse:
1. phase d'effondrement pur.
2. protoétoile. source d'un flot collimé de matière.
3. T Tauri avec disque et enveloppe de poussière. résidus de l'effondrement.
4. T Tauri avec disque seulement.
-5. post T Tauri.
Les deux premières phases sont équivalentes à la classe 1. la troisième représente- rait une transition vers les objets de Classe II. les T Tauris à disque d'accrétion. pour finir avec des étoiles épurées de leur enveloppe poussiéreuse.
Ces descriptions citent les propriétés observées du milieu circumstellaire, quant aux structures qui y apparaissent. En effet, les observations moléculaires (CO en particulier), ont tôt fait de montrer la présence de large Bots bipolaires dont la source semblait être les T Tauris ou des objets de Classe 1 (Snell et al * 1980, Canto et al 1984, Levreault 1988). Les observations optiques montrèrent également la présence de jets très bien collimés et provenant également de ces étoiles en formation. L'observation des disques, jusque très récemment, fut moins directe. Les interferomètres millimétriques, l'optique adaptative et le téléscope -'.~,, [ ,-, V37û ,Ces , ,, 1
. FIGURE 0.3. Comparaison des DES des OSJs sriivant fa classification de Laàa (1985) (incluant les Classes O) et de celles des EAEBEH telles qu'obtenues par Hillenbrand et al (1992). Adaptée de la figure de Bachilier et al (1996) et des graphiques de Hillenbrand. spatial ont permis d'imager directenent les disques d'accrétion autour d'étoiles jeunes (O' Dell & Wen 1994, Close et al 1996, Lay et al 1996 entre autres). Les premières évidences de la présence de disques d'accrétion vinrent de la volonté d'interprétation des excès IR observés. Cohen (1983) fut le premier à proposer l'existence d'une telie structure autour de HL Tau. La modélisation par Adams & Shu (1986)' Adams et al (1987) ainsi que Bertout et al (1988) vint également soutenir cette hypothèse; Il est à noter que ces disques, dits d'accrétion du fait même qu'ils amèneraient la matière du cocon initial vers l'étoile. ont des dimen- sions de l'ordre de 100 AU et sont maintenus par les effets de la gravitation et de fa rotation.
On est tenté de porter un tel schéma de formation aux étoiles de masse intermédiaire, les EAEBEHs. La première difficulté vient du fait que ces objets ne présentent pas la même homogénéité présentée par les T Tauris. De plus. malgré la similarité qualitative de plusieurs de leurs caractéristiques. ces deux types d'étoiles présentent également des différences importantes telles. par exemple. la présence d'un excès UV dans les DES des T Tauris que Leurs consoeurs plus massives ne présentent pas. Pour des revues exhaustives des deux types d'étoiles. on consultera Catala (1989) et Bertout (1989). Herbig (1994) s'intéresse aux EAEBEH. plus particulièrement aux difficultés reliés a la définition d'une classe unique de critères d'appartenance au groupe.
Le premier essai de classification des étoiles de Herbig sépare 47 ob&s dont on a les DES en trois groupes (Hillenbrand et al 1992):
1. Le Groupe I compte tous les objets présentant des excès IR à des longueurs d'onde X 2 2pn, avec une pente en loi de puissance, Afi x A-:. indiquant un disque d'accrétion (Bertout et al 1988); 2. Le Groupe II contient tous les objets dont l'excès IR est plat oii croissant:
3. Le Groupe III est celui des étoiles présentant peu ou pas d'excès IR.
La séquence évolutive inférée va du Groupe II au Groupe III en passsant par le Groupe 1. Les différences sont modelées par des structures circumstellaires dépendant du groupe- Le Groupe I présenterait un disque d'accrétion avec peu ou pas d'enveloppe plus large. Par contre le Groupe II, qui le précéderait dans la chaine d'évolution, nécessite la présence d'une enveloppe importante. en plus du disque d'accrétion pour reproduire correctement les profils spectraux continus observés. Les membres du groupe III seraient entourés d'un disque tenu. relique du milieu originellement dense de l'environnement circurnstellaire. La figure 0.3 montre. côte à côte. des spectres typiques d'étoiles des trois groupes ainsi que des représentations schématiques des DES des trois classes de Lada et de la classe 0: l.376 appartient au Groupe II. LkHo 234 au groupe 1 et BD +65" 16:37 au groupe III. Le parallèle semble assez intéressant ne serait-ce du fait que. pour les objets de Classe 1. les étoiles sont enfouies et indétectables dans le risible. Les objets de masse intermédiaire appartenant au groupe équivalent. !e groupe II. sont observables aux longueurs d'ondes optique. Il reste à démontrer bien sûr que les pics observés dans les images de nébulosités associées à ces étoiles sont belle et bien les étoiles elle-même. En effet, les travaux de Leinert et ses collaborateurs (1991, 1993 entre autres) montrent que, pour plusieurs objets du groupe II. ces "pics stellaires* sont en fait de la lumière diffuse où se retrouvent les mêmes structures observées à grande échelle.
Ces résultats sembleraient concluants ne serait-ce du fait que les solutions trouvées par Hiiienbrand et al ne sont pas uniques. En effet Hartmann et al (1993) ont démontré qu'un disque n'était pas une nécessi té fondamentale pour l'interprétation de I'dure des DES des membres du Groupe II. Ils arrivent en effet à reproduire le spectre continu en modifiant les propriétés de l'enveloppe sphérique. Un résultat que confirme également les travaux de Natta et al (1992. 1993) sur un certain nombre d'étoiles Ae/Be de Herbig, en particulier LkHa 198. Là où des certitudes forment nos connaissances des T Tauri, des inconnues sont ce qui fait les EAEBEH. L'abscence de disque autour d'une T Tauri est vu comme l'exception qui conftrme la règle tandis que pour les proto-étoiles de masse intermédiaire, on en est encore à chercher la réponse définitive à cette question. On consultera avec profit les revues de Bertout (1989), pour les TTS, et de Catala (1989) pour les EABEH ainsi que l'analyse des critères de sélection des EAEBEH faite par Herbig (1994).
Un autre problème important des DES vient du fait qu'elles combinent de l'information obtenue par des moyens très différents et à des résolutions variant de la seconde d'arc dans le visible à quelques dizaines de seconde d'arc dans I'IRL. Les techniques d'observation récentes. interférométrie speckle dans le PIR. l'optique adaptative. l'imagerie dans l'infrarouge moyen. ainsi que la cartographie dans le millimétrique d'objets jeunes viennent montrer les dangers d'une telle approche vu que beaucoup de sources se trouvent avoir des cornpanions IR ou millimétriques (voir. entre autres. Leinert et al 1994. Lagage et al 1993. LVeintraub et al 1994 ainsi que ce travail). Le chapitre 1. en particulier, montre une telle situation.
L'importance des structures circumstellaires est donc primordiale à la com- préhension de la physique de la formation stellaire. La connaissance précise des structures présentes et leur caractérisation sont essentielles au déchiffrage des processus en jeu dans la naissance d'une étoile. 0.2 Les structures du milieu circumsteilaire
Les objets de faible masse, comme nous l'avons vu précédemment. présentent des caractéristiques de groupe homogène. Les structures circumstellaires ne dé- rogent pas à la règle. La plupart des T Tauri ou des objets de Classe 1 sont des sources de flots bipolaires moléculaires, créant une cavité évidée aux deux pôles de l'étoile. Il est également cIair que ces objets possèdent des disques d'accrétion dont la taille est de l'ordre de 100 AU (Bertout et al 1988. Shu et al 1987). la modélisation de leurs DES étant bien reproduites par de telless structures. ainsi que l'interprétation d'observations spectrales. On croit que ces disques sont géométriquement minces et optiquement épais, du moins aux longueurs d'onde du visible a l'IR. Certaines observations montrent également la présence de disque ou de structure applatie, à plus grande échelle. L'archétype de telles structures vient de l'observation millimétrique d'un disque de 2000 AC' de diamètre autour de HL Tau par Sargent S. Beckwith ( 1987. 1991). Les cartes de polarisation. que nous aborderons en détail un peu plus loin dans le texte. montrent également la présence de disques de grandes dimensions. Les travaux du groupe de formation stellaire de l'université de Montréal ont permis de montrer la formation de telles structures lors de l'effondrement de nuages moléculaires allongés (Bonne11 et al 1991). Récemment, des calculs théoriques de Galli & Shu (1993a.b) ont permis également d'identifier un mécanisme qui reproduirait de tels disques: c'est ce qu'ils nomment pseud+disques. Ces structures ne sont évidemment pas supportées par la rotation mais découlent, dans l'optique de Galli & Shu, d'un support magnétique offert à la matière qui s'effondre vers l'étoile via le disque d'accrétion, Des ob- servations millimétriques récentes du disque de HL Tau (Cabrit et al 1996) ont montré la présence de mouvement d'effondrement mélangés à des pertes de masse - montrant que le disque n'est effectivement pas en rotation pure autour de l'étoile mais s'apparenterait aux pseuddisques de Galli & Shu. Le futur dira si une étude plus approfondie du systême confirmera le modèle de Gdi & Shu. En plus des flots bipolaires observés, des jets optiques sont associés aux étoiles en formation, qu'elles soient de faible masse (e-g, hIundt et al 1990 pour HL Tau) ou de masse intermédiaire (e.g., Corcoran et al 1993 pour LkHo 198). Les études théoriques ont produit un mécanisme plausible de formation des jets (Pudritz 1993 et références incluses, voir également Bachiller 1996). Celui-ci nécessite la présence d'un disque dont la perte de matière. collimée par les lignes du champ magnétique stellaire, produit un jet de matière montrant les même caractéristiques que les jets observés. Ces derniers seraient quasi perpendiculaires au plan du disque. Les quelques objets pour lesquels nous avons une mesure solide de L'orientation du disque et du jet montrent effectivement une telle géométrie (voir Bastien S- Ménard 1990). La figure 0.1 représente schématiquement les structures susment ionées.
11 est à noter que. depuis leur identification. les objets de Classe O sont venus bouleverser quelque peu notre compréhension des mécanismes accompagnant la formation d'une étoile. 11 apparait de plus en plus évident que les phénomènes d'expulsion de matière sont quasi contemporains à l'effondrement gravitationnel initial. La revue de Bachiller (1996) en fournit une vue complète et détaillée. Ceci évidemment montre que les phénomènes de projection et d'accrétion de matière sont fondamentaux dans la formation steliaire. Un telle hypothèse a mené à la formulation par Adams k Fattuzo (1996) d'une théorie qui permettrait d'obtenir la fonction de masse initiale à partir de l'observation des coeurs de nuages moléculaires. Ces derniers étant les premiers signes du déclenchement d'une production d'étoiles.
L'étoile en formation et son environnement circumstellaire forment donc un tout. La compréhension profonde de son évolution nécessite ainsi le déchiffrage des différents mécanismes d'interaction ainsi que l'étude des différentes composantes. IFloc Bipolnirt
FIGURE 0.4. Vue schématique des structures majeures du MCS
11 est donc important de connaître en détail la constitution physique du MCS: Les densités. compositions. dynamique... Dans cette optique. la poussière. composante essentielle du MIS. va donc contribuer d'une facon importante à cette environne- ment. En effet. elle détermine l'état de polarisation de la lumière et. grandement. l'émission électromagnétique dans l'infrarouge. Sa contribution croit avec la lon- gueur d'onde, quand celle de l'étoile diminue. Il est à noter que I'intérêt porté aux environnements des étoiles en formation déborde sur la problématique de la formation de système planétaire. Quand l'environnement solaire nous offre un système planétaire déjà formé, donc un problème aux conditions limites. l'obser- ~ztiondes environnements d'étoiles jeunes offre l'avantage de sonder directement la nébuleuse prirnordide. Évidemment, des problèmes d'échelle de comparaison viennent se poser dans la mesure où, à la distance de la pépinière stellaire du Taureau, l'une des plus proches, une seconde d'arc représente 140 AU; le disque est donc non résolu. Sans oublier le fait que l'observation in situ est impossible. Les deux voies suivies pour la recherche faisant l'objet de cet ouvrage décou- lent de ce qui précède. Des observations millimétriques et polarimétriques - de l'imagerie polarimétrique - feront l'essentiel des données obtenues. Quel genre d'information peut-on obtenir de ces techniques?
0.3 Le millimétrique et submillimétrique
L'intérêt porté aux observations dans le millimétrique et le subrnillimétrique vient de ce qui est communément appelé La quête du Graai de la formation stellaire. la découverte d'une protoétoile par excellence. une étoile qui n'en serait pas encore tout à fait une prise au moment où le nuage s'effondrerait. De tels objets sont évidemment très froids, il n'y a encore ni réaction nucléaire. ni coeur contracté suffisamment dense pour avoir des températures assez élevées. On retrouve ainsi de nombreux papiers où les auteurs ont cherché (sans succès) une source très froide
et n'ayant pas encore de jets ou de flots bipolaires (e.g.. Zinnecker et al 1992). Les tous premiers stades de l'évolution sont par excellence des stades aisément détectables dans le millimétrique de par ce fait même. La radiation thermique de ces objets atteint son maximum à des longueurs d'ondes de I'IRL. La découverte des objets de Classe O procède de cette logique (André et al 1993). L'ouverture des fenêtres d'observation dans 1'IRL a permis d'identifier des objets pratiquement indétectables dans le visible ou le PIR. D'ailleurs hs critères d'appartenance à cette classe d'objets font appel aux flux dans le millimétrique (André et al 1993). La Figure 0.3 illustre bien ce propos. Ces observations ont permis de réaliser que l'éjection de matière apparait presque en même temps que l'accrétion. D'ailleurs les objets de Classe O sont les moteurs de flots bipolaires (Bachiller 1996 et références incluses).
L'émission millimétrique astronomique est essentiellement d'origine molécu- laire (e.g, CO, CS, HCO+) pour les raies d'émission et provient de la poussière pour ce qui concerne ie continu thermique. C'est précisément ce qui fait l'intérêt de telles observations pour les environnements circumstellaires. En effet. on peut affirmer que l'émission d'OSJ dans le millimétrique est uniquement due aux grains de poussière ce qui permet donc de les étudier sans se soucier d'autres con- tributions au flux électromagnétique dans cette région du spectre. Yous nous intéresserons donc a l'information que porte l'observation du continu thermique quant aux environnements des étoiles en formation.
Hildebrand (1983) montre la possibilité du calcul des masses de nuages denses à partir des observations millimétriques. En effet. vu l'opacité faible des grains interstellaire à ces longueurs d'ondes, aux densités considérées, les milieux sont optiquernent minces. De plus, il est possible de conndtre la loi d'émission des grains 6, à partir de la photométrique millimétrique multi fréquence. L'émissivité des grains dans ce régime de fréquences s'écrit
Aux températures des coeurs moléculaires. on peut raisonnablement supposer que le spectre d'émission suit la limite de Rayleigh-Jeans, B,(T) sc u2 oii BJT)est la fonction de Planck. Ainsi le flux provenant du milieu observé aurait la forme F, a v2+< Il est donc possible de tirer de l'information sur les propriétés optiques des grains à partir de la mesure directe des flux millimétriques et submillimétriques. Connaissant la température, la dimension de la source, l'index J d'émissivi té des grains et supposant un milieu optiquement mince, il devient possible de déterminer la masse d'un nuage ou d'un coeur moléculaire. Cette approche est appliquée également aux OSJs (e.g., Weintraub et a1 1989, Walker et al 1990, Zinnecker et al 1992). En poussant plus loin les modèles. il est également possible d'extraire de IÏnformation sur les disques et les enveloppes du MCS (e-g., Adams et a1
1990, Beckwith et ai 1990, Ohashi et al 1996). Par contre, ces résultats souffrent de différents problèmes dont le fait qu'ils exigent plusieurs paramètres libres. Les différents auteurs vont généralement considérés un système étoile-disque et
considérer une température en loi de puissance dans le disque. T = To ($q ainsi qu'une loi de puissance en densité. Les rayons internes et externes du disque sont également des paramètres libres du modèle (e.g., Hillenbrand et al 1992). Il faut également noter que les environnements circumsteilaires présentent des structures et des caractéristiques fondamentalement différentes des coeurs me léculaires. La structure en densité en effet pourrait bien mener à des régions optiquement épaisses même aux radiations du submillimétrique. Beckwi th et al ( 1990) explorent d'ailleurs cette possibilité et élaborent sur les effets possibles d'une telle éventualité. II est à noter également que les milieux denses du disque favoriseraient la coagulation des grains et leur croissance bien au delà des tailles du 111s. D'ailleurs. le changement de taille dans le .CI IS en fonction de la densité du milieu est déjà bien observé (Mathis 1989 et références inclusesj.
Les observations millimétriques sont rarement considérées indépendamment de I'ensemble des obsenat ions p hotomét riques des OSJs. Ainsi que présenté plus haut, les modèles tentent de reproduire les DES dans leur ensemble (e-g.. Hillenbrand et al 1992). Les interprétations des observations restent quand même multiples. Pour les EAEBEH, par exemple. les étoiles du Groupe II sont modélisées
en considérant un disque et une enveloppe étendues par Hillenbrand et a1 (1992) mais uniquement avec une enveloppe par Hartmann et al (1993). La présence de compagnons invisibles posent également un problème à l'interprétation adéquate des DES. Le faisceau du Kuiper Airborne Obserntory à 100 prn est de 37". celui du JCMT à 1.1 mm de 18.5". Les observations visibles atteignent typiquement des résolutions de 1" à 2". Or les observations IR récentes ont montré la présence de - compagnons enfouis près d'EAEBEH qui sembleraient contribuer significativement au flux dans 1'IRL (e.g., Lagage et al 1993 pour V633 Cas et le Chapitre 1 de cet ouvrage, Weintraub et al 1994 et 1996, ainsi que Cabrit et al 1997 pour LkHa 234). La cartographie des environnements détoiles jeunes est donc importante, tel que va le montrer le premier chapitre.
0.4 L'imagerie polar imét rique
0.4.1 Une introduction
La figure 0.5 montre un exemple typique de carte de polarisation d'un OSJ. en l'occurence HL Tau, une T Tauri. Les vecteurs représentent la direction et l'intensité du champ électrique. linéairement polarisé. Les vecteurs montrent un certain nombre de profils clairement identifiés. Près du pic. on remarque une série de vecteurs parallèles délimitée aux deux extrémités par deux points de polarisa- tion nulle. en s'éloignant du centre. le profil devient centrosymétrique: d'elliptique proche du centre. il devient circulaire plus loin. Ce type de profil n'est pas l'apanage des OSJ dans la mesure où tout objet présentant des caractéristiques circurnstellaires semblables montre le même genre de profil de polarisation (voir Kastner Sr Weintraub 1995. pour un exemple de carte d'objets plus évolués). Le profil parallèle n'est pas toujours observé et certains objets présentent des profils particuliers. Bastien Sc Ménard ( 1990) ont présenté une compilation de toutes les cartes de polarisation d'OS J publiées à cette date, en décrivant pour chacune les caractéristiques observées.
0.4.2 L'interprétation et les modèles
Les premiers travaux d'imagerie polarimétrique datent du début des années 1980 (e.g.. Draper et al 1985). Cette technique fut rendue possible grâce à l'avè- nement des CCD. La fin des années 1980 vit le début de tentatives sérieuses de compréhension de telles cartes. La confrontation des premières observations polarimétriques d'OS J ( polarimétrie d'ouverture), des cartes de polarisation et 1 1O 5 O -5 -10 R A. Offset (arcsec)
FIGLTRE0.5. Carte de Polarisation dans la bande 1 de HL Tau. retrouvée plus loin dans cet ouvrage. Les vecteurs représentent l'intensité et l'orientation du champ électrique polarisé. Le profi1 de vecteurs est superposé aux contours d'intensité de la nébuIeuse.
des modèles d'extinction dichroique et de diffusion simple ( Bast ien 1987) conduit aux résultat qu'aucun des deux modèles ne pourrait expliquer toutes les observa- tions. Les premiers travaux d'interprétation des cartes furent réalisés à l'Université de Montréal par Bastien & Ménaïd (1988, ci après BM) et Ménard (1989). Il est à noter qu'avec les premières cartes obtenues les différents observateurs - empruntèrent initialement au MIS l'interprétation des profils parallèles de vecteur. En effet, cew-ci considéraient l'extinction dichroique par des gains alignés par un champ magnétique dans le disque comme le mécanisme fondamental de polari- sation. Ce que le modèle BM vint modifier est précisément cette interprétation. Le modèle considère uniquement la diffusion par des grains de poussière distribués suivant une géométrie identique à celle représentée à la figure 0.1. De la difision simple dans les lobes bipolaires et une double diffusion dans le disque permet aux auteurs d'expliquer les observations, du moins dans leur grande ligne. Celà leur permit de conclure que les cartes de polarisation des étoiles jeunes peuvent être expliquées avec la diffusion multiple et une géométrie appropriée.
L'étape suivante, développée par Wénard (1989)' consista a considérer des distributions en densité plus réalistes et à développer un code Monte Carlo pour traiter la diffusion multiple de façon rigoureuse. Le profil parallèle est alors dû à de la diffusion multiple dans le disque et le profil centrosymétrique à de la diffusion simple dans l'enveloppe étendue et optiquement mince. Les points nuls. dans cette optique, sont alors interprétés comme des points de transition entre la région optiquement épaisse et celle optiquement mince du disque. Ils correspondent à une profondeur optique dans le disque r 1. Des calcals de transfert radiatif de la lumière polarisée. aux mêmes fins que celles précédemment citées. et par des techniques Monte Carlo, ont été réalisés par Whitney et al (1992). Sfénard (1993). Ménard et al (1996) et Fisher et al ( 1994. 1996). Les modèles les plus récents, tels ceux de Fisher ( 1996)? de Ménard et al (1996) ou ceux de LF9hitney et a1 (1997) adoptent des profils de densité empruntés aux calculs hydrodynamique de formation stellaire.
Toujours dans le but de proposer une interprétation des cartes de polarisa- tion, Gledhill(1991) opte pour une approche différente. il considère une nébuleuse cylindrique autour d'une source de radiation polarisée. Aucun mécanisme ou géométrie ne sont proposés pour expliquer la polarisation du rayonnement de la boite noire centrale. Le modele a l'avantage de ne recourir ensuite qu'à de la diffusion simple. Ce qui permet à l'auteur d'étudier les effets d'inclinaison de la direction privilégiée d'émission de la source relativement à l'axe de symétrie de la nébuleuse. A la lumière de nos connaissances actuelles. ceci équivaut à l'angle que font l'axe de symétrie du disque et celui du flot bipolaire.
0.4.3 L'utilité des cartes de polarisation
Ces travaux vinrent établir finalement la diffusion comme le mécanisme expliquant la polarisation des OSJ. Bastien k Ménard (1990) effectuent une comparaison des deux mécanismes proposés. l'extinction dichroique et la diffusion multiple. montrant ainsi clairement les avantages d'une interprétation des proprié- tés polarimétriques des OSJ par la diffusion. De plus. les calculs de BSI. réalisés pour différentes inclinaisons du disque par rapport à la ligne de visée. vinrent établir un lien entre cet angle et la présence. et l'aspect du profil parallèle de vecteurs. Ceci permet à Bastien S; Slénard (1990) de calculer les angles d'in- cliriaison des disques de plusieurs objets stellaires par la simple comparaison des cartes de polarisation observées et théoriques. Leurs résultats concordent très bien avec ceux obtenus par d'autres méthodes. Par la suite. Whitney et al ( 1992) caractérisent les structures circumstellaires de R Mon. une EAEBEH. par la modélisation de sa carte de polarisation. une approche également poursuivie par Whitney et al (1997) pour interpréter les cartes de polarisation infrarouges d'objets enfouis du Taureau.
Historiquement, une des utilisations des cartes de polarisation eut pour but de montrer l'association physique d'une nébuleuse de réflexion à l'étoile centrale - (e-g., Aspin et al 1985, pour LkHa 233) ou de localiser la source d'illumination d'une nébuleuse (e.g, Draper et al 1985, pour RN0 138, et plus récemment Weintraub ef al 1994 pour LkHa 234). D'autres utilisations sont ceUes proposées précédemment, où les modèles sont comparés à l'observation afin de tirer de FIGURE 0.6. Cartes de polarisation théoriques calculées pour différentes distributions théoriques de matière autour des OSJ. Figure extraite de Fisher et a1 (1996). l'information sur le milieu circumstellaire. Les modèles les plus récents montrent d'ailleurs le potentiel des cartes de polarisation. La Figure 0.6 est extraite de Fisher et al (1996);elle montre des cartes de polarisation théoriques pour différents modèles de distribution circumstellaire de matière. Pour l'interprétation des abré- viations utilisés (TO, LA, RB ...) on se réferera à l'article de Fisher et al. On voit clairement que les cartes de polarisation offient un moyen puissant de déterminer les particularités de la distribution en densite du MCS! Les profils de vecteurs ainsi que leurs modules montrent des différences notables d'une distribution à l'autre. Malheureusement, les auteurs n'offrent, à part les cartes publiées. aucune grandeur quantifiable et pouvant être également évaluée sur une carte observat ionnelle qui peremettrait une comparaison quantitative de l'observation et de la théorie.
0.5 Les objectifs
Tous les travaux réalisés à ce jour se sont intéréssés à regarder des cartes de polarisation linéaire à une seule longueur d'onde. Aucune étude multi longueurs d'ondes n'est disponible dans la littérature. Les codes Monte Carlo ont jusqu'à ce jour été utilisé pour calculer des caxtes de polarisation employant des modèles de distribution de matière variés et de plus en plus élaboré. tels ceux de Galli S; Shu (1993a, b).
Le potentiel d'une approche où des cartes à plusieurs longueurs d'onde sont comparées est illustré par la prédiction du modèle BM. Le disque dit polarimé- trique est celui mesuré par la distance séparant les 2 points nuls. Si ces derniers - marquent la ligne de visée où la profondeur optique r est égale à un, on devrait s'attendre à observer leur déplacement en fonction de la longueur d'onde. En effet, en allant vers l'IR, le milieu devient de plus en plus transparent à la radiation. Ceci devrait amener, suivant la prédiction de BM, une diminution de la dimension du disque polarimétrique. Les chapitres trois et quatre, deux publications scientifiques à soumettre à Astrophysical Journal, sont consacrés à l'observation et à l'étude de cet effet et ébauchent également une première approche permettant l'extraction de la densité volumique de l'équateur du disque à partir de la mesure du déplacement des points nuls.
Le chapitre deux. un article préparé pour être soumis à l?..lstrophysical Journal, aborde Ia question de la détermination de la position de l'étoile à partir d'une carte de polarisation. En effet. cette approche avait déjà été utilisée par le passé mais uniquement à travers un examen visuel de la carte de polarisation. Un simple examen peut en effet fournir des éléments de réponse intéressants daos des cas où 1'objet stellaire est suffisament enfoui pour échapper à l'observation directe (voir Darper et al 1985 et Kastner et al 1994), mais devient plus problématique dans le cas où l'observation semble montrer la présence d'un pic stellaire. L'obser~ation à très haute résolution, telle l'optique adaptative ou l'interférométrie speckle. a mont ré que ces pics stellaires étaient en fait résolus et constitués principalement de lumière diffuse. HL Tau. observée par Stappelfeldt et 01 (1994) avec HST. est un exemple typique d'étoile T Tauri montrant ce phénomène. LkHa 198. observée en interférométrie speckle par Leinert et al (1991) et, plus récemment par Koresko et al (1990, est un cas d' EAEBEH possédant ces caractéristiques. Dans ces situations. la carte polarimétrique permettrait de déterminer de façon précise la position réelle de l'étoile. La méthode quantitative utilisée est élaborée au chapitre deux et discutée d'une manière plus élaborée au chapitre quatre.
L'observation millimétrique portera sur deux axes, la cartographie du MCS d'étoiles jeunes ainsi que la photométrie multifréquence. L'objectif essentiel vise à tenter éventuellement de comparer les résultats obtenus par la polarimétrie et le millimétrique. La cartographie servira à déterminer l'étendue et, lorsque possible, la morphologie de l'environnement de 1'OSJ observée. Combinée à la photométrie. il est dors possible de conndtre les masses des enveloppes circurnstellaires ainsi que certaines des propriétés optiques des grains de poussière. Le premier chapitre. un article soumis à 1'Astrophysicul Journal, présentera les observations de V633 Cas et de V628 Cas, deux EAEBEH. V633 Cas fait également partie de la liste d'étoiles faisant I'ob jet de l'étude du chapitre 4. Des observations millimétriques sont incluses dans les chapitres 2 et 3. Adams. F.C., Emerson, J.P. & Fuller. G. 1990. ApJ. 357. 606
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Observations miilimétriques de V633 Cas et V628 Cas
Article scientifique soumis à I'Astrophysical Journal V633 Cassiopeiae (LkHa 198) and V628 Cassiopeiae (MWC 1080) Revisited!
Roger Hajar', Pierre Bastien' Star Formation Group and Département de Physique, Université de Montréal and Observatoire du Mont Mégantic C.P. 6128, Succ. "Centre-Ville", Mont réal, Québec, H3C 157. Canada.
Received 13 September 1996; Accepted
lGuest Observer, James Clerk Maxwell Telescope ABSTRACT
We present the results of mm observations carried out at the James Clerk Maxweil teleseope (JCMT) on August 1993 and August 1994 of two Herbig Ae/Be stars, namely V633 Cas and V628 Cas. We follow up on the reported discovery of a probable protostellar object 20" North -West of V633 Cas by Hajjar k Bastien (1994), also reported later by Sandeil & Weintraub (1994). It is the main source of emission in the region at mm wavelengths, and probably the major contributor in the submm up to 100 Pm. We estimate a total mass of 0.12 .Mi, with a dust ernissivity index of 0.35 f 0.21. As seen in maps from Nakano et al (1990), V633 Cas C is at the center of a 13C0 cavity and at the position of observed molecular concentrations. Since it also Falls at the center of a flattened structure seen in HCO'. it is most probably the source of the large scale CO outflorv (Levreault 1988). From the dimension of the CO cavity. tve estimate an age of 28 x 103 years for the object. Based on a first estimate of its SED. ive find a ratio of
Lbol/ L tJmm > I x 104. Based on al1 these facts. tve conclude that \-633 Cas C is not a Class O but an extreme Class I object. The dust emission from MWC 1080 is extended and flat with some leatures correlated with the optical ones. It is not evident that the grains we are sampling are directly related to the stellar environment,
Infrared: stars - Stars: formation - Stars: pre-main sequence - Circurnstellar matter. 1.1 Introduction
Herbig Ae/Be Stars (HAEBES), first identified by Herbig ( 1960). are now commonly accepted to be intemediate mass pre-main sequence stars. They are associated with nebulosity, and are surrounded by circumstellar materid resulting from their formation process. Catala (1989) gives an extensive review of their properties.
The main drive behind most recent studies involves understanding the evo- lution of intermediate mas pre-main sequence stars through the classification of HAEBES and their inclusion in a larger star formation scenario for 2 to 10 Mt3 stars. Thé et al (1994). bwed on accumulated spectroscopic data on Young Stellar Objects in generd - especially T Tauri stars - and HAEBES in particular. proposed a refined set of selection criteria and a classification based on spectral characteristics. similar to the one adopted for T Tauri stars. They try to interpret this classification scheme in terrns of an evolution of the stars towards the main sequence. Herbig ( 1994) reviews possible classifications of this group he established. Both discussions are based on opt ical. mainly spect roscopic. data. Another attempt. following a completely different path has been made by Hillenbrand et al (1992). It relies on the spectral energy distribution (SED) and on its slope in the mid to Far infrared part of the spectrum.
The millimeter and submillimeter windows, now accessible to ground-based observations, have brought their own sets of data to this study. The main concerns go towasds understanding the circumstellar environment of HAEBES since emis- sion at these wavelengths is mainly due to dust gains. The kind of structure needed to explain the observed SEDs is still a hotly debated subject (Hillenbrand et al 1992, Hartmann et al 1993). Models usually assume an accretion disk within a larger envelope (Natta et al 1993, Hillenbrand et al 1992, Francesco et al 1994) for example, with types of dust grains i>r a gap in the disk. The types of dust grains is in itself controversial since photometric uncertainties bring up even larger uncertainties in the deduced dust emissivity, d, dlowing for a spectrum of possible models to reproduce the SEDs. The first classification scheme based on far-infrared excesses is described by Hillenbrand et al (1992). They classify stars in three groups based on the SED dope longward of 3prn and an evolutionary sequence is inferred from the circumsteUar properties of these groups. Group II. with flat or rising spectra, represents the earliest evolutionary stage: these objects are modeled with an active accretion disk and an envelope. Recent observations question sorne of the conclusions of their study relating to individual stars (e-g.. ?latta et al 1992. Hartmann et al 1993). The discovery of IR cornpanions also disturbs the picture we have of these environments (e.g.. V633 Cas-IR . Lagage et al 1993. Corcoran et al 1995).
The wealth of data collected in different spectral cvindows points t o a some- what evident fact. For a better understanding of the evolution and formation of HAEBES and higher mass objects rve clearly have to gather observations obtained at different rvavelengths. The infrared is particularly important since more massive objects tend to spend a larger fraction of their pre-main sequence phases embedded in their cocoon. Imaging in the IR. near and far. is rhen crucial to the understanding of the processes involved in the evolution of HAEBES and the modeling of their circumstellar environments.
In this paper, we present maps of the environments of two such stars. 1-633 Cas and V628 Cas. In the next section, we present our observations and calibration procedure then, in two separate sections we discuss our results for each object. Section 5 is devoted to the review of our main conclusions and their possible implications. 1.2 Observations
The observations presented in this paper were obtained at the JCMT in August 1993 and August 1994. The JCMT has a 15m dish. Continuum data and rnapping were obtained with the UKT 14 bolometer (Duncan et al 1990). Both maps, V633 Cas and V628 Cas, are point-by-point, very deep maps. at a
wavelength of 1.1 mm. Photometry was performed at 0.8. 1.1 and 1.3 mm on detected features in both maps in August 1994 and at 0.15 mm in December 1995. Table 1.1 gives the Log of the observations.
1.2.1 Calibration
Due to the absence of planets during Our observations except for Cranus at the beginning of the shifts in 1993 and llars at the end of those in 1994. ive observed rnostly secondary calibrators. namely W3(OH) and SGC 793Y IRS1. We also made use of the 225 GHz zenith optical depth measured at the Caltech Submillimeter Observatory with a tipping radiometer (rcso ).
.At 1.1 and 1.3 mm we considered the optical depth at each moment to be
proportional to rcso as rai = a x TCSO Since these passbands are relatively close to 225 GHz, this is a fairly good approximation. .\ssuming the standard relation between gain, optical depth, flux and measured CKT11 voltage (see the JCMT Guide for the Prospective User, also Stevens b Robson 1994), we have:
where A is the airuiass, G the gain, V the measured voltage, and F the flux above the atmosphere. A linear regression through calibration data for ail nights of each run allowed us to calculate a and the Gain value. Figure l.la shows the best fit for the 1993 nights and Figure Llb for the 1994 nights.
This proved to be better than going through calibration data the standard way since sky conditions were variable through the nights. Our dues for a and the Gain as given in Table 1.2 are compatible rith advertised values in the JCMT User's Guide. The use of all the calibration data for each run allows one to minimize fitting errors.
Calibration at 0.45 mm was carried out in the same way. At 0.8 mm we used the empirical relation derived by Stevens & Robson (1994)
1.2.2 Sensitivities
Figures 1.2 and 1.3 show the error distributions for both maps. Both are plotted for each run separately. In 1993, we mostly measured the inner Y1 points of the C628 Cas map whereas the 1994 run rvas aimed at extending both maps in the directions where we still detected a signal. The 1993 error distribution of i-628 Cas is double-peaked and shows a very large profile. It cornes from the fact that for the first part of the map. the opacity of the atmosphere decreased significantly within a few hours. We lowered the integration tirne for the second part of the map. Cnfortunately. this affected somewhat the sensit ivity reached t hus producing a second peak in the distribution. Stable skies during the 1991 run produced the nicely peaked distribution for this run. These facts should be kept in mind rvhen interpreting the data. The V633 Cas data is much more homogeneous as seen from the histograms in Figure 1.2. Figures 1.1 and 1.5 show contour plots of the maps, the lowest level is at 34. the chosen u level being the peak value of the error distributions, 40 mJy per beam for the V633 Cas map and 30 mJy per beam for V628 Cas. For the latter, we chose the peak of the 1994 error distribution to emphasize the outer part; it should dways be kept in mind that the sensitivity in the inner part is more around 50 to 60 mJy per beam. 1.3 V633 Cas
V633 Cas, better known as LkHa 198, is associated with V376 Cas which is known to have the largest opticd linear polarization measured arnong pre-main sequence stars, =21-23% (Bastien et al 1989, Asselin et al 1996). It has been associated with a bipolar outflow mapped by Canto et al (1984) and Levreault (1988). Levreault first presumed that V633 Cas is the driver. Recent observations supported the fact that the driving source is closer to V633 Cas than to V376 Cas.
Fuente's (1992) molecuiar observations as well as those of Nakano et al (1990) showing dense mediurns around the star go in the same direction. Previously mentioned molecular line observations assumed that the position of V633 Cas is the one given by the Herbig & Rao catalogue ( 1988) which happened to be offsetted by 17" to the north of the actual position (Lagage et al 1993). Furthermore the discovery of optical outflows supports even more the same result (Corcoran et al
1995). The receot observations of an infrared cornpanion to 1-633 Cas ( iagage et al 1993. Corcoran et al 199#5)reopen the debate since both groups argue in favor of it being the source of the outflow.
The disk itself aroiind 1'633 Cas is still a hotly debated subject. Piirola et al ( 1992) report the observation of a parallel vector pattern in the innermost part of their 1 band polarimetric image, under excellent seeing conditions (0.1" ). indicative of the presence of a disk but it extends only to a 0.5" diameter. .A jet is probably associated with V633 Cas (HH 164, Corcoran et al). This can be regarded as indirect evidence for the presence of a disk. Another jet is also associated with V633 Cas B (HH 161). Both have position angles compatible with the large scde molecular outflow but Corcoran et al favor V633 Cas B as its source because of its "axial" position relative to the optical bubble seen as the signature of an outflow cavity. Aperture polarimetry is also inconclusive and is a clear indicator of the complexity of this region (se Asselin et al 1996 for a thorough review and discussion). Infrared observations, from near IR to far IR, show an infrared excess in the SED of V633 Cas. HiIIenbrand et al (1992) consider it to be a Group II abject. They mode1 this class of systems with a central star or starfdisk system smounded by a large, moderately opaque. dust envelope. They conclude that Group II objects are more embedded and obscured than other groups and thus represent an earlier phase of the BAEBES phenornenon. Natta et al ( 1992). using KA0 observations at 50 pm and 100 Fm, through the evaluation of the envelope dimensions and a detailed modeling of the SED propose a spherically symmet ric dust distribution with a central star/disk component to fil1 the mid IR excess: but they also argue that the presence of PAH and very small grains would relax the need for a disk. Scans at these wavelengths give a 100 Fm size of 33". However al1 authors assumed that measured fluxes are from V633 Cas only: for IRAS data it is considered that the emitters are V633 Cas and V3T6 Cas. The presence of the recently discovered IR cornpanion reshuffles the cards in this area of the sky. Based on the 10 pm silicium optical depth measured and on an assumed resemblance with IRAS 04365+2935 (Beichman et al 1986. .\lors et al l9Y;). Lagage et al note that V633 Cas B is presumably the main ernitter at 100 pm thus relaxing much of the constraints irnposed upon the CS medium of V633 Cas. In our analysis. we will make use of their results and assurnptions to try to understand the effects of Our findings on the current picture of this object and its neighborhood. as rell as possible implications for the understanding of HAEBES.
The 1.1 mm rnap (Figure 1.4) clearly shows that we are deaiing with extended emission resolved with the JCMT beam. The emission peaks at = 20" to the NW of the opticd position of V633 Cas. This discovery was first reported by Hajjar & Bastien (1994) and later by Sandeli & Weintraub (1994), hereafter SaW. SaW mapped the source at 800 Pm, and obtained a I2CO map of the area. We will compare our results to theirs later in the paper. On a large scale, the source seems to be extended in the EW, and in the NESW directions when looking at the inner isophotes. Assuming a gaussian profde in two directions for the source and a gaussian beam with HPBW of 18.5", we find a very large source with a deconvolved extent of 24" x 16". Natta et a1 (1992), based on KA0 scans at 100 prn and the same assurnptions we made. find a source size of zz 33". which is quite comparable to ours. This leads us to assume that they might have observed the new source instead of V633 Cas. dthough this would lead to some constraints regarding the emitters. SaW found a source size of 8" x 6" at 0.8 mm. Both sizes are somewhat hard to reconcile although we would expect to find a larger size with increasing A. In this case. this is a large difference that needs to be confirmed wi th higher resolution mapping and comparable sensi tivi ties at both wavelengths. Hereafter, ive narne the new source V633 Cas C.
1.3.1.1 At the center of a disk?
The positional argument is very important in comparing previous observa- tions with ours. One of the striking results comes from comparing with molecular maps obtained by Nakano et al (1990) (Fig. 1.6). Their HCO+ map shows an extended structure in the NE-SW direction with two condensations on each side of the structure. Its position angle is almost perpendicular to the outflow direction detected by Canto et al (1984) and Levreault (1988), and its "center" coincides with the mm peak. This could be the signature of a large scale molecular disk around our deeply embedded object. Furthermore. the HCO+ velocity map shown by the same authors could be explained by rotation of this structure. the NE lobe is redshifted whereas the SW lobe is blueshifted; if this assurnption is correct, this leads us to consider a torus seen almost edge-on rotating around a very young protostar. This is a strong argument in favor of the mm source as being the driver of the large scale bipolar outflow and confirms the hypothesis of a plane-of-the-sky outflow. The 13C0 hole dso observed by Nakano et al is then located on the mm source; extended in the NESW direction, it has the same position angle as the HCO+ structure. The inner contours of the millimeter continuum emission are aiso extended in the same direction (Fig. 1.6). Ail molecular condensations are related to V633 Cas C but no one is exactly coincident with its position in the velocity range mapped by Nakano et al the molecular peaks avoid this position. It is also remarkable to notice. as it appears from their maps, that there is no detection of a blueshisfted component to the L2C0J=1-0 emission at the position of V633 Cas C and it only gets important in the 2-10 km s-' range. Although as stated by Xakano et al. the 13C0 hole could be due to self-absorption. our observations show that it is most probably a molecular cavity due to the protostar. Assuming the cavity's size given by Nakano et al, less than 30". and a constant stellar wind speed. one can get an estimate of the age of the observed object-of course. on the assurnption that the depression was emptied by mass loss from the central condensation. The -first9 detection of the cavity in L2C0appears in the -5 to -2 km.s-' rnap and disappears in the 2 to 10 km.& map. Thus assurning a 7 km s-' speed range. we deduce a 3.5 km s-' speed for the cavity walls. With a 0.1 pc bubble, this leads to an upper limit of = 3 x 104 years! The dynarnical age of the outflow as measured by Levreault (1988). and recdculated by Nakano for a distance of 600 pc. is 1.1 x 105 years.
The point made above has been missed by SaW. Although they mention the molecular observdtions. they fail to make the association with the new millimeter source. This fact is crucial to the association of the large scale molecular outflow to one of the sources found in the region. Their 12C0 map shows that the region is probably complex in terrns of outflows since V376 Cas has some signs of outflow activity in the optical (Corcoran et al 1995). The elliptically shaped nebula with V633 Cas and its IR cornpanion at the apex also tend to show that one of those is the source of an outflow (Corcoran et al). An HH jet is also attributed to each one of V633 Cas and V633 Cas B. 1.3.2 The V633 Cas C Peak
Mie performed aperture photometry on the newly discovered source in .Aupst 1994 and in December 1995 (Table 1.3). -4s noted above. assuming that the source is gaussian-like in two directions and noting that the JCMT beam is almost gaussian as obtained from beam maps. we find a source size of 24" x 16". There is an error of at least 2" on each of the measured dimensions. The main contribution to the error cornes from the relatively low signal to noise ratio. mainly due to the weakness of the source.
1.3.2.1 Dust properties
Multi wavelength photometry allows us to deduce the dust emissivity index. 3. given the following assumptions: a relatively hot source making it Iikely that at mmfsubmm wavelengths we are in the Rayleigh-Jeans approximation and thus that Bu x v2. and that the mass opacity of dust can be written as K, sr v3. This leads to vFu = UK, Bu x v3+J.A linear regression through the data then yields 3.
Based on the flux values given in Table 1.3, we find /3 = O.%& 0.21 (Fig. 1.8). This fits well with other rneasurements for young stellar objects where we usually find emissivity indexes below 1 (Beckwith & Sargent 1991 and references therein). Of course al1 our calculations assume that the size of the source is independent of A. We note that if the source size increases with A, which goes with the increase of the beam size with A. then a smaller value of 3 will be found.
We have also tried to deduce the dust temperature of the object assuming an isothermal distribution of dust. Based on the 1.1 and 0.45 mm fluxes we find a color temperature of 70 K, for ,8 = 0.5. If we take an emissivity index of 0.35, the derived temperature rises to near 106 K! This proves the rveakness of such an attempt and clearly indicates the need for more precise data.
SaW found, based on their photometry, an emissivity index J = 0.8. They do not have a flux measurement at 450 Pm. Our resdts for the other three common wavelengths agree well with theirs. In fact a fit through these data in our case also yields P = 0.8. O bviously this straightforward linear regression through the data suIfers from serious shortcornings. The presence of a temperature gradient is one of these, since we assume an isothermd medium.
1.3.2.2 Mass calculations
Csing a11 measured parameters. source size. dust temperature. emissireity index. and distance, we can obtain a first estimate of the mass and total flux (Zinnecker et al 1992). We assume a unique constant grain temperature. gaussian beam and gaussian source profiles to estimate the total Aux. Here we approximate the shape of the source by a sphere. We also consider optically thin emission although evidences tend to show that. even at these wavelengths. the inner parts of the circurnstellar environment can be opaque (e-g., Beckwith S: Sargent 1991). Results are shown in Table 1.4. We have assumed a constant source size to calculate masses for different values of P. It is obvious that masses are very low and at the limit of hydrogen burning stars if it represents the mass of the star. The energetics of the large scale molecular outflow Iead us to conclude that we are most probably looking at the outer parts of the stellar cocoon in which a forming
- star is embedded, and not at the star itselt The big question would then be about the nature of the central object.
SaW used an arbitrary temperature of 30 K in their calculations, and a source size of 8" x 6". With their value of 0 this leads to a mass of 0.7 M.=,.Each mass estimate leads to diflerent conclusions OLI the nature of the object. CVe will discuss the implication of each of these results in the next section.
1.3.3 Discussion
The main question raised by this discovery concerns the evolutionary st at us of al1 the objects in this loose cluster. and in particular. the nature of the nedy discovered mm object , V633 Cas C.
One of the ciues lies in a detailed SED of V633 Cas C. Based on a few assumptions and data collected from the literature. ive built a first SED of L.633 Cas C. We first looked at a possible IRAS Aux distribution betrveen the different components found in the area (see Fig. 1.7). We assumed. following Lagage et al. that V633 Cas B's SED is identical to that of [RIAS 04363+'L333 and that al1 the 100 pm IRAS flux cornes solely from it and V633 Cas C; we estimated then the 100 pm flux of the latter as well as those at 60 pm and 25 Fm.Knowing that the source has not been detected in the visible nor in the near IR. upper limits can be inferred from observations at these waveiengths (Table 2 of Hatta et al 1992). We overlayd on the plot, a modified Planck curve of the form Fu = RB,(T)(l - e(-")). where the optical depth + = ($)P and fi is the solid angle subtended by the source. Lk used the temperature and emissivity derived in section 3.2. we cleariy see that Our estimated mid IR fluxes are fair. The source should then be bright at 100 pm and thus might weil be the one detected by Natta et al. This also tends to validate our basic assurnption, indicating that V633 Cas is a weak emitter in the mid-IR. Recent observations by Polomsky et al (private communication) failed to detect V633 Cas C at 10 and 20 Pm. The temperature and the emissivity were determined based on the hypothesis that the source size is independent of wavelength. But, one would expect a decrease in the size with decreasing A since Ive would be probing higher temperatures. Applied to our source. this gives a higher value of B. Nonetheless, if we accept the fact that Natta et al observed the new source, and since the extent of the object at 100 Fm is compatible with its mm size, this is an indication that the size of the source then does not change with X and the ,O value deduced is correct within the errors. However one should keep in mind that the measured size of the object depends on the sensitivi t ies achieved.
A11 this proves that great care should be taken when interpreting SEDs. especially when multi wavelength photometry is made with very different aperture sizes. It is also not clear if emission in the rnrnfsubmm part of the spectrum could be readily interpreted as arising from the protostellar environment. The classification of V633 Cas in the Hillenbrand et al scheme should be revised in this perspective. But what about V633 Cas C?
1.3.3.1 A Class O Object or an extreme Class I?
We now compare the 1.633 Cas C physical properties with that of typical Class O objects in order to learn something about the evolutionary status of V633 Cas C. André et al (1993) first introduced this Class of protostellar sources in order to account for very cold and still forming protostars. Class O objects are still in the process of accreting the bulk mass of their protostellar core. Bachiller et al (1996) give a list of ail known Class O objects to date. They note that some are the sources of very powerful outflows. These objects have ages not exceeding IO4 years. They usually have very low dust temperatures, Td = 20 - 50 K. Extreme Class 1 objects have been proposed by Lada (1991). They are Yearlyn Class 1 objects but seen through an edgedn disk so that they are more deeply embedded than "typical" Class 1 objects, expiaining the fact they can not be observed at near-IR wavelengths. Originally, Lada proposed the Extreme Class I objects in order to account for objects that were latter classified as Class O objects by André et al (1993). To distinguish them from Class 1 objects, André et al (1993) define as a membership criterion that 5 2 x IO4. Considering that LI., = 4*D2F~.~b. where D is the distance, fie3the flux density and Au(= 64GHz) the bandwidth. we find Ll3 = 3.2 x 10-~Lo. We caiculated the bolometric lurninosity by assuming that all the upper limits of Figure 1.7 are in fact detections. and we linearly extrapolated between two adjacent points. We found a bolometric luminosi ty
Lbol RS 180 Lo. This leads to zs 5.6 x IO4. If the SED obtained matches closely 1.3 that of V633 Cas C, then it is not a Class O object. To test the effects of a change in the SED,we used the upper limits determined from Lagage et a1 (1993) (open triangles on Figure 1.7) and ignored the one derived from Table 2 of Xatta et al
( 1992). The effect is minimal on the bolometric luminosity. ive find Lhol =z 14OL :. The peak of emission is around 100pm. ?latta et al report a measurement on LkHa 198 by Harvey et al (1979) of 72 Jy in a 37" aperture. -4ssuming that it is most probably the flux density of Y633 Cas C plus some contributions frorn other objects included in the beam. we lowered the dues attributed to the millimeter object at 100 Fm and 95 Fm so as to fit more closely the calculated modified Planck cuve. We find a bolometric luminosity of - 75 Lrr lowering the ratio 2.3 x 104, still higher than the criterion for Class O objects. L13
The dynamical age of the outflow and the more uncertain age derived frorn the dimension of the 13C0 cavity give also a clue to the nature of V633 Cas C. With an age bracket of 3 x lO4yrs A Class O object is still in the process of collapse and is still assernbling the bulk of its mass. A Class I is still accreting matter through an accretion disk but bas already assembled a large part of its main-sequence mass. Hence the circurnstellar mass around Class O objects is much more important than around Class 1 objects. André et al (1993)defined then Class O objects with the following M criterion ibfe;v < 1. Table 1 of Bachiller (1996) lists the envelope masses of Class O objects. Al1 are larger than the mass found for V633 Cas C. In our case. it is rnost probable then that Me;v > 1. From the preceding three arguments, we define V633 Cas C as most probably an extreme Class 1 object. at the boundary between Class O and Class 1 objects. The geometry of some of the molecular concentrations give some credence to this view. The disk like structure is not to be expected for Class O objects but rather a cored-apple geometry is more likely (André et al 1993). The molecular structures associated with the millimeter source show a Battened geometry ( Fig. 1.6), indicative. as we already pointed out, of the presence of a disk-like geornetry. Class O objects have been found to be associated with bipolar outflow acti- vities which is now thought to begin with the very first stage of star formation. Out flow and collapse are sirnultanmus. The outflow activity continues through the following stages of evolution (Class 1). In the Class 1 stage of evolution. the star is still adding mass to its core but through an accretion disk (Galli & Shu 1993a. 1993b). V633 Cas C is an Extreme Class 1 object. It is then the source of the large scale molecular outflow, although, as pointed out by Nakano et al (1990), the dynamics of the region is more complex and probably includes other outflows from the other sources of the group. This conclusion is drastically different than what can be drawn from the results of SaW. SaW found a mass of 0.7 hlo and assurned, without argumentation, a dust temperature of 30 K. This is in line wi th typical values for Class O objects. But one should note that they fail to associate the object with the molecular condensations of Nakano et al (1990) and they SED for the source. Furthemore, their size measurement gives a much smaller extent than our size estimate. Although our size estimate agrees well with that of Natta et al (1992). we should note that al1 these observations were carried out wi th different bearn sizes and different sensitivitiea. V628 Cas, better known as MWC 1080, is classified by Hillenbrand et al as a Group 1 object. These stars are characterized by power law spectra in the mid to far IR with X FA x X-4P which is believed, from theoretical models (Bertout et al 1988 and Shu et al 1987). to be the signature of an accretion disk. Their SEDs are successfuIIy modeled with such an hypothesis. They are considered by Hillenbrand et al to corne after the Group II stars in the H.-\EBES evolutionary sequence. Grankin ei a1 ( 1992) discovered it to be a binary with an orbital period of 2.SY69 days. Leinert et a1 ( 1991) using near iR speckle observations revealed the presence of a cornpanion at 0.75" to the east of the HAEBES. Since the Group I classification for V628 Cas is based on the IR excess. both these discoveries warrant a reassessment of the contribution of the star to the measured fluxes. Furthermore, other possible stellar cornpanions have been obsewed. One of them has been detected by Cohen and Kuhi (1979), 10" to the north-east of V628 Cas. CCD images of the area by Poetzel et al (1992) indicate also the presence of three objects (including the Cohen and Kuhi star), thus we would be in the presence of a small cluster with at least six cornponents. Hillenbrand (1995) covers in detail this group of stars. V628 Cas is also the source of a well known molecular outflow, observed by Levreault (1988), and found to be spherically symmetric. Poetzel et al took CCD images in [S II] and Ha lines, looking for HH objects and optical jets. They found that V628 Cas is associated with a large emission region and identified 6 knots axound the star. named A to F. in the continuum subtracted [S II! line image. These bots seem to be distributed on the edge of what they interpret as an outflow cavity showing a poorly coliimated outflow. Combining their data with radio observations by Curiel et al (1989), they conciude that the outflow is much more shock-heated to the east than it is to the west. The first evidence arising from the observations is that the emission at 1.1 mm (Fig. 1.5) is quite extended around the star. Taking into account the uncertainties arising from a somewhat variabie sensitivity in the inner part of the map. the first noticeable fact is the flatness of the emission in this area. although a few peaks emerge above the "background emission' . Overall. the emiss ion looks clumpy. possibly indicating clumpiness in the dust distribution. It is also more elongated in the east-west direction than in the north-south one with a p.a of == 100°. LVe have compared Our map with the CCD images from Poetzel et al (1992) to look for possible similarities between relative dust and gas distributions. corn- paring the peaks of emission at mm wavelengths and the position of the HH knots discovered. Although peaks 1 and 2 have no optical counterparts, we found peak 3 to fa11 exactly at the position of knot C. Are we dealing here with a dusty HH object? We looked then for possible morphological similarities. The emis- sion "plateau" at 1.1 mm seems to coincide almost exactly with the Ha emission. Another optical feature is also reproduced by the mm map. In the red CCD image, a "spiral arm" in the N-S direction goes from the star to the south of WWC 1080. - The ridge of emission in the mm, tracing an inverted S over the "plateaun, fits exactly this spiral arm, extending it to the N through peaks 1 and 2. It should be noted that this is an unresolved structure in the JCMT beam. 1.4.2 Dust Properties and Distribution A11 of these discovered rnorphological similarities pose the problern of the origin of the emission. We rneasured Bues at the peaks to deduce some of the dust properties. However, this is tricky since these are point-like features rising above the emission "plateaun. The results are shown in Figure 1-9. The 3 indexes derived are more uncertain than that of V633 Cas C. ,O1 is related to the western peak on the "ridge" of emission, pz is related to peak 2. Photometric data for both peaks are presented in Table 1.5. It is also not clear if the emission is tightly linked to V628 Cas. It could mise from the outer shell of a large excited region around the star. It is then related more to the cIoud than to the CS environment of the star. This would be in agreement with the findings of Di Francesco et al ( 1994). The? have observed MWC 1080 at 100 Fm with the KAO. performing a series of scans on the source. They measure a deconvolved source size of '19" in their Y direction and 13" in the X direction with both directions having a position angle of 2S0° and 11' respectively. It is then elongated in the EW direction and slightly tilted towards the South, with a position angle of s 100'. in agreement with our findings. .bsuming that the -size7 of the emission -plateauq. as measured on our map. gives a fair estimate of the extent of the source, we find a EW size of =: :35" and a YS extent of a 28". Di Francesco et al also find from modeling their observations that a large extended envelope is clearly needed to explain the observed characteristics of emission. This somewhat changes the properties for Group I objects as proposed by Hillenbrand et al ( W2), or changes the classification of V628 Cas. At least this clearly demonstrates the need for a large dust structure responsible for what we are realiy observing. One other hypothesis would Iink the emission to the circumbinary dust distribution, or more accurately to the circumcluster dust distribution and - thus is an intricate mix of influence from al1 stellar fluxes. Looking at the location of al1 known stars, one finds that they spread more in right ascension than in declination. Preliminary results from a polarimetric study of the central binary (Manset, private communication) give an inclination angle of the binary's orbit of 85" and show variations in phase with the orbital period indicating that. in the optical, the inner parts of the nebula are clearly iiiuminated by V628 Cas whereas in the mm this does not seem to be the case. Emission in these two parts of the spectnun are probably not sampliq the same dust grains. 1.5 Conclusions We have presented deep rnaps at 1.1 mm of regions around two HAEBES: a Group II object V633 Cas and a Group I object V628 Cas. CVe also presented photometry on the main peaks found in both rnaps. Both maps show extended emission. The V633 Cas map is dominated by a newly found source. k.633 Cas C. at 20" ?F W of the optically visible star. 1-63Cas. The mass associated with the new source is .- 0.1 Id,,. The emissivity index is very low. 3 = O.EfO.:! 1 provided that the source size does not vary with wavelength. However, the most striking feature is its location at the center of an HCOf elongated structure perpendicular to the large scale outflow direction. and also at the center of a I3CO cavity. This suggests that it is the source of the large scale molecular outflow observed by Canto et al (1984) and Levreault (1988). 4 first SED estirnate for the new source indicates a redistribution of mid to far IR fluxes between the different stars. making V633 Cas a weak emitter at these wavelengths. Based on these results and the SED estimate, we find that V633 Cas C is not a Class O protostar but an belongs to the extreme Class 1 type of objects. For V628 Cas, the emission looks clumpy with a few peaks superposed on a large extended flat structure of size 35" x 28" at a p.a. of 100". One of the peaks coincides with the position of a HH hot found in [S II] by Poetzel et al (1992). This region is a small cluster with at least 6 stellar objects already detected. surrounded by an extended common envelope. Waps with a higher resolution are needed to disentangle the web of influences of the different stelIar Buxes on the envelope. Because of the complexity of the emission, the emissivity indexes of the two measured peaks are uncertain. We note that the size of the Bat mm structure is compatible with the 10pm size measured by Di Francesco et al (1994). It is obvious from both maps that the common interpretation of mm fluxes as arising from accretion disks and circumstellar structure is not obvious at all. SEDs should also be interpreted with great care. maps are. for that purpose. an important tooi to check for the source of emission around an object. LVe would like to thank al1 telescope operators and support scientists ive have met during Our runs. They have heiped make it a smooth experience. R.H would like to thank Henry Watthews for fruitful discussions about calibration procedures and Claudine Kahane for interesting scientific remarks. This mork was supported by the Xatural Sciences and Engineering Research Council of Canada and the Québec Government. The James CIerk Maxwell Teiescope is operated by The Observatories on behalf of the Particle Physics and Astronomy Research Council of the United Kingdom. the Netherlands Organization for Scientific Research, and the National Research Council of Canada. Object Date Map Photometry V633 Cas 12-14 August 93 point-by-point - - 2-6 August 94 completed 1.3,1.1 & 0.8 mm - 23 December 95 0.45 mm V628 Cas 12-14 August 93 point-by-point - - 2-6- .Au gus t 94 cornpleted 1.3.1.1 91 0.8 mm T.4BLE 1.1. Log of Obserntions. 1.1 -4ugust 93 12.04 f0.54 1.36 f0.11 1.1 August 94 12.6 f 0.61 1.93 I0.18 0 -45 December 95 37.4 8.1 24.0 dz 1.6 TABLE 1.2. Selected Gaia and scding factors for difkent runs and wavelengths. 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The y ais is the Iogarithrn of the detected signal in mV scded by the known flux of the calibrator in Jy. -.-A-.....August 93 1 ----August 94 Sum 93+94 i i 1i 7 4 - - O 0.02 O.04 0.06 O 08 O. 1 Error (Jy) FIGURE 1.2. Error distribution for the V633 Cas map. l,fi,ll: 1; O 0.02 0.04 0.06 0.08 O t Enor (Jy) FIGURE 1.3. Error distribution for the V628 Cas map. ,.'l-r-lang 1 '- x optical position of LkHa 190 from Lagage et al (1993) - + center of offsets (HBC position) . t levels: 0.12 to 0.4 fy O - 20 -JO -60 RA Offsets (") FIGURE 1.4. Continuum map at 1.1 mm of V633 Cas showing the location of the new source. The shaded axes gives the beam size. 1 l 1 + optical position of WC1080 levels: 0.09 to 0.25 Jy ? FIGURE 1.5. Continuum map at 1.1 mm of V628 Cas. The shaded area gives the beam size. FIGURE 1.6. These figures are adapted from the molecular maps of Xakano el al (1990) (with kind permission frorn the ASJ). Figure a is a contour map of the integrated T', of 13C0 between -2.0 and 2.0 km s-'1 figure b shows Ti of HCO+ (solid lines), broken lines are for the F=2-1 HCN emission; figure c shows HCO+ and CS J=2-1, solid lines are for blueshifted emission, and broken lines for redshifted emission. The fdled circle shows the HBC position of V633 Cas - (offsetted by 17" to the north from its real position), the filled triangle shows the position of V376 Cas and the filled square, the position of the new source. HCO+ HCN FIGURE 1.6. (b) FIGURE 1.6. (c) gJCMT data Planck Curve: a IRAS data + hypothesis T=70X * upper limits 8=0.35 .upper limits O FIGURE 1.7. First SED estimate of the new source V633 Cas C. FIGURE 1.8. Emissivity index of V633 Cas C. FIGURE 1.9. Emissivity indexes for peaks 1 and 2 on the V628 Cas map. CHAPITRE 2 NGC 7129 Article scientifique préparé pour fin de publication NGC 7129: SVS 13 and Other Stars Roger Haü arl, P ierre Bastien1 Star Formation Group and Département de Physique, Université de Montréal and Observatoire du Mont Mégantic C.P. 6 128, Succ. "Centre-Villen, Montréal. Québec. H3C 137. Canada. and Daniel Nadeau' Département de Physique, Université de Montréal and Observatoire du Mont Mégantic C.P. 6128. Succ. "Centre-Ville". Montréal. Québec. H3C 1Ji. Canada. and Département de Physique. Université Laval and Observatoire du klont-hlégantic SteFoy. Québec. GlK TP1. Canada Received 'Guest Observer, Canada-France-Hawaii Telescope ABSTRACT Polarimetric observations carried out at the Mont Mégantic Obsermto- have revealed that SVS 13 has the highest linear polarization in NGC 7129- It &O shows on polarization maps the largest disk as determined by the angular distance between the two nul1 points. at the edges of the parallel vector pattern. .At the distance of NGC 7129. 1.25 kpc. the disk has a size of * 5600 .AU in the I band. The J band polarimetric map also shows a large disk. The polariza- tion of the star in a wide red bandpass is 7.5%. The measured emissivitp index. deduced from mm and submm continuum photometric measurements. 3 = 0.23 is compatible with what is found for Young Stellar Objects. We also report the discovery of 4 new infrared sources around SVS 13. Three are companions of SVS 13. The fourth star seems to be related to the IR source next to LkHû 234. -411 are clearly visible in the K band image and 2 of them are undetected in the 1 band. Polarization maps show that the closest cornpanion to SVS 13 cm not be the illuminating source of the reflection nebula. The two other companions lie on a shocked gas ridge delimiting the nebulosity in the J and K band images. .A large number of infrared sources is also present around LkHa 234. some detected for the first time. They al1 fa11 on or around a V-shaped enhanced emission ridge in the IR. also visible in Hz (Cabrit et al 1996) giving insight on their formation history. No stars are seen around BD +65" 1638 within a radius of = 15". -4 thin red ridge of emission is present to the East and South of this star with a number of stars outlining it. Based on this information, we confirm the following mode1 of star formation, first proposed by Bechis et al (1978): The formation of BD +65" 1638 trigerred the formation of the other stars through shocked molecular gas by the stellar wind. It is still an ongoing process where now SVS 13 and LkHa 234 are trigerring star format ion. Polarization: visible, maps-Disks: Circumstellar-Stars: pre-main sequence. 2.1 Introduction NGC 7129 is a region of star formation containing two well known Herbig ..\elBe stars (HAEBES) (Herbig 1960). LkHa 234 (V3ï3 Cep) and BD +65' 1637 (V361 Cep). The other two stars surrounded by nebulosity are BD +65O 1638 and one identified as SVS 13 (Strom et al 1976)'. which we will cal1 hereafter Noor. The estimated age of this group is. according to Strom et a1 ( 1976). not more than a few times 10' years. BD +65* 1638 is a main sequence object. and BD +65O 163ï is also considered as one or is at least very close to the ZAMS. It has been classified by Hillenbrand et a1 (1992) as a Group III object. This group contains H.4EBES with very little or no infrared excess. This is of course evidence for its classification as a main sequence object . Bechis et al ( 1978) studied star formation in NGC 7129. They proposed the following scenario to explain the observed CO line and far infrared continuum properties. BD +65O 1638 formed earliest in the cloud and produced an expanding shell which triggered star formation in the compïessed gas. To help understand the region Figure 2.1 shows a CCD image of the region obtained in the R band. with an integration time of 900 S. LkHa 234 has been the object of extensive studies. It has been assigned a spectral type of B5,7e by Strom et al (1972) and 08by Cohen S( Kuhi (1979). Thé et a1 ( 1994) follow the classification given by Strom et al. Recent near-infrared po- larimetric and imaging observations by Weintraub et al ( 1994. 1996) have shown that there is a cluster of at least 8 infrared objects within a radius of .- 10" around the HAEBES. One of them (PSI) has been first identified by analyzing polariza- tion maps and has been detected at 3.8pm. A large scale molecular outflow has been observed (Edwards & Snell 1983). Only the redshifted lobe has been detected in CO. Higher resolution molecular observations by Mitchell & Matthews (1994) '~notherstar labeled SVS 13 is known in the Serpens dark cloud. Since the name often refers to the Serpens star, we chose to change it for the sake of clarity. have shown it to be a collimated molecular jet. Its blueshifted counterpart is an optical jet observed by Ray et a1 (1990). The jet has a position angle of 255". -4 molecular ridge going from LkHa 234 to the northwest (Mitchell Si Matthews 1994) shows a sharp discontinuity to the southwest. towards a molecular cavity observed by Bertout (19Sï). A molecular peak (Peak 1) is observed on the ridge at an offset of (-25". S") from the Herbig star. Both jets are attributed to LkHn Z33. although this has been put into question by the newly detected IR sources. mai& PSI. However. as pointed out by Weintraub et al (1996). the polarime- tric disk position angle. as measured from the parallel vector pattern discernible at the location of PSI, is incompatible with this hypothesis. PSI rather drives an outflow revealed through the elongated shape of the cornetary nebula in a direction perpendicular to the 'disk". Two maser sources are also located near it. Weintraub et al (1991) associate one with LkHa 234 and the other with PSI. .Umost no data is amilable on Noor. The only reported observations we found wre IR photometric observations by Strom et a1 (1976) and Harvey et a1 ( 198.5) done while surveying NGC 7129. Based on observed J-H colors, the latter authors estimate from a simple model, photospheric temperature and spectral types for their observed sample: Yoor is assigned a B8 spectral type with an extinction -4r~= 8.2. 'loor. as well as LkHa 234. is located on the CO ridge. It is also very close to Peak 1. Ml this points to the fact that star formation is still active near LkHû 234 and Noor. This is in favor of the Bechis et al scenario, as pointed out by Mitchell S- Matthews. With the aim of studying the circurnstellar material around the two HAEBES in NGC 7129, we first obtained B,V and R direct images and an 1 band poiari- metric map of NGC 7129. This revealed that Noor has the largest polarization disk of dl the objects in the region, as this will be defined latter. We followed that by aperture polarimetry on Noor, LkHa 334, BD +65O 1637 and BD +65" 1638. We also obtained a J band polarimetric map and a K band direct image, and continuum flux measurements at mm and submm wavelengths to gain some insight on dust properties around Noor and. hally, we obtained a low resolution spectrum of the star to try to determine its spectral type. In section 2, we will present the observations, section 3 will discuss optical and near IR observations. with ernphasis on polarization data, and section 4 will be devoted to millimeter observations. Spectroscopic observations wiil be briefly discussed in section 5. The final section will be devoted to a discussion of the impact of our findings on the star formation history of NGC 7129. 2.2 Observations We observed YGC 7129 in September 1991 at the Mont Mégantic Observa- tory. Direct images were obtained in the B. V and R bands. We took short (a few seconds) and long (a few hundred seconds) exposures. We also took polarimetric images in the I band. For these. we placed a polarizer with the I filter and took unsaturated images at three digerent position angles of the polarizer. namely 0". 60". and 120' relative to the sky. Exposure times were determined to stay within the linear range of the CCD for the whole frame. Then. images ivere combined to get a polarization map and intensity image in the manner described by Asselin (1991). A11 images were taken with a 512 x 320 RC.4 CCD with a scale of 0.48" per pixel at f/8. Images were corrected by dark and bias substraction and were flatfielded using dome flats. We also corrected some cosmetic problems affecting some lines of the CCD. Aperture polarimetry was performed in August 1995 at the Mont Mégantic Observatory with Beauty and the Beast (Manset & Bastien 1995), a two-channel photon counting polarimeter. We measured al1 stars in a wide red bandpass determined by a Schott (RG 645) filter and the cut-off of the two RCA C31031A photomultipliers. The bandpass with this setup has a half-power bandwidth of 2411 8, centered at 7661 A. We also measured a P(A) curve for Xoor. and followed its polarization in the red fdter on different nights through the run. We used HD 161056 and BD +5g0389 as standard polarized stars. HD 154392. a standard unpolarized star, was observed to determine the instrumental contribution to the polarization. It is equal to zero within the erron. Millimeter and submillimeter photometry was performed through Canserv2 at the James Clerk Maxwell Telescope (JCMT) in May 1996. Thexommon user bolometer UKT 11 (Duncan et al 1990) was used to measure fluxes at 1.3. 1.1 and 0.85 mm. Sky conditions were variable and did not allow for 4-50 and 350 pm ob- servations. One calibrator (CRL2688) was observed at three different airmasses allowing the determination of an atmospheric optical depth ( r ) value for the night. One should note that this is not the best way to calibrate the data for atmospheric extinction but this was the only one available for that particular run. For a better approach see Stevens & Robson (1991). Weobserved Xoor through near IR filters with MO'iIC.-\ ( MONtreal Infrared C.\mera) (Xadeau et al 1994) coupled to a polarization module in June 1996 at the CFHT. The polarization module consists of a warm half-wave plate rotated at four position angles. O", ?Mo,45' and 67.j0,and a fixed angle polarizer located in the dewar with the filters. CVe used two polarized standard stars for zero point ca- libration. We made for each position of the half-wave plate a set of sky and object images in the following order {sky, object, object, sky}- giving two images on object and two on sky-. to cover four successive positions in a minimum time and allow at the same time for optimal sky subtraction. The CFHT console makes it possible to control the telescope from the data acquisition cornputer. Each complete turn of the half-wave plate leads then to 8 images of the object at each ZCanservis the Canadian service obaerving program at JCMT of the four angles or their equivalent angles. The sky image is subtracted from object images which are then flat-fielded using dome flats. Care is taken to remove the telescope and dome emission from dl frames. The reduced images aze then surnrned for each position. Appendix 2.8 gives Mueller matrix calculations for this particular setup. On 29 June 1996. we obtained a spectrum covering the 1500 ..\ to ZOO -4 range of Noor at the Mont Mégantic Observator. with a Boller & Chiveos grating spectrograph. -4 150 Ifmm grating was used. The spectrograph was coupled to a IK x IE; Thomson CCD with a 15 pm pixel size. Sky conditions were poor. but nevertheless sufficient. Table 2.1 gives the log of al1 observations. 2.3 Results for the Optical and Near IR Observations 2-3.1.1 Direct Imaging Figures 2.2 and '2.3 show a contour plot of Noor's nebulosity in al1 six bands. and turo greyscales of the J and h: images. The 1 to K intensity maps were obtained from the sets of polarimetric observations. The first striking feature is the change in the nebulous morphology while going from the blue band to the infrared. The other interesting feature is the gradua1 appearance of a number of sources in the nebulosity. The E; band image shows a count of 5 stars including Noor. The southern most object in the frame is in our view more related to LkHa 234 than to Noor as we will show when we discuss the HAEBES. Table 2.2 gives the declination and right ascension offsets of the stars relative to Noor as measured in the K band image. Figure 2.3 identifies al1 these sources. The morphological aspect of the nebulosit y reveals some interest ing features of the object. We noted above that the appearance of the various cornpanions considerably &ers the shape of nebulous ernission. Xoor 1 begins to affect it in the R band image. Thus, looking at B and V we can confidently Say that the major source of illumination in this case is Noor itself. The inner isophotes in these two bands show a fan shaped contour oriented to the West, tilted to the South. The emission is also brighter to the Southwest than to the Northeast of Noor. This could be the signature of a bipolar outflow driven by the star. It is obvious from the following frames that the infrared sources gain more and more importance in illuminating the nebulosity. Noor 1 extends the isophote in the YW-SE direction as is seen in R and I. In the J and K bands. Noor 2 ard 3 contributions give it a "spherical" shape. This *sphere" is bound on the east side by a large arc covering almost 180°, clearly distinguishable in the K band image and slightly detectable in the J frame. Noor 1 and 2 Ml exactly on this ridge. The ridge is clearly distinguishable in Hz images by Cabrit et al (1997) (Figure 3 of their paeer 1- We have cornpared Noor's location to other detected structures in NGC 7129 like the CO ridge observed by Mitchell di Matthews (1994). Xoor is located on the northern tip of the molecular condensation, to its southwestern limit. This explains the brightness contrast between the southwestern and northeastern parts of the nebulosity, since the northeastern part would be embedded in the high extinction CO ridge. Peak 1 in Mitchell & Matthews maps, located at an offset of (-25". 8" ) from LkHa 234, coincides with Noor 1, at an offset of (-35.91t, 9.3") with respect to LkHa 234, from Table 2.2. This somewhat explains the higher than expected temperature of Peak 1, which is 63 K. Although as pointed out by Mitchell & Matthews, a formed star would produce strong outflows and winds, thus large CO lines, whereas observations show narrow lines. The Peak is extended and slightly elongated in the NS direction. Keeping in mind that the observations have a resolution of 14" and that the diameter of the CO condensation is approximately 15". Peak 1 could be related to the K band arc. assuming that the arc delineates a dust overdensity and thus a CO concentration. This arc is thus the -progenitor' of two infrared stars and the possible source of CO emission. .Ail this assumes that the location of the Peak is determined to better than 1". The extent of the Peak encornpasses Noor 2, 3 and Noor itself. Although its NS elongation would favor the two peripheral infrared sources against the central stars. The polarimetric images are in the standard 1 and J bands. Vectors on the map (Figures 2.1 and 2.5) represent the orientation and amplitude of the elect ric field. The 1 images were binned with 2 x 2 pixels bins. The map has thus a 0.96" resolution. CVe used 3 x 3 pixels bins for the J band image. and a displacement of 2 pixels between tw-O vectors. The resoiution of the J map is then 0.5". Seeing conditions were excellent for both observing sites (= 1" for the 1 band at OMM and - 0.5" for the J band at CFHT). In order to determine in a non subjective rvay the centroids of the maps. we proceeded to calculate their positions using the orientation of the polarization vectors. We first determined coordinates for the intersection points of di pairs of perpendiculars to any two polarization vectors. To these (x, y) couples we associated an error (O,. O,). A 5 weighted mean gives a first estimate of the centroid's position. Then, an iterative process allows to reject al1 points over or betow n x u frorn the mean. The iterations are carried on until no points are rejected frorn the origindy calculated data set. A first discussion of this method is proposed in Appendix 2.9 and an anaiysis of the results and implications for the study of YSOs is given by Hajjar et al (1997) where it is used for a large number of objects. Table 2.3 lists the results for ail the polarization maps. Column 3 gives the rejection criteria used. A + sign in Column 4 indicates a weighted mean, a - says that no weight was applied to the data set. The initial number of intersection points and the number of points kept are in columns 5 and 6. Finally, the last two columns give the right ascension and declination offsets From the stellar peak of the calculated rnap centroid. AU the calculations are listed for illustration of the method for the I band image of Noor only. For al1 other stars they show a similar behavior. The results for Noor are compatible with the central star being the source of illumination of the nebula if one considers that the calculated a is an error on the location of the centroid. Overlooking the calculated value of a. we notice that the offset in declination is practically unchanged between the 1 and J band. But the calculated offsets in right ascension point to a possible peak displacement of e 1" to the east. This contradicts what is usually observed for other YSOs (e.g.. Asselin et al 1996 for V376 Cas) where the peak is usually displaced somehow in a direction perpendicular to the orientation of the disk. In our case. the disk has a position angle of 100'. This displacement can be due to calibration errors or to the fact that the profile is distorted in its outer parts due to low signal to noise and to the presence of the two cornpanions in the J band rnostly. In fact. a small error in the determination of the zero point for polarization angles will lead to a systematic shift in the calculations of the centroid. From declination offsets, assuming no peak displacement, the position of the centroid is located almost at the center of the line joining Noor and its closest infrared companion. 'loor 3. Another possible source of illumination is Noor 3, located also near the position of the centroid. yet another explmation coutd be proposed for the location of the centroid between both stars. The line joining Noor and its nearest companion is dmost perpendicular to the measured linear polarization (=17'). This might rnean that we are not seeing two stars but the two reflection nebulosities on both sides of the disk. The second peak would only be visible in the IR because of a possible slight inclination of the disk towords the SW giving a higher opacity on this side of the nebula. However this would be in contradiction with the brighter emission in the visible to the S W relative to the NE. Furthemore, the relative separation of both peaks is = 2.7". making it the largest separation rneasured arnong pre-main sequence stars with a known double ~eak.Such an effect was observed on V376 Cas with a 1" seprtration (Asselin et al 1996). Ignoring the last argument, a way to reconcile the surface brightness contrast between the two nebulous -sides" of Noor and the qualitatively estimated disk inclination, comes by noting that Noor lies at the edge of a high extinction region. the CO ridge. And thus. materiai to the East of the star is plowing into the dense medium without completely dissipating it. A favorable argument to this interpretation comes from the cornparison of the northeastern part extent relative to its southwestern counterpart. We measured on both maps the polarimetric disks revealed by the parallel vector pattern delimited by two nul1 points. Since these points are not always distinguishable on the maps. the following procedure was used to get the size of the disk: going away from the peak in two opposite directions and dong an imaginary line. passing through the peak. with a position angle similar to the linear polarization rneasured in an aperture. the two nul1 points are the two lowest polarization points encountered. The disk size is found to be smailer in the J band than in the 1 band, with sizes of 5.6" in 1 and 3.6" in J. We also estimated the disk inclination by comparing mode1 calculations of Bastien & Ménard (1988. 1990) with Noor's polarization rnaps, with the same method described by Bastien & Ménard (1990). The aspect ratio of the inner ellipsoidal pattern of vectors leads. in both the 1 and J maps to an inclination angle of = 75". Comparing both rnaps to the models leads to an inclination mgle > 80'. 2.3.1.3 Aperture Polarimetry As stated above, we observed through an 8.8" aperture dl stars surrounded with extended nebulosity in NGC 7129 with the RG 645 filter. A P(A) curve has been obtained for Noor. Table 2.4 gives the result of al1 these observations and Figure 2.6 shows P(X) and the position angle dependence on wavelength (@(A)). The central wavelengths and FWHM of each filter used. convolved by the PUT sensitivity curve. are given in column 2 of Table 2.4. They are two interferential filters centered at 1300 hi and 7600 A with a bandpass of 800 -4 (the 4102 .$ and 7607 hi entries in the table). a V band filter (5480 A in the table). RG 645 already described, RG 780, the 8330 A entry in the table and RG Y30 centered at YS4 -4. Columns 3. 4. 5 and 6 give the linear polarization and position a&e measured and their errors. The last two columns give the dates. Theta is constant within the errors except at 7600 -4 where a change of more thm 10' is observed, and at Y330 -4 where the change is larger than 5". Compared to other stars measured in NGC 7129. Noor shows the highest linear polarization of all the measured stars. It is also one of the highest polarizations for HAEBES. Its P(X) curve shows a maximum at 7600 '4. The polarization shows a smooth variation wi th wavelength. It increases from the blue to the maximum then decreases towards the red. Both parts of the curve show different behaviors. The red part of the curve shows a steeper decline than the blue part. Measurements in RG 645 do not show any time variability except for one point showing a noticeably lower value of polarization. Of course. when looking at P(X), we exclude RG 645 (7661 .&) data since they encompass a much larger part of the spectrum thm other filters and are an average over dl contribu- tions to the polarization thus giving a lower P. Both red filters, RG 830 and RG 780, show two different polarization values indicative of a possible time variable behavior . The position angle has an associated absolute error of more than 1" due to uncertainties on the position angle of standard polarized stars. In fact. in our data. zero point determination for each of the two standards observed showed a difference of 1'. Error dues in Table 2.1 are detemined using Serkowski's formula (oe = 28.S0 x 7);they only indicate uncertainties due to photon counting statistics. The @(A) curve shows a l0 on al1 points for which oo < 1'. Apart from the two measurements at 8330 A (RG 780), no time variability is detected at other wavelengths. It is most probable from the data that the star is time variable in the red. Our aperture is large enough to include possible contributions from the infrared companions of the star. Although the polarimetric map discriminates against possible contributions to the integrated polarization of Noor 3. it does not give any clues as to any possible contributions of the other 2 nearby stars. The only possible argument for the effects of the companions is the different behavior of P(X) in the red compared to the blue. 2.3.2.1 Direct Imaging Figures 2.7 and 2.8 show LkHa 234's nebulosity for al1 6 bands. and greyscales for the J and K images. It is also striking to note the change in shape going from B to K and the appearance of a number of sources, some of them aiready reported by Weintraub et al (1994, 1996). Table 2.5 lists al1 observed sources chat are thought to be related to the HAEBES, giving their offsets relative to it; they are also identified on the K band greyscale image in Figure 2.8. For the sources not Listed by Weintraub et al, we follow their naming convention. We should note that the first traces of IRS 2 in the isophotes is in the V band image, IRS 4 and IRS 13 are visible in the long R exposure (Figure 24, whereas IRS 3 is only detected in the K band. IRS 9 and PS 1 are seen as two extended nebulosities. two cometary nebulæ. Polarimetric evidence (Weintraub et al 1994) point to the presence of an embedded object illuminating this northwestern cornet ary nebula. In their K Band polarization map, a polarimetric disk is detected at the position of PS 1, with a position angle of a 45". A centrosymmetric pattern, pointing to PSI as the most likely illuminating source, extends to about 10" from PS 1 except that it follows rather closely the nebulous object IRS 9 and breaks on the YW cometary nebula (se Figure 3 in Weintraub et al 1996). hlthough this could be due to a low signal to noise in the K band polarimetric image. it may indicate peculiarities in the dust distribution around the Herbig star. The presence of a polarimetric disk would suggest that PS 1 is the driving source of the jet and the outflow except that. as pointed out by Weintraub et al, the polarimetric disk is almost parallel to the direction of the outflow. Ail models of dust distribution around YSOs lead to a polarirnet ric disk perpendicular to the out flow orientation (e.g.. Bastien S; Ménard 1988. 1990. Whitney S( Hartmann 1999). The recent high resolution mid-IR observations by Cabrit et al ( 1996) have detected an H2 jet which is not pointing towards PS 1, and which they rename IRS 6. They also have detected PS1 at 10 Pm. They speculate on the presence of yet another more deeply embedded source at = 2.8'' to the west of PS 1. The morphological aspects of the nebulosity add sorne dues to the above rnentioned problems. The B and V images show fan-shaped profiles extending to the SW. LkHa 234 is obviously the sole illuminating source at these wavelengths. It is an indication of a bipolar outflow, or at least the remnants of it, driven by the HAEBES. The K band image shows some interesting features: the two previously described cometary nebulæ. On closer look, one sees two long nebulous "feathers* extending in two different directions. The one related to PS 1 extends to the NW at a position angle of ;i 300' then bends toward IRS 12. The IRS 9 nebulosity has a SW orientation at an angle of zs 210". It bends to the south near IRS 11 then follows a path passing slightly to the west of IRS 13. IRS i and IRS Il fdl on it, IRS 6 is very close to it. The optical jet has a position angle of 25.5". aimost exactly at the bisector of the angular sector formed by the two plumes of nebulosity. Both features form a cone with its apex close to LkHa 734. almost half way between PS 1 and IRS 9. Looking at Figure 2.1, one sees that a thin lace of ernission extends from LkHa 234 to the North then forms a curve around BD +65O 1638's nebulosity. A nurnber of faint stellar points lie on or next to it. We count 9 objects, including IRS 13. Their coordinates, as measured on the R image are given in Table 3.6: the Table does not include IRS 13. Al1 authors who studied XGC 7129 have also singled out LkHa 234 since it is thought to be the youngest visible star in the cloud. Mitchell Sr Matthews (1994) showed that the outflow is collimated and that the star lies on the edge of the CO ridge they found. .A comparable feature (an ernission enhancement ) is also seen in the maps of Bechis et al ( 197s). Studying star formation in this region. they find a density enhancement to the east of the reflection nebulosity. as rneasured from 13CL60emission. .A '*CO temperature increase is observed to the north. south and east of the nebulosity. The width of the arc delineated by this emission is unresolved in their 70" beam. Both the density increase and the high temperature ridge closely follow the R band lane. LkHa 234 coincides aith a I2CO emission peak and is the brightest far IR continuum emitter in the nebula. They attribute the flux density enhancement at 100pm relative to the weak mid IR Buxes to a large dust envelope emptied in its inner part by a stellar wind. The discovery of d the IR sources relaxes this constraint. The contribution of each one to the IR flux and the energetics of the region is not known. The 1 band polarimetric map of LkHo 234 gives a hint as to why the measured aperture polarization is low relative to the one measured on Noor. -4 close look at the inner parts of the map (Figure 2.9) shows two paralle1 vector patterns almost perpendicular to each other. The first one has a position angle of .- 120" and is centered 2" to the east of the HAEBES. almost at the position of IRS 8. The second one is more a distortion of the centrosymmetric pattern. It is discernible by two adjacent vectors, both parallei to the declination axis, and located Y to the west of LkHa 234. This could be the first signs of the presence of the embedded source PS 1. Exomining the centrosymmetric pattern in the -outer" parts of the nebula. we see that its radius of curvature is larger than what would be expected if the Herbig star was the only illuminating source of the nebulosity. This points at least to the fact that the central source is extended. Yevertheless. two points delimiting what can be viewed as a disk are identified on the map. The western null point is probably afTected by the distortion of the vector pattern from centrosymmetry, probably due to the presence of a nurnber of embedded sources. The two null points define a size of - 5.5'' with a position angle of z 130°. This is in agreement with the orientation of the jet and probably indicates that LkHa 934 is most probably the source of the large scale molecular oütflow observed. Its contribution to the illumination of the reflection nebula is probably rapidly decreasing in the infrared, in favor of other younger sources formed in the wake of the shocks due to the bursting material from the HAEBES. The second known HAEBES in NGC 7129 is considered to be very near the main sequence, and has a spectral type of B2,3e (Racine 1968). As seen in Figure 2.10, the nebula around it is slightly elongated in the NS direction and shows almost the same size in dl bands. .-2 number of sources are detected only in the R and 1 bands. A jet-like structure is visible in the R band, probably due to the fact that the R band inchdes emission lines, and leaves traces in the B and V bands. This filamentaxy stmcture, named RF by Miranda et al (1993). contaios one HH object (NGC 7129fHH 1) discovered by the same authors. Imaging and spectroscopy of the structure lead them to conclude that it is a reflection nebulosity produced by an outflow from a star they name SI located at the northeastern tip of the structure. It is not visible in our images. Their interpretation of the narrowness of the structure relative to other observed illuminated nebulosities relies on the fact that it is at the edge of the rnolecular cavity and that Bertout ( 1987) had detected motion along the walls. thus giving an explanation to the collirnating mechanism. In the 900 s R exposure (Fig. 2.1). another structure similar to the previously mentioned -jetg' is seen to the northwest of BD +6j0 1637 at a position angle of 225". Its direction passes by the location of S 1. Two stars are seen embedded in it. This may indicate that it is at the origin of the structure. One should note however that the two stars seem to contribute to the illumination of the structure as rvell as the Herbig star. Miranda et al note that in their CCD images, the inner parts of RF are darker than its frontiers. This is in agreement with the previous hypotheses. Both RF and this structure are the two walls of the collimated outflow originating in SI. The polarimetric map of BD +65" 1637 (Figure 2.1 1) in the 1band is peculiar. It does not show any of the "usud" features seen in other maps. Despite the apparent absence of a centrosymmetric pattern, the position of the central illu- rninating source is compatible with the position of the peak. It means that other objects in the nebula have no significant contributions on large scales. Their effects are restricted to their immediate neighborhoods. Figure 2.12 shows the isophotes of the B. V. R and I images of this main sequence star. The nebulosity, at least for its outer contours. looks -spherically symmetric". ..\lmost al1 stars seen in the frarne fd1 outside a radius of ir 15". except for one possible source located to the north of the star. +t an offset of (-O.*ja, 5.8") and is barely detected in the 1 band. It is not seen in the R exposure. Furthemore, the brightest stars in the area tend to show an arc extending more than lYOO. The berpart of the arc towards the star is completely devoid of stellar sources. The polarimetric map (Figure 2-13) shows a centrosymmetric pattern to a large distance from the star. Its syrnmetry is slightly broken to the South and completely lost to the northwest. It looks noticeably disturbed at the location of the two bright objects to the south and southwest of BD +6.j0 1638. The inner parts of the map do not show any particular pattern of vectors. This is due to the fact that the star fdls on a bad line on the CCD which we were not able to properly correct with the cosmetic image. We note also that the centroid of the map seems to Ml far from the peak for a weighted meân but is closer to it for non-weighted calculations. This dso is easily explained by the position of the star on the CCD. The central vectors are incorrect but they also carry the highest signal to noise so that they contribute the most to the weighted mean whereas, for a non-weighted mean, the outer vectors will prevail. As seen in Table 2.4, the star is probably time variable in linear polarization. However, there is no indication of any variability in the position angle. 2.4 Millimeter and Submillimeter Photometry We have observed BD +65O 1637 at l.lmm in 1993. We only have an upper limit for its flux. Noor was observed in May 1995 through Canserv. We have obtained fluxes at 1.3. 1.1. O.S. and 0.45rnm. An East-West scan at 1.1 and 0.8 mm through Noor allawed us to effectively subtract any possible background con- tributions from the data. Five measurements se~aratedby 3'. with a chop throw equal to the separation. were obtained. We calculated Noor's flux by bootstrapping al1 five measurements. Table 2.7 shows al1 the results for our observed stars. The Bu': of BD +6j0 1637 is a 30 upper lirnit. We only have a very crude estimate of the 450pm flux. Calibration data were not useful at this wavelength. We estimated the opacity by assurning that the gain was equal to the published value by Stevens & Robson (1994). Figure 2.14 shows calibration fits for the 3 longest wavelengths. Using only the three well calibrated points, we estirnate the dust emissivity index of the object to ,L? = 0.12 k 0.67. Figure 2.15 shows the fit to the data. It is obvious that. one of the three data point is off its exact value. At 1.3mm. the JCMT beam is not well coupled to the 65" aperture used for photometry. It is the most probable source of the observed discrepancy. However, within the error. the ernissivity index is still found to be < 1, typical of what is found for YSOs (e.g. Beckwith k Sargent 1991). During the observations. a number of consecutive five-point measures were obtained on Noor, mainly to center the telescope on the peak of emission. Centering dways gave large offsets and showed and extended emission region. The photometry given in Table 7 has been obtained at the nominal position of Noor (SVS 13) given by St rom et a1 (1976). We also derived its right ascension and declination from the 1 band image assuming that the position of LkHa 234 given in the Thé et al (1994) catalogue is well known. This showed that the Strorn et al coordinates are accurate to better than 1". Nevertheless, we tried to follow the measured offsets to try to locate a probable peak of emission. They point to an offset of (5 Mlf,- -12") from the optical position. This is almost coincident with IRS 2 in the vicinity of LkHa 234. One should note that 80 Fm observations by Bechis et a1 (1978). and 30 and 100 pm maps by Harvey et al (1985) point to the HAEBES as the strongest emitter. However, the 100 pn opticd depth map in the 1983 paper shows two opacity peaks: one on LkHa 234 and the other one near Xoor. This region is also coincident with the CO ridge observed and Peak 1. the CO peak at T= 63 K measured by Mitchell & Matthews ( 1994). 2.5 Spectroscopy of Noor The spectrum siiown in Figure '2.16 is a first attempt at getting the spectral type of Soor. The slit position angle is equal to 90'. The width used gave a FiVHM of 31 .A on the CCD. This means that almost al1 lines observed are unresolved. Y'o quantitative measurements can thus be made on the spectrum. Some lines are identified on the spectrum. The strongest emission lines corne from [O I] and are most probably nebular in nature. Ha and HO are seen in absorption. and are photospheric, as well as Na 1. Based on the criteria used by Cohen S: Kuhi (1978), the presence of both H lines and Na 1 hints to a possible binary nature of Xoor. Possibly an -4 or B type star with a K or SI cornpanion. The dificuit task of defining a continuum makes it difficult to ascertain the reality of some of the observed structures. It is also obvious that even with the large FWHM, some midentified lines are resolved. They show width in excess of 3000 km si1. Ha and [O I]X6300 A tend to show a P Cygni profile, dthough the [O IIX55i7 a feature does not. All this makes for an intriguing object. All the evidences discussed in this paper show that Noor is a YSOs: but its spectrum does not show Ha in emission as expected for a pre-main sequence object. We should note that the resolution of our spectras being very poor. much remain to be said about Noor. and much higher resolution spectras are a must to solve the enigma of Noor. 2.6 Discussion 2.6.1 Interpreting Polarization Maps Table '2.8 lists dl the polarization maps in this work dong with the map of LkHa 231 by Weintraub et al (1994). Colurnn 2 gives the wavelength, and colurnn 3 the resolution of the map, and in parenthesis the binning used to calculate each vector if it is different from the resolution. A nurnber of models have already been calculated md theoretical maps published. Most are based on Monte Carlo calculations of a radiative transfer model. The first published model. by Bastien & Méoard (BM) ( 1958), was based on a two component structure of the circumstellar environment, an optically thick disk and two optically thin bipolar lobes. The model assumes two scatterings of photons near the disk and single scattering in the lobes. Gledhill's (1991) model is based on ar unresolved polarized source. Gledhill also addresses the effects of a tilt in the position angle of the bipolar nebula relative to the symmetry axis of the central polarized object. Most recent calculations are Fisher's models (Fisher et al 1996). They use theoretical density profiles for disks and envelopes to obtain polarization maps. It is obvious from these that we should be able to get the density from the inner pattern of the polaxization map. Although one should note that al1 actual models assume a single central object and none addresses yet the presence of a binary or multiple system illuminating the nebula. Apart from this. a common feature of dl rnodels is the parallel or eiliptical pattern of vectors in the inner parts of the map. for disks seen edge-on or close to edge-on. Another cornmon feature is the presence of tw~polarization null points marking the edges of this pattern. In the BM model, they are interpreted as transitional points between the optically thick disk and the optically thin envelope. They delirnit what is commonly called the polarimetric di&. This hypothesis predicts a change of disk size with wavelength. The BM model makes no quantitative estimate of this effect. According to tbis interpretation. one should be able to deduce the density structure of the disk if the optical properties are given. Knowing that: for r = 1 at the polarization null points. where a, is the column density in g~rn-~and K, is the opacity in cm'g-'. The major obstacle is this coupling between dust properties and densities. One should point out that polarimetric disks are usually of the order of a few thousands .\Us and should be seen as Rattened structures rather than centrifugally supported disks whose sizes are of the order of 100 .W. Noor shows a change in disk size with wavelength from 5.6" in the I band to 3.6" in the J band. The position angle of the disk also seems to change from - 120" in the I band to 105' in the infrared. One should keep in mind that the resolution in the J band map is twice that in the 1 band. The precision on the disk size estimates should then be viewed accordingly. The difference in position angle can be due to calibration errors but can also be the result of the changing contributions of the different sources to the illumination of the nebula. An andysis of the method mentioned earlier is presented by Hajjar & Bastien (1997) and Haiar et al (1997). We will not expose it in much detail here since the number of ndpoints measured is insuflicient to fully exploit the effect described earlier. It is mostly important to note that the change of the disk size with wavelength is observed for Noor and for a number of other YSOs (Hajar & Bastien 1997. Hajar et al 1997). 2.6.2 Individual Objects From the imaging data. it seems that Xoor is the center of a 'quiet' outflow phenornenon. The shock front observed in the K band and in Hz (Cabrit et al 1996) is a clear indication of the physical association between the star and the molecular ridge ( SIitchell S; SIatthews 1994). It is not clear if Noor 2 and Noor 3 are a product of the shock. They are not detected in H2. Noor 4. the closest cornpanion to the star. does not seem to affect the polarization pattern, although it might be too deeply embedded to produce any significant distortion in the map. Figure 2.17 shows a plot of alI detected IRS objects around this HAEBES. The two plain lines show the orientation of the two IR feathers described in section 3.2.1 and the broken line the orientation of the jet. It is noteworthy to see that the jet is the bisector of the cone formed by the two plumes of nebulosity. The location of ail stars fdls on or very near the two cavity walls. The cavity region is devoid of any observable star. The origin of the jet seems to be the subject of a debate in view of the recent findings of Weintraub et al (1994) and Cabrit et a1 (1996). Nevertheless, it looks highly probable that an outflow event has been triggered by LkHo 234. The cavity wds. result of a shock. were then the scene of very recent star formation. The number of stars associated with the northern cavity wail is smaller than the one related to the southern wail. as seen in Figure 2.17. This could be explained by the fact that the major density enhancements are located to the east and south of NGC 7129. Noor. probably formed at the same epoch. has also contributed to the dispersion of cloud material to the XW of the HAEBES. 2.6.3 A Star Formation Scenario Yo stars are visible to the east of BD +6s0 1638. -411 are to the west. north and south. The first detectable objects near the central star lie on a arc at a radius of - '20". LkHa '3%. Yoor. and BD +6s0 1637. ais0 seem to form an arc roughly paraIlel to the first one. Thus we propose the folloiving scenario for the star formation history of NGC 7129. The scenario follows closely the ideas put forward by Bechis et al (1978). The formation of BD +65O 1638 and its stellar wind triggered the formation of the other stars in the group. Noor. BD +6j0 1637 and LkHa 234 formed as a consequence of an advancing shock front into the dense molecular gas. The outflows of these objects triggered amund them similar events, as pointed out by Mitchell & Matthews and by our observations. The blueshifted outflow from LkHa 234 forced its way through dense materid and thus is not visible. The redshifted lobe burst open into the already fomed molecular cavity delimited by the R emission ridge thus showing this connection between the IR feathers observed and the shock front from BD +65O 1638. In fact. the Cabrit et a1 (1996) images do show that the R ridge is detected in H2. 2.7 Summary and Conclusion We have presented polarization maps for the main objects of NGC 7129 (see summary in Table 2.8). In addition, we have detected a number of new infrared stars around the major memben of NGC 7129. Our main results can be surnmarized as follows: 1. Our primecandidate Noor look younger than previously thought. and much more embedded. although its spectrum apparently does not show Ha in emission. It shows the highest linear polarization of al1 stars measured. -- 7%. and its polarimetric map indicates the presence of a large scale disk seen almost edge-on. It drives an outflow that produces a shock front when plowing into the dense medium to the east of the star. probabiy triggering the formation of two objects. It has 1 detected IR companions. two of which are probably related to a CO condensation. 2. LkHa 234 has at least 14 IR companions within a radius of z 20". Some of them are reported for the first time. The distribution of these objects relative to two plumes of emission delineating an outflow blown cavity indicates a possible shock triggered formation process. The polarimetric rnap of the object shows that the nebulosity is illuminated by more than one source with the relative contributions being wavelength dependent. 3. The overail distribution of stars in NGC 7129 favors a star formation scenario where BD +65O 1638 formed first, triggering star formation by a wind sweeping the cloud material around it. This formed the presently observed molecular cavity and triggered star formation on the advancing shock front. YGC 7129 is stiU a very active site of star formation and is a prime candidate to study shock t riggered star formation. We would Like to thank ail TACS for giving us telescope time to pursue this project. Special thanks to Dr. Lorne Avery who was responsible for the Canserv observations reported in this paper. This project was financially supported by the Xatural Sciences and Engineering Research Council of Canada and the Québec Government. The James Clerk Maxwell Telescope is operated by The Observatories on behalf of the Particle Physics and Astronomy Research Council of the Cnited Kingdom. the Xetherlands Organization for Scientific Research. and the Yational Research Council of Canada. This research has made use of the SIMBAD database operated by CDS in Strasbourg, France. APPENDIX 2.8 Mueller Matrix Calculations The polarization module consists of a rotating half-wave plate analyzer and a futed polarizer, assernbled with the filter. The retarder matrix (Serkowski 1968) for an ided half-wave plate takes the form where O is the angle of the plate axis to an arbitrary zero-point. It is obvious from there that d = &î + n x r/2, where n is an integer. is equivaient to O = O. The polarizer mat rix. following Serkowski. is given by where a is the angle of polarizer's axis to the same zero-point as the half-wave pIate. Let be the Stokes vector representing the incident Iight beam and its dueafter passage through the polarization module. They are related through the following matrix equation: In an astronomical context, the zero-point of polarization measurement is the declination axis and the angles are measured counter clockwise relative to i t . The IR chip measures only the first Stokes parameter, the intensity I. We. then measure it at four different positions of the half-wave plate separated by 22..j0. For each position. the intensity is related to the first three Stokes parameters. from equation ('2.3). we then der ive: for 0 = 3. il = [Io + Qo cos(4P - 24+ Iosin(4d - .'a)]x 112. for tp = ,fi + 22.5", 1, = [Io - Qosin(@ - 2a)+ (I, cos(4Q - %Y)] x 112, for Q = /3 + 45.0°, r3 = [IO - QOCOS(~P - 2~) - Cro sin(@ - 2a)]x 112, for Q = p + 67.j0, 1.= [L+ Qo sin(@ - 2a) - Llo cos(4p - Za)]x 112, where 4 is the angle between the sky zerepoint and the hdf-wave plate zero position. Thus and a are two fixed constants of the module. Once (48 - Za) is measured, it is fixed unless we lose track of the half-wave plate rotation. From equations (2.4) to (2.7), we can hdboth Uo and Qo, as well as Io: .-\Iternatively, by observing a standard polarized star, we can get the angle going into the trigonornetric functions: where Ii, = li - Ij. 2.9 Polarimetric Map Centroid Determination Polarization maps show 'far" from the central source a centrosymmetric pattern and close to it an ellipsoidaily shaped ~rofilewhich shows sometimes in its inner most part a number of parallel vectors. The ellipsoidal pattern has been interpreted by Asselin et al (1996) as due to an extended source illuminating the nebula. Using perpendiculars to the vectors one should be able to point to the source to an accuracy depending on the signal to noise in the nebula since the precision of the angle of polarization 8 is, according to Serkowski (1968), 48 = 28.S0 x y. 2.9.1 Caleulating the centroid For p vecton of polarization we would have rn = 9unique pairs of perpendiculars each contributing one intersection point to the set of data. We have perforrned our calculations in the reference frame of the CCD. For each vector. we have its location (2,y), its polarization angle and its error. 6 is determined in the usud astronomical convention, relative to the declination axis and counterclockwise. North being oriented in the y direction on the CCD. the dope of the normal to the vector is thus where B is expressed in radians. The equation of one line being y - pi = ai x (x - xi). where (xi,yi)are the coordinates of the polarization vector position. the coordinates of an intersection point and their associated errors are: -4 1/u2 weighted rnean and associated error is calculated for the whole set of intersection points. The iteration process first rejects al1 points above or below n x o from the mean and then recalculates the weighted mean and associated dispersion. It stops when no more points are being rejected. 2.9.2 Main Contributors As seen from equations (2.16) and (2.18), both errors are inversely proportie na1 to the slope difference (ai- aj)and thus for adjacent polaxization vectors. the error is large whereas for two perpendicular vectors the error is very smail. The mean will then be deterrnined essentially by intersection points of perpendicular -rays9. So the weighted rnean effectively removes the geometrically error proue determination and keeps the most solid ones. One other problem is -noisev polarization. If the sky was effectively sub- tracted. we may safely assume that noise polarization is random and thus al1 contributions to the centroid determination from these vectors will cancel out in the mean. The iterative process will then remove al1 possible glitches in the data. From Table 2.3, one sees that a 1.50 threshold is too restrictive, at least for our data. We suggest taking a rejection factor between 2 and 3 o. Name X Date n x t(s) PL Direct Imaging NGC 7129 B Septembre 91 10 x 10 N B 1x100 ?l V 10x3 Y v 1x100 Y R 10x40 R 1 x 900 Y 1 10 x 10 Y SVS 13 J June 96 26 x 1S2 Y K June 96 12 x 1:3..j2 Y S pect roscopy 5liIIimeter and Submillimeter P hotomet ry TABLE 2.1. Log of Observations. ------'P indicates whether or not a polarization map was obtained. In that case integration times are for individual positions of either the polarizer (in the opticai) or the hdf-wave plate (at J and K). 2Tirnes for each frame are coacids of shorter integrations. roi Noor's Cornpanion Stars Name RA offset (") DEC offset (") TABLE 2.2. 'loor's cornpanion stars. Name A(AA) A P (%) ap 0 06 Date Xoor 4402 (736) 5-04 0.13 96.7 2.1 70 August 95 7607 (916) 7.27 0.19 110.3 0.6 20 August 95 7661 (2411) 6.86 0.12 98.2 0.5 16 August 95 0.4 18 August 95 0.7 0.5 20 *August 9.5 0.5 1.5 11 August 9.5 1.2 0.9 16 August 95 1.6 30 August 91 1.1 16 JIugust 9.5 :3.3 130 August 91 '1.0 16 August 9.5 2.8 30 .-2ugust 91 TABLE 2.4. Aperture polarimetry of XGC 7129 stars. Name RA Offset (") DEC Offset (") -- - LkHû 234 O O IRS 2 -11.27 -0.91 IRS 3 0.92 -12.78 IRS 4 1.71 8.41 IRS 6 -2.74 -6.07 IRS 7 -4.13 -9.87 IRS S IRS 9 IRS 10 IRS i 1 IRS 12 IRS 1:3 PS 1' Soor 1 TABLE 2.5. LkHû 234 cornpanion stars. 'This position is the hed of the cometary nebula where PS 1 is embedded, the uncertainty on this position is higher than for the others sources and is zz 0.5". The same is valid for IRS 6 and IRS 9. Cabrit et ai have renamed PS 1 as IRS 6. We chose to keep the nomenclature of Weintraub et al since they clearly detected the source at 3.8 pm (Weintraub et al 1996). Name a (hh mm ss) 8 (O ' ") NGC 7129/R 1 21 41 56.34 65 52 16.6 /R 2 21 41 54.40 65 51 55.1 /R 3 21 41 51-28 65 52 00.6 /R 4 21 41 49.32 65 51 55.1 /R 5 21 11 49.13 65 51 12.4 /R 6 21 41 46.95 65 51 38.6 /R 7 21 41 43.24 65 51 28.7 /R 8 21 11 41.82 65 51 26.9 TABLE 2.6. Coordinates of sources tracing the R boundary line. TABLE 2.7. Millimeter photometry of sources in YGC 7129. Object X (pm) Resolution(bin) (") Disk Size (") Reference Noor 0.9 0.96 5.6 1 1.25 0.5(0.75) 3.6 1 LkHa 234 0.9 0.96 -5.5 1 *> 3 dm- 0.1 -1- BD +65" 1637 0.9 0.96 1 BD +65" 1638 0.9 1.9'2 1 TABLE 2.8. Polarization maps and related parameters. References: 1) This work, 2) Weintraub et al (1994). Asselin. L.. Bastien, P., Ménard, F., Monin, J.-L. & Rouan. D. 1996. ApJ, 172. 349 Bastien. P. &. Ménard, F. 1988. ApJ, 326, 334 Bastien, P. & Ménard, F. 1990, ApJ, 364. 232 Bechis. K.. Harvey. PM.. Campbell, M.F. k Hoffmann. W.F. 197s. ApJ.226.139 Beckwith. S.V.W. & Sargent. -4.1. 1991. ApJ. 381.250 Bertout, C. 1987. in Circurnstellar Matter. IAl? Syrnp. 122. 1. Appenzeller S- C. Jordan eds.. '23 Cabrit. S.. Lagage. P.O.. McGaughrean. S1.J. S: Olofsson. G. 1997. ;\&A. 321. ."j3 Cohen. M. Sr Kuhi. L.V. 1979. ApJS. 41. 7-43 Duncan. W.P.. Robson, E.I.. Ade, P.A.R.. Griffin, M.J. S- Sandell. G. 1990. MNR-AS, 243, 126 Edwards, S. St Snell. 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Gatley, 1. 82 Merrill. K.M. 1996. ApJ.468.1.5 Weintraub. D., Kastner. J.H. 9r Mahesh. A. 1994. .\pJ.4'20,87 Whitney, B.A. & Hartmann. L. 1992, ApJ. 39.5. 529 NGC 71 29 FIGCRE 2.1. R band image of NGC 7129. Al1 stars appear double-peaked because of a guiding error. North is up and East is to the left. This is the full CCD image with a dimension of 4' x 2.5'. Noor (=SVS 13) is the northern most star surrounded with nebulosity. LkHn 234 (V373 Cep) is located to its southeast. and BD +65O 1637 (V361 Cep) to its southwest. 10 O -10 IO O -10 10 O -10 RA Offset (") FIGURE 2.2. Contour plots of Xoor in al1 6 bands observed. The letter in the corner gives the passband. Offsets are from Noor. FIGURE 2.3. The J and K images of Noor in grayscde. The numbers in the h: band image identify the cornpanion stars of Noor. O - 10 R.A. Offset (arcsec) FIGURE 2.4. 1 band polarization map of Noor. Polarization vectors axe overlayed on the intensi ty contours. R.A. Offset (arcsec) FIGURE 2.5. a) J band polarization map of Noor and b) Polarization vectors overlayed on the J band intensity contours. 5 O -5 RA. Offset (arcsec) FIGURE 2.5. (b) FIGURE 2.6. P(A) and @(A) curves for Noor. O -10 -20 O -10 -20 O -10 -20 RA Offset ("1 FIGURE 2.7. Same as Fig. 2.2 but for LkHa 231. FIGURE 2.8. Same as Fig -2.3 but for LkHa 234. The nurnbers are for the IRS objects. h 1 - 20 10 O - 10 - 20 R.A. Offset (arcsec) FIGURE 2.9. Same as Fig. 2.4 for LkHa 233. 20 O -20 20 O - 20 RA Offset (") FIGURE 2-10. Same as Fig. 2.2 for BD +65" 1637. R.A Offset (arcsec) FIGURE 2.11. Sarne as Fig. 2.4 for BD +6j0 1637. 20 O -20 20 O -20 RA Offset (") FIGURE 2.12. Same as Fig. 2.2 for BD +6j0 16.35. RA. Offset (arcsec) FIGURE 2.13. Same as Fig. 2.4 for BD +65" 1638. FIGURE 2.14. Cali bration fits for 3 millimeter and submillimeter wavelengths. FIGURE 2.15. Fit through the photometric measurements of Yoor. 3 is the dust ernissivity index. FIGURE 2.16. Nom's Spectra. RA Offset (") FIGURE 2.17. This figure shows the location of al1 detected stars around LkHa 234 relative to the K band detected cavity walls and the [S II] jet observed by Ray et al (1990). CHAPITRE 3 HL Tau, une interprétation des cartes de polarisation Article scientifique préparé pour fin de publication HL Tau: A polarimetric map interpretation Roger Haiar1 Sc Pierre Bastien1s2 Star Formation Group and Département de Physique, Université de Montréal md Observatoire du Mont Mégantic C.P. 6128, Succ. "Centre-Ville", Montréal, Québec, H3C 1J 7. Canada Received ; hccepted - - - 'Guest Observer at the James Clerk Maxwell Telescope 2Guest Observer at the Canada-France-Hawaii Telescope ABSTRACT We present a new high-resolution I band polarization map of HL Tau. Con- sidering previously published J. H, and K band maps, a change in the size of the polarimetric disk is measured and an interpretation based on one of the prediction of the mode1 by Bastien Sr Ménard (1988, 1990), and the Monté-Cor10 calculations of Ménard (1989). is explored. We calculate from the displacement of the pola- rizatioo nul1 points a density distribution in the disk. assumed to be a flattened structure of infalling material raining down on the r =z: 150AU centrifugally sustained disk. We calculate a disk densi ty profile for HL Tau and find a power-law volume density at mid-plane p(r) = po(ro/r)a with 1.1 5 a 5 1.9. We calculate a mass of = O.llt&, based on a constant thickness di& and a disk radius of 2000 .W. h deep point-by-point map at 0.8 mm, obtained at JCMT, is also presented. The peak is unresolved but superimposed on weak extended emission. This shows the presence of an extended envelope. Based on a published dust emissivity index Mgas of 3 = 0.6. a dust temperature of 35 K and an interstellar ratio for . we dust calculate a circumstellar mass of 0.05 M,3. 3.1 Introduction HL Tau is one of the most studied young steiiar objects of the Taurus cloud. It has ben used as a test bed and very often it is one of the first objects to be observed with ernerging techniques. Most recently, interferometric observations at mm wavelengths have been obtained showing the presence of a centrifugdly supported disk around the star (Mundy et al 1996. Lay et al 1994). In the opticd and near infrared, the emergence of adaptive optics have been used to obtain vey sharp images of the central parts of the nebula (Close et al 1997). HST observa- tions (Stapelfeldt et al 1995) first revealed the absence of any marked stellar peak in the inner 1" in the visible but showed a C shaped iliuminated profile attributed to the star light illuminating the walls of a cavity created by an outflow. The shadow of the disk has been observed in the near IR at a similar resolution as well as the C shaped profile (Close et al 1997). The location of the source seems to have been identified in a Ii band image (Close et al 1997). IR obserntions by Weintraub et a1 (1995) and Beckwith & Birk (1995) showed a displacement of the peak position of 0.59" from J to K. K being most probably the position of the star since Beckwith & Birk find no peak displacement between H and K. This is explained by considering that the observed light shortward of K is reflected light in a wind blown cavity (the C shaped profile seen in HST and adaptive optics images). Weintraub et al find that the J.H.Ii positions are al1 different and assign. based on their polarimetric maps a location for the star about 0.5" from the K position. Beckwith Sr Birk based on the color gradients and rneasured peak displacement conclude to the presence of a steep density gradient within a fer hundred AUs explaining the high extinction at the star and the low one in front of the nebulosity scattering star light. One should note the presence of two different structures in the extended nebula. The first one. resolved only in high resolution optical images, has an extent of less than 1" and is delineated by the two illuminated arcs, making a C-shape profile, probably due to an inner cavity blown by the jet. The other is a large scale nebulosity showing a wide opening angle (Beckwith & Birk), also seen in high-resolution single dish CO observations (Monin et al 1996). -411 recent studies show that HL Tau is most probably still deeply embedded or at least hidden behind an optically thick disk hampering any detection of the star at wavelengths shortward of K. Molecular obserntions of the star formation group (Monin et al 1996 and Cabrit 1996) indicate that HL Tau is probably at the edge of the original molecular core, with a disk inclination of around 70". the north-eastern lobe of the flow tilted towards the observer. The redshifted part of the outflow is still digging into the dense molecular gas. leaving no signature in the opticai. The disk of HL Tau, and more generally of T Tauri stars. is now well established. HL Tau shows two major disk like structures. The centrifugally sustained accretion disk with an extent of about 100 to 150 .AU has been readily detected at a plethora of wavelengths; Mundy et al (1996) at 2.7 mm. Close et al (1997) in the near infrared (Ii band), Rodriguez et al (1994) at centimetric wavelengths. Lay et ai ( 1994) in mm interferometry. A larger disk-like structure. extending up to a radius of 2000 AU, has also been detected in molecular obser- vations (Sargent S: Beckwith 1987, 1991). The authors conclude to the presence of rotation in this large scale disk. Hayashi et 01 (1993),observing with the Nobeyama array. conclude to an accretion dominated motion in the structure. The last such observations, by Cabrit et a1 (1996), also find no evidence of rotation and di sagree with a pure accreting mode1 of the elongated structure. Other "indirect" evidences come from linear polarization measurements, de- tection of highiy collimated optical jets and polarization maps. In polarization maps, our main point of interest, the signature of a disk is the presence of a parallel or eilipticd profile in the inner parts of the map, delimited by two null points. In the outer parts, the nebula shows a centrosymmetric pattern due to single scattering of star light on dust grains. Numerous models have been developed to interpret polarization maps. The first such model was worked out by Bastien &. Ménard ( 1988. 1990, hereafter Bhl). The model was a two component model: a disk where photons were scattered twice. and a bipolar cavity where on- single scattering occured. For the first time, polarization maps were interpreted correctely by scattering alone. without any need for aligned grains. Monté-Car10 calculations performed by numerous authors confirmed the initial results of the BM mode1 (e.g.. Ménard 1989. Whitney & Hartmann 1992. Fisher et ai 1996). They dl consider a disk and outflow cavity geometry. Another interpretation was proposed by Gledhill (1991). His model assumes a polarized source emitting light mostly towards the pole and then scattered in the nebula. His model do not provide a mechanism to explain the polaxized light of the central source. hl1 the models produce the general patterns described previously. None of the rnodels provide detailed quantitative tables on the polarization maps obtained. Bastien & Ménard (1990) have calculated disk inclination angles from a visual comparison of model maps with observed maps. They derived an inclination of 70' for HL Tau. Gledhill's paper provide a thorough discussion of the effect of tilted disk axis relative to the outflow axis and Fisher et al calculations show that polarization maps can provide a powerful tool to get a hondle on the densi ty distribution of the circurnstellar material. In this paper. we propose a novel way to obtain density information from polarization maps. It is based on a prediction of the BM model which interpret null points as a transition between the optically thick (multiple scattering) region of the large scale disk and the optically thin (single scattering) materid of the cir- curnstellar environment. This leads to the fact that we should expect the position of the null points to change with wavelength. The polarimetric disk gets smaller with longer wavelengths. We will discuss the results of this approach on HL Tau, based on our high-resolution I band polarization map and published J. H and K maps (Weintraub et ai 1995), and the 7500 -4 rnap by Gledhill Sr Scarrot (1989). CVe wiY also present a deep continuum map obtained at the James Clerk Maxwell telescope (JCMT) at 0.8 mm. 3.2 Observations B. V. R and 1 band images were obtained at the Canada-FranceHawaï telescope (CFHT) in 198'7 December 23-28. The I band polarimetric map was obtained during the same run. We used an KR polarizer placed in the filter wheel with the 1 filter. at a position angle of 0° relative to the North-South axis. We obtained a set of images at three different position angles of the polarizer. namely 0'. 60°. and 120'. For that purpose. we rotated the telescope's bonnette at 60° and -60".The CCD used is the 640 x 1024 RCA 2 with a plate scale of 0.108"lpix at f/S. An overscan of 32 lines is added to the raw image allowing for a zero order bias subtraction. The seeing was zz 06". Five exposures were taken at 0'. and four at each of the other two orientations. In each set. three had integration times of 600 S. with XZ Tau saturated, and one was a 120 s image. The additional frame at 0' position angle is a 180 s image. More details are given in Table 3.1. Images were calibrated the standard way by removing bias and dark and flat-fielding with dome Bats. The 1 band image were dl flatfielded with a set of flat frames obtained at O" since the polarizer position angle did not change relative to the CCD axis. The polarized frames were then corrected for aimas. We used an extinction per airmass value deduced from available data at the CFHT and UKIRT web sites, T zz 0.1. A sky fit was subtracted from the individual frames which were then shifted and added to produce a final set of three images. The polarization maps were then extracted from this data. Totai exposure time for the 0° polarized image was 2100s and 1920s for both 60' and 120'. We corrected the latter two for integration time before producing the final map. 3.2.2 Submillimeter The Y00 pm rnap was obtained at JCMT with the UKT 14 common user bolometer (Duncan et al 1990). The JCMT has a 15 rn dish and can operate at millimeter and submillimeter wavelengths. The 800 Pm map was obtained in August 1994 and completed through the Canadian service observing program Canserv in December 1994. It is a point-by-point map with an integration time of 300 s per point. reaching an average sensitivity of 0.05 Jy. Figure 3.1 shows the distribution of the errors for al1 the points in the map; the nmow peak is around 0.05 Jy. Cdibration was pelformed in the standard way. We observed Mars and one secondary calibrator. NGC 7538 IRSl (Sandell 1994). and assumed a constant opacity for each night. Pointing errors were sornetimes large. rr 5". Some parts of the maps were reobserved in order to account for the large pointing offsets that were found. Table 3.1 gives the log of submm observations. 3.3 Polarization 3.3.1 Polarization maps Figure 3.2a shows the 1 band polarization map and 3.2b the same vectors overla-d on intensity contours. The lowest contour is at 3 sigrnas from the background. At the northwestern edge of the map, the vectors seem to deviate from the generd pattern seen in the other parts of the map. This is due to a ghost image probably resulting from reflection on the lens or filters. This reflection contaminates the image in the long exposures. It does not affect the inner profdes nor the results deduced from the map since we are mostly interested in the inner parts of the nebula. We clearly distinguish the two null points on the map. They define a polari- metric disk 4.1" in diameter. Two other ways of showing more clearly the position of the ndpoints are displayed in Figures 3.3a and 3.3b: 3.3a) contours of constant polarization position angles, and 3.3b) contours of constant P x 1. Two nodes are clearly seen on these maps: in 3.3a, they are the two intersection points of al1 the equal position angle curves, and in 3.3b, they mark the strangulation in the inner contours due to the very small polarization. They identify, within the errors. the position of the nul: points. Table 3.2 gives the measured size of the disk (column 3) along with the wavelength and resolution (coiumns 1 and 2) of al1 the available rnaps in the literature including the I map of this work. The last column gives the reference for each of these polarization maps. Within the limits determined by the polarimetric disk. a parallel vector pattern is clearly distinguishable. Its position angle is equal to the disk's orienta- tion, 14 1'. in very good agreement with previous linear polarization measurements and our own detemination from the map (Table 3.4). Moving outward. the vector pattern turns elliptical up to the limits of Our map. This is indicative of an ex- tended illuminating source for the nebula. It is not clear from our map where it becomes circular. At the position of the peak, vectors are clearly deviant from the otherwise - paralle1 pattern seen between the two nul1 points. We measured the linear polari- zation in two different circdar apertures centered on the peak. In a 2" opening, the polarization is as high as 16.6% with a position angle of 151". whereas in a 8" aperture, we îmd for both quantities 12.05% and 139". Both measurements are inciuded in Table 3.4. To test for the BM model hypothesis. we measured the polarimetric disk size on the near IR maps published by Weintraub et al (1995) (Figure 3.4). and on the map of Gledhili Sr Scarrott (1989). The results are included in Table 3.2. The precision of these measures is obviously less then the 1 band estimate since the near IR maps have a lower resolution because of pixel size. They rvere also not optimized for the determination of null points on the map. The K band estimate is just an upper limit since the Weintraub et al rnap can not show the polarization of the inner 1". It is clear from this set of data that the size of the optically thick region is srnaller with increased wavelength. This confirms the prediction of the BM model. Although predicted. none of the amilable models has already dealt with such an issue. This is also an important test for the validity of competing models. especially since recent calculations by Matsumura S. Seki ( 1996) showed that single scattering by aligned grains could produce the observed patterns (centrosymmetric with an -embeddedn parallel pattern and null points). 3.3.1.1 Polarimetric Centroid We have calculated the polarimetric centroid of the 1 band rnap using the rnethod described in Hajjar et al (19Wb). Basically, we calculate the intersection points of all perpendiculars to the polarization vectors and determine a mean position through an iterative process, rejecting too deviant points. Table 3.3 gives the results for wious rejection thresholds; offsets are given relative to the position of the peak. Figure 3.5 shows the two dimensional histogram of dl intersection points. It is interesting to note that the profile matches closely the 2-D histogram obtained from the intersection points of the perpendiculars to an ellipse (Hajar et al 1997a), with one of the axes of the "cross" pardel to the position angle of the disk. We note here that we were unable to reproduce the result of Weintraub et al (1995) relative to the position of the polarimetric centroid as well as the precision with which it was determined in their paper. We do think that great care should be taken when considering this kind of calculation since it is prone to a number of systematic errors. We leave the full discussion to the paper by Hajjar et a1 ( 1997a). 3.3.2 An aperture polarimetry review A number of linear polarization measurements axe available in the literature. We also measured it on Our images in two different apertures and one measurement comes from data obtained by Manset (private communication). hl1 are listed in Table 3.1. The first column gives the wavelength or filter used, the other colurnns give repectively the circular aperture used for the measurement, the linear pe larization and its error. the position angle and its error, and in the last colurnn the reference for the given measurement. No definite conclusions can be drawn lrom the data since there is no complete set, from the UV to the IR. that was taken in a short time interd if not simultaneously. Furthemore. no overlap is available between the different authors. If two or more measurements are available at the same wavelength they are taken with different apertures. This is the case for the two values of the polarization at 5895 A or in the K band. Hodapp (1984). based on the data available at that time. found that the P(X) curve for HL Tau is suitably fitted by a modified Serkowski law. One should note that he mixed data taken with similar, albeit different, apertures. Two measurements are availabie at 5895 and in the I band. It is interesting to note that in both cases the smallest aperture yields a larger value of the polarization. The position angles do not show the same behavior at both wavelengths. The aperture in the I band were centered on the peak: the small aperture has been matched to the extent of the deviation of the pardlel vector pattern seen in the map. If the pattern found in the I band is sirnilar to the one in the V band it's then normal to expect that the position angle within a 4" aperture is different than in a 2" aperture. This explains the discrepancy in the behavior of the position angle with aperture. 3.4 Far infrared dust distribution Figure 3.6 shows the deep Y00 pm map. The shaded circle to the right gives the beam size. 13.5". The cross gives the optical position of the source from the Herbig Sr Bell Catalogue (1988). Two components seem to contribute to the emission at this wavelength. The peak of emission is unresolved but looks superimposed on an extended plateau of emission extending to the South and slightly in a EWdirection, with a position angle of l%OO. For this anaiysis. we have assumed a gaussian beam profile with HPBW of 13.5" and looked at cuts through the peak of ernission at different orientations. They al1 show an unresolved core with significant emission over the wings of the gaussian fit. Figure 3.7 shows a gaussian fit to an azimuthal plot of the source. The curve is the best fit gaussian to the data. One clearly sees a wider component evident from a spread of the data points. Another feature is the -arm9 extending to the South then to the East of HL Tau. It remains to be seen how it correlates with features observed in interferometry. In fact rnaps by Cabrit et a1 ( 1996) as well as the maps of Sargent 8r Beckwith (1991) do show some emission extending to the South of HL Tau. Based on the emissivity index value found by Weintraub et al (1989). and a color temperature of 35 K derived from 1.1 and 0.8 mm photometry, we calculated a mass of .- 0.05 Mo for the circumstellar environment. We assumed that the source is unresolved at al1 wavelengths and isothermal dust. This is in agreement with the value obtained by Weintraub et a1 (1989). 3.5 Discussion 3.5.1 Nul1 points reviewed -A paradigm for the interpretation of polarization of starlight by the circum- stellar dust of YSOs has emerged in the past few years. The general characteristics of the observed polarization maps are well reproduced by scattering alone. multiple and single scattering. The circumstellar geometry now agreed upon requires the presence of an optically thick disk with an extended envelope (outflow) (Fisher et al 1991. 1996, Whitney 9c Hartmann 1992, Bastien Sr Menard 1988, 1990). The fine details of the geometry and dust distribution are still being investigated. as welI as the exact nature of the dust grains. One recent work involving the last topic is the work by Chrysostomou et al (1996 & 1997) where the authors need small metallic grains (ruling out pure silicates) to interpret their circular and linear polarization observations. They also conclude to the presence of a large scde disk on top of the centrifugally sustained accretion disk to account for the degree of circular polarization seen in the nebula. The authors used a Gdli % Shu ( 1993) densi ty distribution in their Monté Carlo code cdcuiations. However. up till now no direct measurement of the density in this large scaie structure has ever been performed. We argue that null point determination provides. in a simple and elegant way? çuch a possibility. And we will try to apply what follows to the case of HL Tau. Monté Car10 calculations by Fisher et al (1994, 1996) show a clear depen- dence of the inner pattern of polarization vectors on the density profile. The BM model, although based initially on very simple assumptions, predicted a depen- dence of the null points loci with wavelength, an effect which is clearly detected for HL Tau and for a nurnber of other objects observed (Hajar et al 1997a,b) at different wavelengths. The BM mode1 also argue that null points fa11 at the position were the line of sight optical depth is unity. Following this line of reasoning, one could easily show that: where a, is the column density dong the line of sight (in g cm-2) and K, the opacity per unit mass of dust gains. Obviously, one needs a dust grain model. The approach proposed by Chrysostomou et al (1997) seems to bear fruit for the particular problem of dust optical properties and composition. along aith other observations. We will use, for the sake of our calculations and for cornparison purposes. a number of dust models taken from the literature. The most recent one has been developed by Henning 9: Stognienko (1996) designed for protoplanetary accretion disks. The two other models are due to Pollack et ai ( 1994) and Preibish et a1 ( 1993). It is not straightforward to convert oc to a volume density since Ive are sampling in our line of sight different parts of the circurnstellar environment. In a cylindrically symmetric distribution of circumstellar matter. the volume density is a function of r and z. p(r. z). A general form of n, is given by: B G(rd = ds)k (13.2) where A and B are the intersections of the line of sight with the boundary of the cylindrical environment around the star. ds is a path element along the segment [AB]. Since (ds)* = (dr)2+ (rd@' + (d~)~, and since z and û can be expressed in terms of r, we can write ds as a function of dr only, Replacing ds by its expression as a function of r. this leads to Figure 3.8 shows the definition of al1 the variables and parameters in the preceding equations. In Our current understanding of the 'layering" of the circumstellar environ- ment of YSOS, it is safe to assume that the highest density occurs at ro on a given line of sight. Disks around young stars are considered geometrically thin. Obser- vations also show that they show a large density gradient along z. For HL Tau. the works of Beckwith St Birk (1995) have shown the presence of a steep density gradient along the z axis. near the star. They estimate it at not more than a few hundred .Ali's. Knowing from interferometry that the radius of the large scale disk is about 2000 .4Cs (Sargent Si Beckwith 1987. 1991) and that the structure is unresolved along the direction of the outflow. this is evidence for -thin disks" around YSOs, with ratio a 5 10 at least, where R is the radius of the disk and H its scale height. It is safe to assume that if a disk is slightly tilted the extinction is mostly due to a narrow region around ro. If the disk is seen edge-on. the Iine of sight will probe straight through the disk giving a very high degree of extinction. If the estimate of .4\-(~23) towards HL Tau. at the estimated position of the star. is any indication, over the whole radius of the disk the optical depth in V is T 2 20. In the case of an edge-on disk, measuring at different line of sites we will probe deeper and deeper in the disk structure. This approach can help convert oc to actual disk densities p. It can at least give a better knowledge of its structure. Let us assume that we have a finite number of measurements for n wavelengths, XI, ..., A,, where XI is the shortest wavelength. For the i-th determination. we have from equation 1.5: R r , dr. (r2- r:)~ This is a weil known inversion problem, the Abel integral-equation which has an andytical solution. To give an idea of the solution we will rather try to fit a power law to the column density and use the exact solution to get an idea of the actual disk densities. It will obviously be a power law also. The column density is written: Knowing that the exact solution of the Abel integral is given by: this leads to: CVith the change of variables r u=- tH we find: putting we end up with where B(z, w)is the beta function, an Euler integral of the first type (Gradshteyn & Ryzhik 1980), which can be evaluated in terms of Gamma functions as (Abramowitz & Stegun 1965). The volume density in midplane of the disk then a+l takesthefomp(r)=po(?) where ..ssuming then an exponential vertical structure such that : P with H = Ho (6), leads to a total disk mas: This should be viewed as a first attempt to use the technique to get the structure of the circumstellar material. It is clear that uncertainties on grain models and nui1 points hamper, for now, any reiiable inversion of our data. The nurnber of actual data points amilable and the discrepancies in resolution are some of the major sources of error. the other being the dust model chosen. as can be seen from Table 3.5. 3.5.1.1 HL Tau From Table 3.2 we can get the size of the HL Tau disk at different wavelengths. Applying the previousiy developed assurnptions to these points ive are able to get the densities of the disk at different radii. The results for three different grain models are listed in Table 3.5. Column 1 gives the wavelength. and column 2 the radius in arcseconds. The next three columns give, for each grain model. the surface density calculated with equation 1.5. We note here that we used an interstellar ratio of gas to dust masses, = 100. The mode1 of Henning Sr Stognienko (1996) is concerned about the effects of dust aggregates and thus includes composite and Buffy grains as well as ice mantles around the grains. The model of Pollack et al (1994) is sirnilar to the grain model of Henning & Stognienko but uses a different chernical composition of the basic grains. Preibisch et al (1993) built a three component model made up of amorphous carbon grains, silicates and ice coated silicates with amorphous carbon inclusions. We have used their data for the ice coated silicates. In fact, assuming a photospheric temperature of 5000 K and a power law temperature gradient T = To(r/ro)-0-5.typical for a Bat spectrum source like HL Tau (Adams et al 1988), we find at 80 ..\ü T z 10 Ii, well below the emporation temperature of the ice mantles considered by Preibisch et al ( 1993), T = 11-5 K. We are then justified to use coated grain models for al1 our data points since they aii fdl at radii of at least 80 .AC'. Figure 3.9 shows logarithmic plots of the colurnn density versus the radius with linear regression fits to the data for each of the grain models used. The radius is calculated assuming a distance of 160 pc to the Taurus cloud (Kuhi 1964). CVe find a power law index 0.1 5 a 5 0.9 (cc= (ro/r)")for the column density. This leads to a volume density distribution at midplane of p(r) = po(ro/r)' with y = 1.4-1.9, depending on the grain model, with po given by equation 1.12. The errors on the size of the disk are of the order of the resolution of the maps or higher in the case of published maps since we had to estimate visually the position of the nul1 points. The resolution is given in Table 3.2 for al1 the polarization maps used. This shows the effect of grain models on the density profile obtained. CVe assumed a unique grain rnodel for al1 the disk. and also asswned that the disk of HL Tau is seen edge-on. Grain properties may well change in the disk since temperature and density change with radius. Furthermore, the inclination angle of the disk of HL Tau is - iOO (e.g.. Sargent SL Beckwith 1991. Bastien & Ménard 1990). A rigorous inversion of the column density measurements requires a larger dataset and more precise determinations of disk sizes. Aowever, to test our model we calculated the disk mas based on the following assumptions: a disk radius of 2000 AU. roughly the size of the large scale structure observed by Sargent & Beckwith (1987. 1991); a constant scale height H = Ho = 80 AU (/3 = O in equation 1.13) almost equal to the peak displacement observed by Weintraub et al (1994) and Beckwith & Birk (1995) between the J and K band images, and consistent with the density gradient inferred by Beckwith & Birk from their observations; a constant volume density in the inner 0.1 AU of the disk, p(r 5 O.1AU) = p(O.lAU). Table 3.6 shows the results of these calculations. Column 2 gives the value of a for the given grain model. and po are the column density and volume density at a radius of 1 AU, and the last two columns give the calculated disk mas. is calculated from the linear regression fits to the data shown in Figure 3.9. The mass estimates are in the range of values found with other methods. Lay et al (1994) model their millimeter interferometry observations with a disk and find a lower limit to the inner disk mas zz 0.03 Mo. Yundy et al (1996) find a mas of 0.05 to 0.07 Mg. Their model also suggest a shallow radial dependency to the surface density in the disk, S x r-' or less. This seems to contradict our results but both could easily be reconciled by choosing a suitable vertical scde height to the disk since S = J:: p(r. :)dz However. ive should note that their model. as well as ours have many ad hoc assumptions. Cabrit et al ( 1996). on the basis of CO observations. attribute an upper limit of 0.2 M., for the circumstellar material around HL Tau. A mass of 0.03 Ma is adopted by Hayashi et a1 ( 1993) for the disk: they find a range of disk masses .LfD = 0.032-0.11 hl5. Other estimates based on millimeter and submillimeter continuum observations give 0.11 !di(Sargent S- Beckwith 1991). 0.04 M.;,(Weintraub et al 1989) and 0.05 Mx. from the Y00 pm map and flux. It is interesting to note that the total mass is not so sensitive to the grain models, at least for the ones we used. Since we assumed that the disk is seen edge-on, we are underestimating the true volume density. The mass should then be viewed as a lower limit to the actual mass, obviously in the limits of the assumptions considered above. The volume density profiles calculated for HL Tau, for the chosen grain models are compatible with the theoretical profile for the protosolar nebula (e.g., Cassen & Surnmers 1984, and references therein). A fdl comparison of Our results with other calculatioos and models for young stars and the protosolar nebula is developed by Haüar et al (1997a). The mass found is weii above the minimum mass for the protosolar nebula (0.01 Mo). The gain model obviously plays a big role in determining the density profile of the disk. We assurned it to be in the form of a power law but a thorough analysis requires a reliable inversion of the .4bell equation, an aim unrealistic with the present data. Errors on the disk size are of the order of the K band disk size and the number of points is small. Xevertheless, we note that a power law seems to be a very good approximation to the density profile for all grain models. Grain rnodels should also be investigated more thoroughly. We assumed a unique grain model for al1 the disk. unaffected by changes in density and temperature. This is most probably not the case since. close to the star. grains should at least lose their ice mantle. Furthermore. clustering of srnaller grains is expected to be dependent on density. as well as other disk parameters. A change in the grain size distribution should then be expected with disk radius. 3.6 Conclusion Based on published polarization maps and a new 1 band polarization map. we have proposed a new method to extract the disk density from the displacement of the nul1 points as a function of wavelength. This is the first method altowing to measure 'directly* the density in the disk. We have used, to test the method, three published gain rnodels and obtained a power law density profile for the disk p(r) = po (F)' with 1.4 5 7 5 1.9. This is compatible with the theoretical profile - usually used for the protosolar nebulae and is in line with theoretical models of accretion disks around YSOs. We calculated a disk mass for HL Tau based on oufindings, using a constant scde height for the disk and a radius of 2000 AU height taken from the literature. The mass estimate we find, 0.09-0.15 Mo, is comparable to other mass estimates from millimeter photometry and molecular line observations. With the densities found at 1 ?LU for the different dust models and assuming the dust opacity used by Hildebrand (1983) and commonly used in interpreting millimeter data (- 0.1 cm2 g-' at 250 pm. Hildebrand (1983)). we would expect the inner parts of the disk to be optically thick even to millimeter radiation. To invert the column density measured at the position of the nul1 points. ive need a larger sample of points through a greater wavelength coverage in polariza- tion maps. More precise maps are also needed: ideally by observing simultaneously the ordinary and extraordinary beams thus removing the problem of the fluctuations of the sky t ransparency. We also presented an S00pm map of HL Tau obtained at the JCMT. The rnap shows an unresolved peak superimposed on weak extended emission. The weak emission is more extended to the South of the star then it is in the East- West direction. The mass of the disk calculated from these observations agrees well with the mass calculated from the polarization data. since the assurnptions made to calculate the mas from the submillimeter flux are very crude. In fact. changing ,3 from the adopted dueof 0.6 to 1.5 Leads to a mass of z 0.13 Mc,. We would like to thank the CFHT and JCMT T.4Cs for the generous al- location of telescope tirne. Thanks also to the Canserv observers who provided us with needed data. This work was supported by the Natural Sciences and - Engineering Research Council of Caaada and the Québec Government. The James Clerk Maxweil Telescope is operated by The Observatories on behalf of the Particle Physics and Astronomy Research Council of the United Kingdom, the Netherlands Organization for Scientific Research, and the National Research Council of Canada. This research has made use of the SIMB.4D database. operated at CDS, Strasbourg. France. with access to the database provided by the CADC. Victoria. BC. Canada. Filter Date X (Frames) Exposure Time (s) Polarization - -- - - 1 23-23 December 1987 3 600 OO. 60'. 120° 1 120 0°, 60". 120° 1 180 O" 800 pm 2-6 August 1994 9 x8 points 300 - December 1994 Completed to 300 - a 9x9 map TABLE 3.1. Log of Observations Filter Resolution ('l) Disk Size (") Ref. TABLE 3.2. Polarimatric disk size at different A. References: 1) Gledhill & Scmott (1989). 2) this work. 3) Weintraub et al 1995. TABLE 3.3. Centroid calculations. A (A) Aperture(") P(%) ISP 8 18 Ref. unfiltered 3700 5895 5895 7.543 1660 1 I J H K K K TABLE 3.3. Linear polarization measurements on HL Tau References: 1- Vrba et al 1976. 2- Ménard k Bastien 199'2. 13- Bastien 198.5, 4- Bastien 1982. 5- Manset. private communication. 6- this work. 7- Hodapp 1984. S- Tamura QG Sato 1989. 9- Grasdalen et al 1984. Pollack et al Henning & Stognienko Preibish et al X Radius (D/2) CC OC OC 1.25 (J) 9- 2.4% 1.11 2-82 1.6 (H) 1.1 3.05 4.69 4-14 2.2 (K) -< O..? 1-13 -5.12 6-75 TABLE 3.5. 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Offset (arcsec) FIGURE 3.2. (b) 5 O -5 R.A Offset (arcsec) FIGURE 3.3. a is an overlay of the vector pattern on a equal polarization angle contour plot of the object. Contours are from O" to 180' with a step of 20". The two null points are clearly identifiable on the map, they act as two illurninating point-like objects for the nebulae. The pattern looks like a interference pattern of plane waves. The inner parallel pattern is identified by the "rectangular" contour between the null points and the effect of ghost images on the CCD is seen in the northwest; b) shows a contour plot of PxI. The nul1 points are also very well determined on this plot by the central strangulation of the contours. R.A. Offset (arcsec) FIGURE 3.3. (b) 1O 5 O -5 -10 RA offset (arcsec) FIGURE 3.4. J. H and K polarization map of HL Tau (a. b a~dc repectively), from Weintraub et al (1995). A quick look at these maps shows the change in the position of nul1 points with wavelength. FIGURE 3.4. (b) 10 5 O -5 -10 RA offset (arcsec) FIGURE 3.4. (c) I L 9 40 2 0 O - 20 -40 RA Offsets (") FIGURE 3.5. Two dimensional histogram of the intersection points. Offsets are from the position of the peak. RA Offsets (") FIGURC 3.6. The 800pm map of HLTau. The first contour is at 38. Following contours are at 30 from each other. The greyed area shows the JCMT beam size at this wavelength. It is clear that emission is extended. FIGURE 3.7. An azimuthal plot of al1 the pixels in the 800pm map. The plain curve shows the best fit gaussian to the data. FIGVRE 3.8. Schematics of the geometrical relation of the three cylindrical coordinates on the light path: a) projected on the disk b) projected in a plane perpendicular to the plane of the sky. - - - a=-0.86rO.13 (Preibisch et al) ; l,,,l.#*llP IL* 1, 1' 1''I 1" III l"1' FIGURE 3.9. Logarithmic plots of the column density versus the radius. The straight lines show the results of a linear regression fit to the data. CHAPITRE 4 Cartes de Polarisation Infrarouge ..\rticle scientifique préparé pour fin de publication Infrared Polarization Maps of some Herbig Ae/Be Stars Roger Haüar' & Pierre Bastien' Star Formation Group and Département de Physique, Université de Montréal and Observatoire du Mont Mégantic C.P. 6 128. Succ. "Centre-Ville". Montréal. Québec, H3C 1J 7. Canada and Daniel Nadeau1 Département de Physique. Université de Montréal and Obsermtoire du Mont Mégantic C.P. 6138. Succ. 'Centre-Villen. Montréal. Québec. H3C 1JT. Canada Received - 'G uest Observer, Canada-France-Hawaii Telescope ABSTRACT We present a set of near IR polarization maps of a number of Herbig ..\e/Be stars (HAEBES) which have published optical polarization maps showing a large polarimetric disk. The sarnple dso contains other types of Young Stellar Objects (YSOs). r\ll stars in the sample are assumed to have inclination angles for the disk of = 90". For al1 the stars in the sample. we observe a change in the polari- metric disk size which we interpret as due to the dependence of optical depth on wavelength. We use the technique presented by Hafiar &- Bastien (1997) to deduce from these observations densi ty information on the large-scale disks observed around YSOs. We find that. to a first order. dl disks show the same density profile. independent of age or type. with p(r) û: r-'e4 to -'-'. 4.1 Introduction Herbig Ae/Be stars were first defined by Herbig (1960) as the higher rnass counterparts of T Tauri stars, with a mass range of 2 ME, 5 Il. 5 lOM,,. The most recent cataiogue of HAEBES is the work of Thé et al (1994) where they review the membership criterias and espand the range of spectral types to early F-type stars from the original A stars used by Herbig. .A recent review of this class of Young Stellar Objects (Herbig 1994) shows that there is no unique set of criterias to select members of the class and that a multi wavelength as well as detailed analysis of the stars should be achieved to be able to really assess the pre-main sequence status of a given star. Catala's review (1993) also points to the difficulties in deciphering their common characteristics. This state of affairs, as shown by Catala (1993). pervades the circumstellar environment of HAEBES. Whereas T Tauri stars show common structures in their circurnstellar material. Herbig stars are still controversial. The presence of a disli is now well demonstrated for the low rnass pre-main sequence objects. Its absence is somewhat the exception that confirms the rule. The situation is completely different for Herbig Ae/Be stars. No clear trend for or against the presence of disks is observed. The interpretation of spectral energy distributions ( SEDs). for exômple. leads to different and equally valid models for the same objects (see Hillenbrand et al 199%. Hartmann et al 1993, .latta et al 1992.1993). It is especially important to study the circum- stellar structures of YSOs since these are important in identifying an object's evolutionaxy status. The sequence of evolution determined by Hillenbrand et al (1992) shows this. They classified the Herbig stars into three groups based on the shape of their SEDs. Modeling showed rnarked differences in the circurnstellar structures needed to explain the energy output. This lead the authors to infer an evolutionaxy scheme based on the importance of the circurnstellar rnaterial since the models showed that it is the main contributor to deviations of SEDs from a pure blackbody. Hillenbrand et ai derive by modeling the following properties for HAEBES: Group II stars. considered the earliest on the track towards the main sequence, possess an accretion disk and are embedded in a large scale envelope; Group I objects, having IR excesses following a power law oc A-:. are believed to have blown their envelope but are still surrounded by an accretion disk: Group III members. showing no or marginal IR excess have dl but completely blown al1 rnaterial around them. 4.1.1 Disks and Pseudo-Disks As pointed out in the previous section. disks are still the exception rather than the rule arouud HAEBES. Albeit this fact. depending on their dynamical characteristics, disks can be classified in two -groupsq. Accretion, centrifugaily supported. disks are known in the literature as, simply, Disks. Whereas. recently, Galli Sc Shu (1993) coined the term, now widely in use, pseudo-disk, to denominate largescale flattened structures. Disks have a size of = 100 AU in radius. Pseude disks cm be larger by an order of magnitude or more, up to a few thousand .\Us. These structures are loosely bound to the forming star. Evidence for both structures is both observationâl and theoretical. The modeling of the SEDs of T Tauri stars provides one evidence for accretion disks, spectroscopic data also point to the presence of these closely bound circumstellar structures (Bertout 1989). Recent high resolution rnillirnetric observdtions have provided a direct mean of resolving the disk (e.g, Lay et al 1994 for HL Tau), as well as diffraction limited observations in the IR. using adaptive optics techniques (Close et al 1997 for HL Tau), where, in K band images the shadow of the structures is visible against the brighter refiection nebula. The only evidences for large highly Battened objects corne from observations, mostly, millimeter interferometry and polarime- tric imaging. Recent calculations by Gdli & Shu (1993) have shown that these could be formed by magnetic support of the material raining down from the large protostellar cocoon on the accretion disk. The largest body of observations of pseudo-disks cornes from the spatially resolved polarization maps of YSOs. Some of these maps show a pattern of parailel vectors in the central parts of the map. Two nul1 points delimit this observed pattern; both are interpreted as indicating the size of the disk. Typically sizes are of z 1000 AUs. Another recent evidence comes from the HST observations of the Orion star forming region where obscured lanes are observed around YSOs: they are interpreted as arising from an edge-on pseudedisk (e-g. O'Dell Sr Wen 1994) 4.1.2 Interpreting Polarization Maps It is now accepted that polarization in reflection nebulae is solely due to scattering (see the review by Bastien 1996). Schematically speaking, one considers multiple scattering in an optically thick disk and single scattering in the optically thin envelope or bipolar outflow. -411 recent models. except Gledhill's calculations ( 1991). use a Monte-Carlo code to generate polarization maps for different cir- curnstellar geornet ries and different density distributions. A concise review of the different models is given by Bastien (1996) and references therein. A discussion on the possibili ty of extracting disk geometrical parameters from these observa- tions is found in Ménard et al (1996). The first attempt at extracting quantitative information from polarimetric imaging has been successfully completed by Bastien & Ménard (1990), where the authors deduce the disk inclination from a comparison of observed maps with theoret ically calculated ones. Their results are in agreement with determinations based on other independent means. Another interesting use of polarization maps is due to Chrysostomou et al (1996a,b). They obtain linear and for the first time circular polarization maps of GSS 30. The authors are able to constrain dust properties from their observations. In this paper, we follow up on another way to extract quantitative information from the observations. We have first proposed the method in another paper (Hadar 6; Bastien 1997. HB). Our work is based on the predictions of the Bastien & Ménard mode1 (BM) (1988. 1990). Based on the fact that the parallel vector pattern in the inner parts is due to multiple scattering and the centrosymrnetric pattern in the reflection nebula to single scattering, BM interpret the nul1 points as transitions between the optically thick and opticdly thin material. The nd points should then have a wavelength dependent position; the polarimet ricdy- deduced size of the disk being smder with increasing wavelength. This allowed us (HB) to devise a method to extract density information from multi wavelength polarization maps. In the following we will show that this effect is observed in a number of objects (al1 of our sample). The next section will present the observations and the reduction and calibra- t ion procedures. Section 3 will be devoted to the results and their analysis for each observed object. Sections 1 and 5 dlbe devoted to the results on disk density and centroid position. two techniques to extract information from the maps that will be discussed in section 6. The last section will summarize our main results and discuss possible implications. 4.2 Observations 4.2.1 The Sample In order to test further the BM prediction, we chose a sample of HAEBES with already published polarization maps in the optical. We wanted objects having a disk seen edge-on or very close to it. One other requirement was that the size of the disk, as seen in the maps, should be of at least a few arcseconds to allow for the detection of the change with wavelength. We did not try for now to address the issue of completeness so that no statisticdly valid conclusions can be drawn from the observations. It is obvious that a complete polarimetric imaging survey of HAEBES in the optical is needed to begin a thorough study of the effect.. Our aim here is to add a few more blocks to our previous paper in order to show that the predictions of the BM mode1 is confinned for different types of stars. Table 4.1 gives some of the known parameters for the selected objects. The resolution and wavelength of the polarization maps are also provided dong with references. It includes the maps obtained for this work. Column 3 gives the binning used to calculate each vector on the map in parenthesis only if it is not equai to the resolution. 4.2.2 Observations The infrared observations were carried during two runs at the Canada- France-Hawaii telescope. in August 1995 and in June 1996. The observations were obtained with Monica ( MOKtreal Infrared CAmera), which is based on a 2.56 x 2.56 HgCdTe infrared array (Sadeau et al 1994) and a polarization module made of a hot rotating half-wave plate and a cold polarizer. Mueller matrix calculations for the setup and the observational procedure described in Hajjar et al (1997). The pixels have a size of 0.238" at the f/8 focus. We have taken sky frames dong with object images for al1 our extended objects. For point-like polarimetric standards. polarized and unpolarized, sky frames were obtained by calculating the median of dl masked object frames. Exposure times were chosen such as to not saturate the individual pixels. We also used in Our second run. three different positions on the chip, allowing for a better correction of persistence problems for very bright sources. Because of different instrumental problems we were unable to have both J and K polarization maps for al1 our objects. Column 5 of Table 4.1 gives the passbands for which we were able to complete the maps. Some of the objects observed during these runs have been included in another paper (Hajjar et al 1997). 4.2.3 Reduction and Calibration The object images were sky subtracted individudy. using, whenever possible. a sky image constructed from the closest sky frarnes taken. They were then flatfield corrected by dome Bats. Care was taken to account for the telescope and dome emission by obtaining lights-on and Lights-off exposures. These were subtracted to get a flat field image. Finally, before shifting and adding. each exposure aas corrected for bad pixels. To calibrate Our zero point and check for instrumental polarization. we observed on the first run CRL 2688, and HD 155197. as polarized standards. and BD +32" 3739 as an unpolarized star. The instrumental polarization was found to be nu11 within the errors. HD 155197 was unreliable due to its brightness which sometimes saturated the chip. This was also the case for our second run. We had also used HD 161056 which also happened to suffer from the same problem. This along with weak polarization values in the infrared made them unreliable for calibration purposes. The first run aas successfully calibrated with CRL 2688. We assumed a position angle of 1%' in a 25'' aperture and a polarization intensity equal to the one we measured on Our data. The second nin Ras calibrated using Par 21. since it was observed on both runs albeit in different wave bands. E; and J for the first and second run respectively. We assumed a constant polarization angle. the polarization being equal to what we measure by centering an 8" aperture on the star. 4.3 Results for Individual Objects For each of the program sources we will briefly review its known properties. We will present and discuss our results and infer implications on what is known. 4 out of 5 objects are HAEBES. RN0 138 is not visible in the opticd and, as we will show remains deeply embedded in the infrared. It also shows a large polarimetric disk in the map of Draper et al (1985b). Data on disk size with wavelength are gathered. for al1 objects in Table 4.2 where column 2 gives the wavelength. column 3 the disk size estimate and column 4 the column density at the nul1 points as determined from the Stognienko & Henning (1996) dust model. Centroid calculaiion results are al1 given in Table 4.3. 4.3.1 Parsamian 21 This star is given a spectral type .43 in the catalogue of Thé et al (1994). based on a spect rophotometric classification by Dibai ( 1969) and Neckel8; Staude (1981). But more recent work by Staude 8~ Neckel(1992) points out to a drastically different spectral type. They find from spectral data that Par 21 shows a11 the characteristics of FU Ori activity and is a F5 Iab star. They also suggest the presence of an .4 type cornpanion to explain the broad Ha emission. Their estimate of the systemic velocity yields a distance of 1.8 kpc. a radical change from the previous distance of 0.1 kpc. The. also report on the detection of a Herbig-Haro outflow. with three detected knots. They also find that the luminosity peak in the nebula is not stellar implying that an important fraction of its light is due to scattering. It agrees well with infrared images by Eiroa gi Hodapp (1990). These authors also find no noticeable peak displacement with wavelength. although they do not give any indication on the way relative positions where measured. Hillenbrand et al (1992) classify Par 21 as a group II object, being associated with an accretion disk and a substantial dust envelope to account for its rising SED in the infrared. In the case of an FU Ori star, a disk is definitely present. The large scale pola- rimetric disk observed in the Draper et al (1985a)map was assigned an inclination angle i = 80 - 85" by Bastien & Ménard (1990). In view of al1 these facts, the best mode1 for Par 21 is still the one proposed by Elskser & Staude (1978) with a few minor refinements. Figures 4.1 and 4.2 show our J and Ii band maps for Par 21. A rapid exami- nation of the maps clearly show the change in the size of the disk wi th wavelength. The disk size at J carries a larger uncertainty than in the K band. This is due to the difficulty in determining the exact location of the western nul1 point. We measured no peak displacement with wavelength in the images. We should note that. due to the obserkational procedure. the usable field for such a determination contained too few stars to make for a reliable astrometric measurement . But. an examination of the centroid calculations in Table 4.3 shows that. on average. the stars location shows a larger offset from the peak in the J band than in the K band. -4 difference of = 2" in declination offsets is found. -4lthough both positions are compatible within the errors, for reasonable values of the rejection criteria, it is interesting to note that this difference is about the same for al1 the calculations. This would agree well with the fact that the peak seems extended. showing a larger FWHM than other stars in the field. This indicates that the peak is due to scattering of starlight by circumstellar dust rather than a direct viea to the star itself. This star is assigned a Aïea spectral type by Thé et a1 (1994). It is located at the center of a reflection nebula showing a bipolar structure, as delineated by four spikes of emission extending from the central object (e.g, CCD images in Aspin et al 1985). No molecular outflow has ben detected at its position. Based on the shape of its SED in the far IR, Hillenbrand et al (1992) have clâssified this HAEBES as a Group II object. Leinert et al (1993) have made use of speckle interferometry in the near IR to obtain information on the most inner parts of the nebula. They hdthat the bipolar structure extends to within 0.13" and. from polarimetric data, probably inside the inner 100 AU. They aiso show that a large fraction of the peak is in fact scattered light in an extended halo. so that the observed stellar peak in normal optical seeing conditions is in fact non-stellar. They conclude that the previously measured extinction. Ar- = 2.6 (Calvet & Cohen 1978) is largely underestimated (up to a factor of 10 to 20). The star shows a very high degree of linear polarization. P= 10.8% in the optical (Hillenbrand et al 1992). Based on the V band polarimetric map by Aspin et al (1985). BM find that the disk around this HAEBES is seen edge-on. The results of Aspin et al demonstrate the presence of a disk or a dust torus around the star. They also find that the Yorth-Eastern lobe is much brighter than the SW part of the bipolar structure; a 1 mag difference. The degrees of polarization observed in the V band. in the extended nebulosity. tend to agree well with the disk inclination deduced by BM. The value of the extinction deduced by Leinert et al shows the same situation as for Par '21 (Staude 8. Neckel 1992) and HL Tau (Weintraub et al 1995 and Beckwith S; Birk 1995). arnong possibly other pre-main sequence objects. Al1 of the above agrees well with the presence of an optically thick disk. seen edge-on. in the environment of the HAEBES. Figure 4.3 shows the K band polarimetric image. It is identical to Figures 4.1 and 4.2. -4s is the case for al1 published maps, the visual determination of the disk size is not always easy. The estimate for the V band rnap of LkHa 233 suffers from this problem. Where the usuai error on the nul1 points is equal at least to the resolution of the map. it should be viewed as larger in the case of an already published map. The errors, however, do not make the two sizes, at V and K7fit within one a. 4.3.3 V376 Cas The region containing V376 Cas and V633 Cas is cornplex. It contains a millimeter protostar. LkHa 198 C (Haüar S: Bastien 1994. 199Cb). V633 Cas. also known as LkHa 198, has an infrared companion 6" to the North of its position (Lagage et al 1993). A large scale CO outflow has been detected in the region ( Levreault 1988). The source of the outflow has long been the subject of debate. attributed sometimes to LkHa 198 and sometimes to V376 Cas. Strong arguments point now to the millimeter object as the source of the outflow (Haüar &. Bastien 1997b). V376 Cas was found CO have the highest koown linear polarization among YSOs. = 23% in the red (Asselin et al 1996) and, from imaging polarimetry, is probably surrounded by a disk seen edge-on. Imaging data in the optical. up to the I band. show that V376 Cas is extended with a position angle of 135". almost perpendicular to the direction of the disk, as revealed by the parallel pattern in the maps. Speckle interferometry and polarimetry by Leinert et al (1991) have shown that this structure extends to the inner parts (inside the inner 0.2") of the nebula. Scattered light makes an important contribution to the peak of the nebula. which is then probably non-stellar. A number of Herbig-Haro objecte have been detected in the vicinity of the HAEBES by Corcoran et a1 (1995), although their relationship with the star has yet to be demonstrated. One other important feature in the optical images of V3ï6 Cas is the sickle-shaped nebula to the West of the star. Our data show that it is still prominent in the J band (Figure 4.6) but does not appear in the K band (Figure 1.1) image. Asselin et al (1996) find that their K band image shows a slight elongation of the isocontours at a position angle of BO0. Our image shows that the outer contours are rather extended in the NESW direction. It seems that the da& region separating the nebular material in the immediate surroundings of V376 Cas and the sickle-shaped profile. brightens up in the IR. longward of J. Unfortunately, we have no H image to be able to check for that transition. In the inner parts of the nebula, the elongated profiles in the optical (V to 1) are no longer seen in K, instead, as seen in Asselin et al the contours show an almost circular shape. The transition can be readily identified in the J image where the inner contours show a triangular shape. One possible due to the understanding of these observations cornes from our K band polarimetric image (Figure 1.1). -4 number of polarization images are already published in the literature. Asselin et a1 (1996) have observed it in the I and V bands: their published maps are of various resolutions: 1'' and 1" for the 1 images, 1.5" and 4'' for the V band maps. Another 1 band image has been studied by Piirola et al (1992), with a resolution of O.?". The last one is in a paper by Leinert et al (1991). .4s pointed out by Asselin et al. their measurements do not always agree with Leinert et al. They attribute some of the changes to a probable wiability in the illumination of the nebula. A close inspection of the V and 1 maps obtained by Asselin et al show important differences in the vector patterns. Whereas in the I band image one could easily pinpoint the nul1 points at two opposite directions from the peak. the V band image shows a characteristically different vector profile near the peak. The nul1 points are not at al1 obvious in this image. The rneasured peak displacement. FZ 0.6" between V and 1. do not explain the apparent discrepancies. Furthermore. a look at the disk sizes measured. clearly show the importance of resolution on determining the position of the nul1 points. Another important discrepancy is found between both the Asselin et al and Leinert et al maps on one side and Piirola et al on the other. This is probably due to the much higher resolution of the last map. A match between the contour image and polarization map of Piirola et al show that the peak is at the edge of what con be considered the parallel vector pattern, whereas in Asselin et al and Leinert et al maps, the peak is at the center of this region. They also speak only of one nul1 point, observed at 0.8" to the S W of the peak in the 1 band, and 1.2" in the V band. They seem to assume that nul1 points should be symmetric relative to the peak, so that the disk sizes measured, 1.6" in 1, and 2.4" in V, would be in sharp disagreement with Asselin et al. The location of the peak wanants a revision of these estimates. A casual examination of the Piirola et al 1 map indicate a disk about 3" in size. Our K band polarimetric image (Figure 4.4a, b and c) has a resolution of 0.75" x 0.75". A higher resolution map (0.5" x 0.5") is not different apart from the number of vectors drawn. The centrosymmetric pattern is clearly identifiable in the outer parts of the nebula and up to a small distance from the peak. In the inner 10'' . the map is noticeably different from al1 other optical rnaps. The peak falls at the northern edge of a parallel vector pattern oriented NS. Although it appears to be the illuminating source of a part of the reflection nebula. it is evident that some parts of the rnap (mostly the SW and deviations in the SE and NW profiles) tend to show that the illuminating source fails on the parallel pattern described before. This looks Like the situation for LkHa 231 where in the 1 band. a disk appears to fa11 at the position of the star but with deviations in the ceutrosymmetric pattern (Hajjar et al 1997); whereas in the K band. it is clear that another deeply embedded source is illuminating major parts of the nebula. Imaging at longer wavelengths should give the definitive answer on V376 Cas. 4.3.4 V633 Cas V633 Cas is an extensively studied HAEBES. Located 37'' to the south of V376 Cas, it is close to the apex of an elliptically shaped nebula. It is most probably the source of a jet (Corcoran et al 1995). Its infrared companion, V633 Cas B, has been detected at 10pm as well as in the K band. A close look at the Piirola et al rnap shows the presence of the infrared companion in the I band. It is also related to feature E, in Asselin et al nomenclature, which is coincident with a polarization nul1 point in the otherwise centrosymmetric pattern of the 1 band map. The star is also the source of a jet (Corcoran et al 1995). According to integrated linear polarimetry data by Leinert et al (1991), V633 Cas shows a sharp change in the polarization position angle, in the infrared. Based on their poiari- zation map. Piirola et al (fig. 4 of their paper) consider that the parallel vectors they see in the inner 0.5'' is due to an unresolved disk of about 0.2'' . It is not clear if this is real or due to other observational effects; for example. a saturated peak in one frame would possibly give such an effect. In fact. an examination of the surrounding contours would make us expect an orientation of the paralle1 profile at 90' to the one seen. But it would also be the first signs of the change in polarization angle. -4s pointed out by these authors, the integated polariza- tion observed depends on the geometry of the close environment and the optical thickness of the disk. We estirnate from their map a disk size of about 3" with nul1 points at equal distances from the peak, and a position angle 5 70'. This agrees with an estimate made on their lower resolution map ( fig. 1b of Piirola et al 1992). Our K band polarimetric image (Figure 4.7) clearly shows an elliptical pattern of vectors. The null point is also clearly present in the map and coincides with a very small cometary nebula. pointing to the W. associated with the IR source. The polarization vectors deviate from centrosyrnmetry on dl the extent of the cometary shape. The null point is most probably the result of the addition of two perpendicular polarizations. One due to the scattering of light coming from the IR source by the cometary nebula. the other from the illumination due to V633 Cas. From a casual inspection of the image. we notice that, at the position of the null point, the intensity of the nebula is two times that of its surroundings. The contribution of the IR source is about equal to that of V633 Cas so that polarizations are equal in percentage. Vectors on the same radial direction, and close to the cometary nebula. give a polarization of 4.5% and an angle of 140". So that the polarization of light from the IR source by the materid in the cometary nebula is about 1.2% with 0 50°. This agrees very well with the polarizations found on or near the peak where, naturally, it is the contribution from the IR star that dominates the luminosity. If an obscuring disk is present at the position of the star. the measured polarization on the peak should be lower than in the nebula. exactly what is seen in our map. No disk is meaurable around LkHa 198. The general profile is elliptical and the inner vectors. although distorted. are more reminiscent of an imperfect ellipse rather than the usual parallel pattern expected from the disk. We put an upper limit on the size of a disk of about 0.5". the resolution of our map. The peak of emission for the HAEBES fa11 at the center of the elliptical and centrosymmetric patterns. RN0 138 is located in the southern part of NGC 7129. NGC 7129 contains two known HAEBES, LkHa 234 and BD +65" 1637; these two stars and some other ones are studied in detail by Hajjar et al (1997). RN0 138 is a srnall patch of nebulosity showing a bipolar structure with a recently identified central stellar peak (Miranda et al 1994). Draper et a1 (1985b) obtained the first polarization map of RN0 138. They observed no stellar peak against the bipolar nebula. Based on their results, they conclude to the presence of an optically thick dust torus around the illuminating star blocking it form direct observation. They found that the south eastern part of the nebula is brighter than the NW lobe. This lead thern to assume that the nebula is tilted towards the observer. More recent observations (Miranda et al 1994) point to an outburst between 1988 and 1993 of the illurninating star. They observe in their 1993 image a stellar peak almost at the center of the bipolar nebula. The peak was not present in 1988. They also obtain the first spectrum of the peak. Some characteristics of an FUori type outburst are present, but the absence of other signs, like double-peaked emission lines. makes its classification still uncertain. The Draper et a1 (1985b)map shows a centrosymmetric pattern in the outer parts of the map and a parallel pattern in the imer parts. Two null points are clearly present. We estimate a disk size of 10" at their effective wavelength. The map has a resolution of -2.4". The size estirnate should then be taken as an upper limit to the disk size at the wavelength of observations. We obtained a polarized J image (Figure 4.8) and an H and h: intensity images (Figure 4.9). Our Ii band map has a too low signal-tenoise to be of any use. Only a few vectors are seen on the map. They confirm the J data but no serious analysis can be performed. The polarimetric map has a resolution of 1". We can clearly identify a centrosymmetric profile. and a paraliel vector pattern. The position of null points can be estimated by extrapolating the profiles seen. The bright peak observed is also extended. We rneasured its position relative to SL'S 6 in the three IR bands. The peak shows a clear displacement to the East with longer wavelength (Figure 4.10). This is identical to what has been observed in other YSOs (e-g.. V376Cas. Asselin et al 1996). This points out to the fact that this is probably not the star but is probably due to scattering in the circumstelIrtr environment. It is also an evidence for the existence of a dense optically thick disk seen near to edge-on if we rely on the K band image. The disk is substantially denser then Say the HL Tau disk (HB), since its K size (in arcseconds) is much larger than that of the T Tauri star. RN0 138 is most probably younger than HL Tau. Can we then make use of measured disk densities to have an age estimate of the YSOs studied? 4.4 Disk Density Structure Table 4.2 gives size determinations and densities for al1 the objects in Our sarnple. Null points were chosen by searching for the two lowest polarization values on or near a line passing through the peak and parallel to the direction of the integrated lineax polarization. We have added data for Noor from Hajar et al (1997). md HL Tau from HB. The most 'useful" data in Our sarnple are for HL Tau. It is the only object for which we have 5 different size measurements. The dust opacities used corne from Henning & Stognienko (1996) which includes coagulation of the basic dust grains and ice mantles. They calculated the opacities with their grain model as well as with Pollack's model (1994). For the sake of cornparison. and to give an appreciation of the effect of the grain mode1 on the results. we used the calculated opacities for Pollack's grains and for the iron-rich model of Henning &L Stognienko as detaiied in Figure 5 of their paper. It is clear from Our data. that the change in the polarimetric disk size is in fact a basic mechanism. independent. at least. on the YSOs type. be it a TTS. a FC Ori or an HAEBES. This shows that it re1at.e~to basic circumstellar structures present in al1 forming objects. Following on this idea. Figure 4.1 1 shows a logarithmic plot of the colurnn density versus the scaled disk diameter for ail the objects of Our sarnple. and HL Tau, and Noor. We chose two scaling diameters. the I band and K band sizes. Both suffer from a number of shortcomings. Measurements at 8500 A were considered to represent fairly closely the 1 band; except for HL Tau, dl size estimates at this wavelength are of poor quality because of the resolution achieved, they are al1 based on already published polarization maps. In the K band, aithough more homogeneous since most of the data was obtained by us, some of the sizes are just upper lirnits. Figure 4.11 shows that the column density profile can be approximated by a power law. The enor on the slope is larger than what is given on the figure since we did not take into account the uncertainties on the position of the nul1 points. Assurning that oc(r)x $, inverting the Abel1 equation gives p(r ) x r.1 The values found for a result in a density profile p(r) a r-'.', in the case of Pollack's (1994) dust mode1 and p(r) a r-'" for the He~inget al (1996) rnodel. This agrees well with the commonly used profile for the protosolar disk. p(r) x r-l.5 4.5 Calculating the Centroid Position We are interested in this section to cmy further the method developed in Hajjar et al (1997) for the determination of the center of illumination of nebulae from polarization maps. It makes use of the fact that in a cent rosymmet ric pattern. al1 perpendiculars to the polarization vectors pass by the illuminating source. We thus determine the intersection points of al1 couples of the perpendiculars. The rnean of this distribution should point to the real position of the star. Looking at the set of polarization maps. the so called centrosymmetric pattern is in fact elliptical. this is clearly seen in maps of Par 21, V633 Cas or HL Tau (HB) for example. For the sake of cornparison we calculated the distribution of the intersection points of the perpendiculars to an ellipse. The result can be seen in Figures 4.12 and 1.13 where the first figure shows contour plots of the density of points in 4 x 4 binning boxes, in the arbitrary units of the graphies. The major axis of the ellipse is pardlel to the abscissa mis, and for two different ellipses:4.12a) with a minor to major axis ratio of 0.7, and 1.12b) with a ratio of 0.2. The second figure is a plot of the points in the sarne arbitrary units. One notes that the central position (O, O) is almost devoid of points. The distribution, as expected, is symmetric relative to both the x and y axes. The highest densities show a losange shape around the center of the ellipse. We used 100 equally spaced points sampling the major axis of the ellipse, leading to 200 distinct points on the ellipse. The distributions shown in Figure 4.12 are to be compared to the distribution found for our objects (Figure 4.11). For most of our objects. the two- dimensional histogram show the general patterns of the ellipse's histogram. The inner parts are more confused. RN0 138 and LkHa 233 show major deviations from that general profile. For each of the observed maps, we calculated a weighted and a non-weighted rnean using the formalism of Hajjar et al !1997). We performed the calculations for different rejection thresholds. determined in units of the standard deviation O (n x O). The aim is to look at the behavior of the calculated mean with the value of n. This pmvides important information on the kind of distribution of intersection points found. as can be seen from Table 1.3. Generdly. below n= 1.8. the number of points left drops significantly and the mean shifts towards different values then what is obtained for the higher rejection thresholds. For n 5 1.5. al1 points are rejected. It is easily noted that al1 calculated means fall at the center of the distri- butions seen in figure 4.11. For small values of the rejection criteria. the mean position shifts towards the peak of the distribution which does not necessarily indicate the location of the source since for an ellipse. the central position is marked by a depression in the 2D histogram (figure 4.12). It is to note that these calculations are particularly sensitive to the observational procedure and the accuracy of the calibration. An error of just a few degrees in the position of the zero point of polarization angles will introduce a systematic error which is impossible to evaluate. 4.6 Discussion 4.6.1 The Disk The following discussion should be considered carefuily and is just an in- dication about possible implications of the observations. The number of points available for each object and the uncertainties, which are large. especially on nul1 points determination from published maps, make our conclusions a first attempt at extracting information from the available dataset. The wavelength of 8500 for RN0 138 and LkHcr 233. as well as the rnap at 7500 A of HL Tau. are in fact the peak response wavelengths of the instrumental setup. Al1 these maps were obtained without a filter by the respective authors. The effective wavelengths of the maps may be quite different since they are affected by the response curve of the instrument. the emitted continuum spectra in the sensitivity region of the instrument and on the P(X) curve at the position of the measurement. The effective wavelength could thus be a function of the spatial position on the map since there are good reasons to believe that the grains have different properties in the diffuse extended envelope (outfiow) than in the disk. Another possible source of error. as already pointed out. arises from the way we measured disk sizes on published maps. This makes for an uncertainty on nuIl pointe larger than the resolution achieved, which. at least. from the previously mentioned maps. is quite large. In the column density fitting of section 1, we assumed a priori that the profiie is independent of age and type of the star. The initial results seem to justify our initial hypothesis. It is tempting to note that evolutionary models of the protosolar nebulae by Ruden & Lin (1986) show very smdl changes with time on the density profile at rnidplane of the disk. Following this line of argument, we will compare the calculated profile with models and observations already available in the Literature. We note from the start that, almost ail modeling and calculations on disks have considered sizes of 100 AUs or smdler. unresoived at the distance of al1 of our objects except for HL Tau. This is the typical size of a centrifugally supported accretion di sk. Two different modeling paths can be found in the literature, at least for what is pertaining to protestellar environnements. Protosolar models are mainly concerned by *reversing9 the history of the solar system in order to gain knowledge of the original nebula. Obviously, the current conditions of the sun and planets make up for the boundary conditions. Models of disks around YSOs aim at reproducing the SEDs found for different types of stars. T Tauri. FU Ori ...etc ... they also t ry to predict the outcome of stellar format ion and the final masses of the stars nhen they reach the main sequence. Calculating accretion rates and modeling the viscous transport of angular momentum are some of the major concerns. .\dams et al ( 1988) developed what is now considered a standard model of accretion disks around YSOs. They use a geornetrically thin disk with a power Iaw surface density C(r)x r-p with p = 1.5 for a Bat spectrum source. Xoting that around this due. the SED is not sensitive to the surface density distribution. they use. based on stability arguments. a value of p = 1.75: this is a value also used by others in modeling millimeter and sub-millimeter observations (e.g. Adams et al 1990). Beckwith et a1 (1990) consider a disk with p = 1.5 to fit their observa- tions. Other values of the power law are also used. Berrilli et a1 (1992) compare near to mid IR observations of HAEBES with model calculations for an index p = 0.5 while Henning et al consider a constant density model (p = O) to fit their 1.3 mm observations of YSOs in Charneleon. FU Orionis type stors are modeled by Weintraub et al (1991) with a range of 1 p 5 1.75. Yamashita et al (1990) have modeled sub-millimeter continuum observations of GGD 27 IRS with a power law dust distribution in a disk geometry. In their case, they do consider large scale structures as opposed to the previous examples of 100 .4C accretion disks. They find that the best fits give a range of 1.8 < p 5 2.2. These results compare well with those of obtained on theoretical and observational studies of the protosolar nebula and the wlar system. Weidenschilling (1977) show that the present planetary mass distribution is cc r-i, a result also dernonstrated by Hayashi (1981). In their review of solar nebula models, Cassen k Surnmers (1983) also show that for weak viscous transport of ongular momentum the solar nebula reaches a power law density distribution with an index zz - f . And. as we pointed out previously, Ruden &- Lin ( 1986) find little change in the mid-plane density profile with evolution: a profile well approximated by a power law with an index =y --3 2 ' It is interesting to note that our result for the density profile agrees ver' well with these theoretical and observational results with the advaatage of being one of the few "direct- ways of measuring the disk density. The density profile of the inner. centrifugally sustained disk. extends to the outer parts of the nebula. This is somewhat expected if one considers that the Battened structure is the result of an infall towards the accretion disk. In fact Shu (1977) calculating the dynarnic collapse of a sphericd distribution of matter have shown that the distribution of free infalling matter towards a centrally forming stellar core has a power law index of -2. and outflowing material will lead to an index of -2. A reasonable assumption would then be to expect some kind of middle ground since recent studies of the disk of HL Tau (Cabrit 1996) indicate that the large scale disk shows mixed signs of infalling and outflowing motions. Our a priori assumption of the independence of the density law on type or age is strengthened by the results seen in Figure 4.11 as well as by theoretical results. One major concern is obviously the completeness of the sample. To further test those results and conhor infirm some of these conclusions, on needs to look more closely at the effects of dust rnodel. It is also important to see if observations of disks at different inclination angles still agree with Our findings. 4.6.2 The Star's Position Table 4.3 lists al1 centroid calculations for d our polarization maps. Figure 4.14 shows the density distributions of intersection points. As we already pointed out. al1 calculated means are located near or at the "peak' of the distributions. Except that results for different values of the rejection threshold n. and for a puticular object ,generally show two "major positions': one for n 2 1.9 and one for smaller values. This can easily be understood in terms of the expected distribution of intersection points for an ellipse (Figures 1.13 and 4.13). The highest densities being located on a losange shape with a hole. the mean value will fa11 suitably near the real position of star. When n gets smaller. the selection box will not include al1 of the high density region and the mean will be more sensitive to local Auct uations in the distribution of intersection points. This is understandable by looking at some of the actual distributions of Figure 4.11. The polarization maps clearly show deviations from the purely elliptical profiles. Furthermore. the presence of a binary illuminating the nebula. as is possibiy the case in V376 Cas. will greatly affect the resulting distribution. The number of amilable vectors is also an important factor in the precision of the results as well as the distribution of position angles. In RN0 138, for example, al1 the vectors have a very narrow range of position angles, this makes for very large errors as shown in Hajar et al (1997). Calibration uncertainties will introduce a systematic error in the coordinates of the intersection points that is impossible to evaluate. This being said a few words should be said on the values of a calculated for the means. These nurnbers should be viewed more like a measure of the spread of points in the distribution than an absolute error on the determination of the star's position. This is clearly due to the fact that the polarization vectors have an elliptical rather than a circular pattern. The calculated cr is then a sum of the actual error on the position of the illuminator. due to the factors described above. and a measure of the spread of points in the distribution which depend on the ellipticity of the polarization patterns. 4.7 Conclusion We have presented infrared polarization maps of some Herbig .-Ze/Be Stars and other YSOs. A published polarization map in the optical and a di& seen nearly edge-on were the main selection criterias. We measured a displacement of the observed nul1 points as a function of wavelength in the maps. as predicted by the BM model. Interpreted as transitional points between the optically thick (multiple scattering) region of the disk and the optically thin (single scattering) region of the large scale envelope (or bipolar outflow). we were able to obtain information on the density structure of the disk. king the opacities calculated for two grain models which incliide coagulation (Henning et al 1996). we find that the density distribution in the disks follows a power law p(r) = r-'-' t" -L-4. very close to profiles currently used in modeling accretion disks around YSOs. This profile also seems independent of the YSO's type (T Tauri. HAEBES. Fr Ori). and not dependent on the age of the object. This study gives further evidence for the presence of disks around H4EBES. We also analyzed the method described in Haüar et a1 (1997) to obtain the position of the star from polarization maps. In general, the density distribution of intersection points follows that obtained for elliptical polarization patterns, with major differences in the inner parts of the distribution. Optimal rejection criterias for the cdculation of the mean are between 2.0 and 3.0 o,provided that the polari- zation pattern is symmetric leading to a symmetric two-dimnesiond distribution of intersection points. The major source of error cornes from the observational procedure and the calibration. The solution lies in the observation of both the ordinary and extrmrdinary bearns simultaneously. This will dow for a better correction of the sky and will overcome the problem of sky transmission. hnother improvement lies in a broader coverage in wavelength of a larger number of objects, with different disk inclination angles. This will help obtain statisticdly valid results. We would like to thank al1 telescope operators that we have met during Our runs. They have helped make it a smooth experience. This work was supported by the Natural Sciences and Engineering Research Council of Canada and the Québec Government. This work has dso made use of the SIMBAD database operated at CDS. Strasbourg, France. We would like to thank the CADC for providing access to SIMBAD for canadian astronomers. Kame Filter Resolution (bins)(") Disk Ref. D(pc) Ref. (ofA) Inclination (O ) V633 Cas 1'376 Cas Par 21 RYO 135 LkHû '233 TABLE 4.1. Program stars and related parameters. Inclination angles are from Bastien & Menard (1990) except for V376 Cas (Asselin et al 1996). References: 1 ) Asselin et al (1996), 2) Piirola et al (1992), 3) Leinert et al (1991), 1) This work, 5) Draper et al (1985a), 6) Draper et al (1985b), 7) Aspin et al (1985). 8) Racine (1968), 9) Neckel & Staude (1984), 10) Staude & Necke1(1992), 11) Calvet & Cohen (1978). Pollack et al (1994) Henning et a1 (1996) Yame Filter (or A) Size (") CC OC g cm-* - - V633 Cas 1 K V316 Cas V 1 HL Tau 7500 .;\ 1 J H K Par 21 8500 .i J E; RSO 1:3S 8.500 .A J Noor (SVS 13) 1 J LkHa 133 V Ii TABLE 3.2. Disk sizes and density calculations for ail Our program stars. HL Tau and SVS 13 are from Hajjar & Bastien (1997) and Hajjar et al (1997) respectively Xame X (pn) NtOtd * Weighted (n x o)l Nleft ' ~a(a~)(")~A6(0.5)(")~ - 1.5 O \'6:33 Cas 2.2 -58.31 1 + 3.0 35764 12(4.1) lJ(3.5) - 3.0 49.840 0.7(3.4) 1.1(44) + 1.0 25.853 1.5(2,9) '2.8(2. l ) - 2.0 4 i,343 0.6(3.2) 1.1(3.5) + 1.9 20.230 1.9(2.*5) 3.1(1.1) - 1.9 39,386 OA(3.1 ) 1.1(3.4) TABLE 4.3. Calculation resuits for centroid determination. 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ApJ. 400, 556 Thé, P.S., de Winter, D. & Pérez, M.R. A&AS, 104, 315 Weidenschilling, S.J. 1977, Ap&SS, 51, 153 Weintraub. D.A., Kastner, J.H. & Whitney, B.A. 1995, ApJ, 452. 1-11 Weintraub, D.4., Sandell, G. & Duncan, W.D. 1991, ApJ, 382, 270 Yamashita, T.,Shuji. S.. Kaifu, N. & Hoyashi, S.S. 1990, ApJ, 365, 615 R.A Offset (arcsec) FIGVRE 1.1. J band polarimetric map of Parsamian 21. a) shows the vector patterns. where the vector length is proportional to the polarization in percentage and its direction, the position angle of linear polarization. b) and c) overlay the vector pattern on intensity isocontours and equal angle contours respectively. The equal angle contours provide an elegant way to locate the position of the nuU points. RA. Offset (arcsec) FIGURE 4.1. (b) A.d--T-- TT- -. - 10 l I IO 5 O -5 - 10 R.A Offset (arcsec) FIGURE 4.1. (c) i 50 X I 1, 4 5 O -5 R.A Offset (arcsec) FIGURE 4.2. K band polarimetric map of Parsamian 21. The three figures are identicai to figure 1.1. RA Offset (arcsec) FIGURE 4.2. (b) 5 O -5 R.A. Offset (arcsec) FIGURE 4.2. (c) R.A. Offset (arcsec) FIGURE 1.3. Sarne as figure 4.1 but for the K band map of LkHa 233. -6 - 1 i -.- 6 4 2 O -2 -4 -6 RA. Offset (arcsec) FIGURE 4.3. (b) - C 1 1 6 4 2 O -2 -4 -6 R.A Offset (arcsec) FIGURE 4.3. (c) \\k - -\- -20 - ----. ;<. 50 X -I [I. -\\-,- 1O O - 10 - 20 R.A Offset (arcsec) FIGURE 4.1. Same as figure 1.1, for the K band map of V376 Cas. i O O - 10 - 20 R A Offset (arcsec) FIGURE 1.4. (b) 10 O - 10 - 20 R.A. Offset (arcsec) FIGURE 4.5. (c) 1O 5 O -5 -IO R.A Offset (arcsec) FIGURE 1.6. Contour plot of the direct image in the J band of V376 Cas. The sickle-shaped nebula is clearly seen whereas it is not distinguished in figure 1.4. 5 O -5 RA Offset (arcsec) FIGURE 4.7. Same as figure 4.1, for the K band map of V633 Cas. The low polarization region within the centrosyrnmetric pattern is at the position of the cornetary nebula illuminated by the IR source (Lagage et al 1993). as seen in b). 5 O -5 R.A Offset (arcsec) FIGURE 4.7. (b) 1 1 5 O -5 R.A. Offset (arcsec) FIGURE 4.7. (c) 5 O -5 R.A Offset (arcsec) FIGURE 4.8. Same as figure 4.1. for the J band map of RN0 138, a bipolar nebula in NGC 7129. 5 O -5 R.A. Offset (arcsec) FIGURE 4.8. (b) R.A Offset (arcsec) FIGURE 4.8. (c) FIGURE 1.9. Contour plots of the direct images in H (a) and h: (b) of RN0 138. FIGURE 1.10. Peak displacement with wavelength of RN0 138, offsets are relative to the position of SVS 6, assumed to be independent of A. O.? FIGURE 4.1 1. A logarithmic plot of the column density versus the scaled disk diameter for the sources given in Table 4.2. DI is the disk diameter in or very near the 1 band, and DK its size at 2.2pm. The numbers give the calculated dope and its error. FIGURE 1.12. 2D histogam showing the density of intersection points of perpendiculars to a pure ellipse: a) for an axis ratio of 0.7. and b) for an axis ratio of 0.2. Units on the graphic are arbitrary. -200 O 200 FIGURE 1.12. (b) FIGURE 4.13. Plot of the intersection points of perpendiculars to an ellipse with an axis ratio of 0.7. 10 O -10 40 20 O -20 -40 LkHa 233 (K) RN0 138 (J) 20 10 O - t0 -20 V376 Cas (K) 50 O -50 20 O -20 Par 21 (J) Par 21 (K) RA Offsets (arcsec) FIGURE 4.14. Density distribution of intersection points for al1 our polarization maps. CONCLUSION Xous nous sommes fixés certains objectifs au début de ce travail. Certains ont été atteints. d'autres ne l'ont été que partiellement. Passons avant tout en revue les principaux résultats de ce travail awtde les mettre en perspective des objectifs fixés et d'analyser leur potentiel vis-à-vis des ouvertures et possibilités de recherche qu'ils offrent. Xous avons présenté trois cartes millimétriques et submillimét rique de trois objets différents: deux E-AEBEH, V633 Cas et V628 Cas. et une T Tauri. HL Tau. Les deux premières cartes ont permis de montrer. du moins pour les objets concernés. les difficultés de l'interprétation des DES dëtoiles jeunes. Cne nouvelle étoile en formation a été détecté près de V633 Cas. et tout indique qu'elle est reponsable d'une grande partie de l'émission dans I'IRL de la région. Toutes les évidences accumulées montrent également qu'elle est à l'origine du flot bipolaire de la région. La carte de MWC 1080 illustre les difficultés d'attribution d'un flux millimétrique à une étoile en particulier. L'étendue de la carte ainsi que sa morphologie particulière complique la détermination des différentes contributions au flux des étoiles appartenant au groupe de MWC 1080. Ceci montre l'importance d'une cartographie dans I'IRL des MCS. L'analyse des cartes de polarisation à plusieurs longueurs d'onde de plusieurs OSJ a permis de développer une méthode de détermination directe des densités dans les disques des étoiles en formation. Nous avons donc démontré les possibilités qu'offrent une telle approche des cartes de polarisation. Nous avons déterminé le profil en densité des disques et avons trouvé qu'une approximation en loi de puissance était plausible pour ces structures. La loi de puissance trouvée. est comparable à ce que considère les modèles de la nébuleuse protosolaire ainsi que les modèles de disque de protoétoiles. De plus. le Chapitre 4 semble montrer que le profil en densité du disque est. en un certain sens universel. II serait indépendent du type d'OSJ étudié (T Tauri, EABEH, FU Ori ...) et indépendent de l'âge de la protoétoile. Évidemment, ces résultats sont une première tentative d'extraction de ce genre d'information des cartes polarirnét riques. Une première piste de recherche découlant de ce travail vient du fait que la détermination de la densité dans le disque permet de lever un des paramètres libres des modèles d'interprétation des DES, permettant ainsi une meilleure contrainte des autres paramètres importants. Nous avons proposé et analysé une méthode de détermination du centroide des cartes de polarisation. Le calcul de la moyenne est sensible aux particularités des cartes de polarisation et aux particularités de la distribution en deux dimen- sions des points d'intersection. Nous avons montré que ces distri butions sont très semblables dans leurs cuactéristiques générales à celle des points d7intersection des perpendiculaires à une ellipse. L'analyse détaillée des distri butions devrait produire une estimation plus solide de la position du centroide et pourrait contenir de l'information sur la morphologie de la source d'émission. Les sources majeures d'erreur viennent de la résolution des cartes ainsi que de la technique obser~tionnelleutilisée. En effet, l'observation simultanée des deux polarisations d'un faisceau, les rayons extraordinaire et ordinaire, permettra d'annuler toute variation de transmission atmosphérique. ce que ne permet pas la méthode utilisée actuellement et décrite au Chapitre 2. Des inst niments permettant de telles obsenations sont actuellement en développement, entre autres à l'univer- sité de Montréal. La résolution des cartes est également cruciale dans la mesure où elle décide de la précision avec laquelle nous pouvons obtenir la dimension des disques. Elle affecte également la dimension du disque dans la mesure une mauvaise résolution va étaler le profil parallèle de vecteurs. Whitney et al ( 1997) aborde d'ailleurs ce point. C'est de la détermination précise de la position des points nuls que dépend le succès d'une procédure d'inversion permettant d'obtenir les densités dans le disque. Vne première carte de polarisation obtenue à très haute résolution (0.2") grâce à l'optique adaptative est décrite par Close et al (1997). Elle montre la faisabilité d'une telle entreprise. Sotre dernier objectif concerne la comparaison des résultats obtenus par l'observation millimétrique et subrnillimétrique et ceux obtenus par via la polari- sation. Malgré le peu de données et le peu de recoupement entre les objets observés par les deux techniques. il reste que certaines conclusions importantes sont tirées de cet exercice. La plus importante concerne l'opacité du MCS à la radiation millimétrique. Les densités de colonnes obtenues a un rayon de I .4U de HL Tau montre qu'à cet distance et aux opacités généralement admises pour les grains de poussière (0.1 cm2g-' à 950 Pm, Hildebrand 1983) le chemin optique est de loin supérieur à 1. Les estimations de masse obtenues à partir de flux millimétrique devrait donc être plutôt considérés comme des limites inférieures. Les émissivités des grains également tirées de la photométrie à ces longueurs d'ondes devraient également être revues à la hausse dans la mesure où les lois de puissance des DES sont un mélange d'une loi de Planck à la limite de Rayleigh-Jeans, a u2, et de l'émission du milieu optiquement mince. a u2+?.ou d est l'indice d'émissivité des grains. L'effet de déplacement des points nuls avec la longueur d'onde est maintenant bien établi mais certains résultats proposés. tels celui concernant l''universalité- du profil de densité des disques nécessitent une base statistique plus solide. Il est donc important d'identifier un échantillon complet d'OSJ qui permettrait détablir plus solidement de telles conclusions. L'échantillon devrait également inclure des objets présentant des angles d'inclinaison différents. Ceci devrait permettre d'étudier la structure verticale du disque et de voir si. tel qu'on devrait s'y attendre, les disques polarimétriques changent de dimension avec l'inclinaison. Ceci évidemment nécessite un échantillon statistique suffisamment significatif. L'émergence des techniques d'optique adaptative dans le visible et le PIR ainsi que l'interférométrie dans le submillimétrique va, permettre d'atteindre des résolutions jusque là inégalées dans ces régions du spectre. Ceci va permettre une comparaison plus intéressante des observations aux différentes longueurs d'ondes et une meilleure intégration des données obtenues par ces différentes techniques. Il deviendrait alors possible de mieux connaître la nature des graius de poussière du MCS. un élément jusque là très mal connu de ces environnements. L'exploitation de l'information contenue dans les cartes de polarisation en est encore à ces débuts. Les conclusions observationnelles obtenues dans ce travail doivent encore être reproduites au niveau théorique. Ceci va évidemment amener une meilleure connaissance des géométries du MCS. De plus, il reste encore beaucoup d'informations à tirer des cartes observationnelles. En effet, I'ellipticité du profil de vecteurs de polarisation pourrait contenir de l'information sur la dimension et la morphologie de la source d'illumination de la nébuleuse. La mesure de cette donnée et sa comparaison avec celle à obtenir des modèles pourraient donc porter encore plus d'iaformation sur la géométrie et les propriétés optiques du MCS proche de l'étoile. Une telle information pourrait être difficile à obtenir par d'autres moyens dans la mesure où même les résolutions rendues possibles par l'optique adaptative sont incapables de résoudre le disque d'accrétion à des distances de plus de 1 kpc. Close. L.M.. Roddier, F.. Hora. J.L., Graves. J.E., Northcott. M.. Roddier. C.. Hoffmann. W.F.. Dayal, A.. Fazio. G.G. & Deutsch. L.K. 1997. ApJ. sous presse Hildebrand. R.H. 1983. QJRAS. 34. 267 Whitney. B.A.. Kenyon. S.J. 6; Gornez. M. 1997. .4pJ. sous presse REMERCIEMENTS Un jour d'Août 1990, je suis entré dans un bureau au quatrième étage de l'aile D du pavillon principal. Ignorant tout de l'astronomie - tel que je l'ai découvert plus tard en suivant le cours d'Astronomie Galactique -. Pierre Bastien a bien voulu m'accepter comme étudiant dans son groupe de formation stellaire. je le fut pendat près de sept ans. Je lui dois d'avoir réalisé le rêve de devenir astronome. la chance d'avoir visiter - et quatre fois s'il vous plait - le site astronomique le plus important de la planète. le sommet du Mauna Kea à Hawaï. et le bonheur de travaiiler au CFH qui. pour I'adolescent que j'étais au début des années 80 symbolisait l'astronomie dans toute sa splendeur. Ln jour de Mars 1992. déclarant devant la famille réunie que je ne m'engagerais dans un doctorat que si j'obtenais une bourse. mon père me foudroya du regard. La bourse. je l'ai obtenue. mais je ne serais pas là sans le support moral et financier de mes parents. Après tout. les meubles de la chambre de Nour. les billets d'avion et notre séjour au Liban. toutes mes réalisations. de mes premiers gaga à ma thèse de doctorat, sont le fruit de leur éducation. Mes parents. c'est aussi mes beaux-parents. mon beau-pére qui a accepté de passer trois mois aux frites et aux oeufs rien que pour nous donner la chance d'avoir parmi nous au Canada. mon adorable Tante Adèle sans qui nous serions, Roula et moi, encore en train de se demander comment donner le bain à un bébé de trois jours! Un jour de Mai 1992, une femme bien assise dans son siège au bureau de Mkallès au Liban, accepta de me prendre comme époux. Ensemble, nous xxiv construisons la plus belle des réalisations. une famille dont le premier élément est ?Jour. notre petite étoile, et qui le restera même si la communauté scientifique n'adopte pas Noor comme nom de cette étoile jeune que Papa a étudiée. A toutes les deux. je dois mon bonheur. un soutien et une présence de tous les jours et une raison de perséverer même aux jours où ma thèse, mon travail et tout mon être glissaient vers l'abime. Un soir d'Octobre 1989. je suis arrivé au Canada. Je m'installais finalement au Québec. à Montréal plus précisement, le 31 décembre 1989. Et l'aventure de I'immigré que j'étais commença. Je létais, oui. car je ne le suis plus. grâce à l'accueil extraordinaire que j'ai reçu au sein du groupe d'Astronomie de l'université de Montréal et du Centre du Mont Mégantic. Je remercie tous ceux avec qui j'ai passé mes journées. étudiants et professeurs qui m'ont accepté parmi eux comme l'un des leurs. Une réalité qu'exprime au mieux l'annonce qu'a faite Nicole St- Louis d'un séminaire que présentait "notre Roger national"! Ca fait chaud au coeur! Le Québec. Montréal. le Groupe au sein duquel j'ai vécu. les Rencontres du Centre. les Bistros Brahé. les .Ateliers Étudiants. mon fameux billet d'infraction de 3 15 $. mon acquittement du dit billet et tous ceux que j'ai plus ou moins connus resteront gravés à jamais dans mon coeur. Un grand merci à Luc Turbide qui a bien voulu que je branche mon ordinateur au réseau quaad je n'avais plus d'appartement. qui a mis au point le fichier b'QX udm-these. sty pour nous faciliter la vie et la sienne aussi. Merci pour son incroyable talent d'administrateur de systême et parce qu'il a demandé qu'on le fasse dans ce petit manuel qu'il a écrit au sujet du fichier cité plus haut! Je m'incline devant le talent de Linus Torvalds qui, étudiant de premier cycle à Helsinki. a produit le systême d'opération Linux qui m'a permis de compléter ma thèse dans le comfort douillet de mon domicile sur un Pentiurn 166 plus rapide que les stations de travail SUW du groupe! Khaled. Hanane, Alain. Nadine, Roland, Dany, Nazha, Sami, Sleiman, Elie, Sabil. Georges. Michelle. Jean-Marc, Stefanie, AdeI, Samer, Aimée et tant d'autres sont de ces prénoms qui ont laissé leur trace sur ce travail. Ils ont contribué à la belle expérience que j'ai vécue durant ces années. Je les remercie et espère pouvoir être à mon tour. à leur côté, ce qu'ils ont été pour moi. Ainsi prend fin un travail qui a mis plusieurs années à mûrir. puisse Dieu me donner la chance de continuer à faire de la recherche le restant de mes jours (On oublie toujours de le remercier Lui!). Disons qu'au nombre de postes disponibles. une intervention divine est inévitable pour en décrocher un! Dussé-je pouvoir un jour réalisé le rêve de redonner à mon pays une partie de ce qu'il m'a offert. malgré 15 ans de guerre! L'astronomie n'était-elle pas également cet appel de l'infini. ce goût de l'élan vers les hauteurs. que nous portons en nous! IMAGE EVALUATION TEST TARGET (QA-3) 11111Ei 11111~ Y II, 1.6