phys.astr.co
UF5C, Brésil Roberto Cld Fernandes 
  
Aos Meus 
  
Roberto Cid Fernandes 
  
Grazyna Stasinska 
  
Aos Meus Colaboradores 
  
William Schoenell 
  
Jean Michel Gomes 
  
Abüio Mateus 
  
Laerte Sodré 
  
Marielli Schlickmann 
  
Juan Pablo 
  
Torres Papaqui 
  
Luis Vega 
  
Caroline Deggerone 
  
Christophe Morisset 
  
Yuri Izotov 
  
Ana Gomez 
  
Danielle Michoud 
  
Myriam Rodrigues 
  
Stéphane Tho Mas 
  
Virginie Hababou 
  
David Lépine 
  
Jean-Yves Giot 
  
R Cid Femandes 
  
N Vale Asari 
  
J R Torres-Papaqui1 
  
W Schoenell1 
  
R Cid Femandes 
  
  
  
  
  
  
  
  
Uncovering the Chemical enrichment and mass-assembly historiés of star-forming galaxies

Keywords: galaxies: évolution -galaxies: statistics -galaxies: stellar content galaxies: active galaxies: évolution -galaxies: statistics -galaxies: stellar content, 1 galaxies: évolution -galaxies: statistics -galaxies: stellar content, 1 stars: AGB and post-AGB -galaxies: active -galaxies: general, 1

We analyse galaxies from the Sloan Digital Sky Survey witli our spectral synthesis code starlight. STARLIGHT finds the combination of simple stellar po pulations of different âges and metallicities that best. fits a galaxy spectmm. Tliis allows us to dérivé several galaxy properties, such as mass in stars, star formation history and Chemical évolution. From the residual spectmm (obtained by subtracting the modelled from the observed spectrum), we measure émission lines, from which we obtain properties of the gas inside galaxies.

We study the évolution of star-forming galaxies. We find that the more mas sive galaxies hâve formed stars faster and enriched their gas more quickly. The évolu tion of stellar metallicity is directly studied. We calibrate the current star-formation rate derived from the synthesis with the one from Hcr luminosity.

We dérivé the stellar mass-stellar metallicity relation (MrZ*) for different redshifts. Tins was the first time the évolution of M+-Z+ relation was derived for the same set of galaxies. We find that the metallicity évolution is compatible with that expected from a simple closed box Chemical évolution moclel.

For galaxies classifiée! as LINERs, we show that their observed Hcr luminosity is compatible with the number of ionizing photons emitted by old stellar populations.

Tliis resuit implies a deep review of the rate of nuclear activity in galaxies in the local Universe.

• À CAPES, ao CNPq, e ao projeto internacional CAPES-COFECUB pelo apoio financeiro;

• Ao Nelson Canzian, pelos bate-papos sobre podcast.s de ciência;

• Ao Marco e à Aléa, que compartilharam comigo as aventuras dos nossos doutorados paralelos em Astrofîsica;

• Aos artistas que fazem müsica, cinéma e literatura;

• Ao Peter Wilson, que desenvolveu a charmosa classe memoir para DTgjK;

• A minlia mâe Clarissa, ao meu pai Nelson (in memoriam), aos meus avos Rossini e Enoi, e ao meu irmâo favorito, Marcelo;

• A todos da minlia famflia e da famflia do Fabricio;

• Ao Fabricio, meu companheiro de jornada.

Olhando essa longa lista de pessoas extraordinârias que tornaram esses anos alegres, possfveis e menos traumâticos, estou convencida de que tenho mais sorte do que jufzo por tê-las na minlia vida.

Résumé

Nous analysons les galaxies du Sloan Digital Sky Survey avec notre code de synthèse spectrale starlight. starlight trouve la combinaison des populations stellaires simples de différents âges et métallicités qui reproduit au mieux le spectre d'une galaxie. Il permet de dériver divers paramètres associés à une galaxie, comme la masse totale des étoiles, l'iiistoire de la formation stellaire et l'évolution chimique.

A partir du spectre résiduel (en soustrayant le spectre observé du spectre modélisé), on mesure les raies d'émission, qui permet d'obtenir les propriétés du gaz dans les galaxies.

Nous étudions l'évolution de galaxies à formation stellaire. Nous constatons que les galaxies les plus massives ont formé leur étoiles et leur métaux plus rapidement.

L'évolution de la métallicité des étoiles est étudié directement. Nous calibrons aussi le taux de formation stellaire récente obtenu par notre synthèse à l'aide de celui obtenu à partir de la luminosité de Ha.

Nous dérivons la relation masse stellaire-métallicité stellaire (M*-Z*) à des redshifts différents. C'était la première fois que l'évolution de la relation M*-Z* a été calculée pour les mêmes galaxies. Nous avons constaté que la métallicité stellaire observée a une évolution compatible avec un modèle simple d'évolution chimique en boîte fermée.

Pour des galaxies classées comme LINERs, nous trouvons que la luminosité observée en Ha est compatible avec le nombre de photons ionisants émis par les po pulations stellaires vieillissantes dans ces galaxies. Ce résultat implique une profonde révision du taux d'activité nucléaire dans les galaxies de l'Uiiivers local.

Mots clés: galaxies: évolution -galaxies: statistique -galaxies: contenu stellaire -galaxies: actives.

Résumé (version longue)

Il n'y a pas si longtemps, l'adjectif astronomique se rapportait uniquement à la quantité d'étoiles et de galaxies inaccessibles à l'observation astronomique. Cela a changé dans la dernière décennie, et les astronomes d'aujourd'hui sont fiers de la quantité astronomique d'observations disponibles.

Cette révolution s'est produite grâce aux mega-relevés qui ont balayé de grandes parties du ciel pendant des mois et des mois. L'un d'eux est le Sloan Digital Sky Survey (sdss), dont le plus d'un million de spectres de galaxies, quasars et étoiles ont fait basculer l'astronomie dans une ère nouvelle.

Notre programme de synthèse spectrale semi-empirique de populations stel laires starlight est un des fruits de cette nouvelle ère de l'astronomie, starlight trouve la combinaison des populations stellaires simples (SSPs) de différents âges et métallicités qui reproduit au mieux le spectre d'une galaxie. Il permet aussi de dériver divers paramètres associés à une galaxie, comme la masse totale des étoiles, l'histoire Dans cette thèse, je présente quelques études de galaxies du SDSS qui élargissent le panorama de l'évolution des galaxies dans l'Univers local. Le Chapitre 1 est une courte introduction à l'évolution des galaxies et présente le contexte scientifique de cette thèse. Le Chapitre 2 présente les données du SDSS que nous utilisons, notre algorithme de synthèse de populations stellaires STARLIGHT, ainsi que notre code de détection de raies d'émission.

Dans le Chapitre 3 nous obtenons l'histoire de formation stellaire et l'évolution chimique pour un échantillon de plus de 80000 galaxies à formation stellaire (SF) dans le SDSS. Nous étudions ces galaxies en fonction de leur position sur le diagramme [N il]/Ha vs. [O iii]/H/3 en les groupant en bins de métallicité nébulaire (Zneb). Nous les étudions également dans des bins de masse stellaire. Nous constatons que les galaxies les plus riches en métaux (et les plus massives) ont formé leur étoiles et leur métaux plus rapidement. Les galaxies de plus basse métallicité (et moins massives)

sont plus lentes dans la formation des étoiles et la production de métaux. C'était la première fois que l'évolution de la métallicité des étoiles a été étudié directement.

Le Chapitre 4 explore plus en détail les propriétés des galaxies SF du SDSS.

Nous trouvons une corrélation entre Zneb et divers paramètres physiques et observati onnels, comme la niasse stellaire, l'âge et métallicité moyenne stellaire, la luminosité et la quantité de poussière. Nous calibrons le taux de formation stellaire récente ob tenu par notre synthèse à Laide de celui obtenu à partir de la luminosité de Ha. Cette calibration peut être utilisée pour étudier les galaxies hôtes de AGNs, où la lumino sité de Ha ne peut pas être utilisée comme indicateur de formation stellaire. Nous approfondissons également l'étude des histoires de formation d'étoiles de ces galaxies, des distributions de ces histoires et des effets de sélection, en regardant en détail la distribution des histoires de formation stellaire à l'intérieur d'un bin en métallicité, et étudions les causes de biais possibles.

Le Chapitre 5 traite de l'évolution de la relation masse stellaire-métallicité stellaire (M*-ZJ au cours du temps. A partir de l'évolution de la masse stellaire et de la métallicité obtenue par starligi-it, nous dérivons la relation M+-Z+ à des redshifts Le Chapitre 7 résume les principaux résultats de cette thèse et propose des di rections futures. Nous avons présenté de nouvelles méthodes pour étudier l'évolution des galaxies en utilisant les propriétés obtenues par la synthèse spectrale de popu lations stellaires. Nos méthodes ouvrent la porte, entre autres, à des études plus approfondies sur les relations entre les AGNs et la formation stellaire.

Our semi-empirical stellar population spectral synthesis code starlight is in line with tliis new era in astronomy. starlight finds the combination of simple stellar populations (SSPs) of different âges and metallicities that best fits a galaxy spectrum. Tliis allows us to dérivé several galaxy properties, such as the mass in stars, the star formation history and the Chemical évolution. Ail tliis is done in a highly automated fashion, suitable to analyse linge samples like tliose coming from the SDSS.

Furthermore, from the pure nebular residual spectrum (obtained by subtracting the modelled from the observed spectrum), we measure émission lines. We tlien obtain several gas properties, such as Chemical abundances and nebular extinction.

Emission lines also provide important information ou the ionization sources in a galaxy through line-ratio diagnostic diagrams. For instance, galaxies form two "seagull wings'1 on the [N Ii]/Ha vs. [O m]/H/5 diagram: Galaxies whose gas is ionized by a barder source, such as active galactic nuclei (AGNs), live on the right wing; those whose gas is ionized by young stars are found on the left, wing.

Tliis thesis encompasses some studies on SDSS galaxies wliicli put constrains on the évolution of galaxies in the local Universe. Chapter 1 is a short introduction to the évolution of galaxies and the scientific context of this thesis. Chapter 2 reviews the SDSS data, our stellar population synthesis algorithm STARLIGHT, and our emissionline détection code.

In Chapter 3 we obtain the star formation history and Chemical évolution for a sample of over 80000 star-forming (SF) SDSS galaxies. We study these galaxies according to their position on the [N Ii]/Ha vs. [O Iii]/H/? diagram by grouping thern into nebular metallicity (Zneb) bins. We also study them in stellar mass bins. We find that the more metal-rich (and more massive) galaxies hâve formed stars faster and enriched their gas more quickly. On the other hand, the less metal-rich (and less massive) ones are slower in forming stars and in polluting their gas with metals. This was the first, time the évolution of stellar metallicity was directly studied.

Chapter 4 explores the properties of SF galaxies. We show how Zneb correlates with some physical and observational properties, such as stellar mass, mean stellar âge, mass-weighted stellar metallicity, brightness and dust content. We calibrate the currcnt star-formation rate dérivée! from the synthesis with the one from Ha luminosity. This calibration can be applied in future studies of AGN hosts, for which

Ha cannot be used as a star formation indicator. We also study more deeply their star formation historiés, distributions of historiés and possible bias and sélection effects.

Chapter 5 deals with the évolution of the stellar mass-stellar metallicity re lation (M*-Z*) with look-back time. From the stellar mass and metallicity historiés of SDSS galaxies recovered with STARLIGHT, we dérivé M*-Z* for different redshifts.

This was the first time the évolution of M*-Z* relation was derived for the same set of galaxies. We find that M+-Z* steepens and spans to lower mass and metallicities with look-back time. We find that the metallicity évolution is compatible with that expected from a simple closed box model. This means that, as a first approximation, the relation for galaxies with stellar mass greater than ~1010 M0 may be explained by variations in their star formation historiés, and not by inflows or outflows of gas.

Chapter 6 focuses on LINERs, galaxies classified as AGN hosts, but which could in fact be ionized by hot old stars. For a fraction of the galaxies classified as LINERs in the SDSS, we show that their observed Ha luminosity is compatible with the number of ionizing photons emitted by SSPs older than 107 5 yr in those galaxies.

In addition, photoionization models using the ionization field obtained for old stellar populations almost completely cover diagnostic diagrams, such as the [N Ii]/Ha vs.

[O mJ/H/3, which means that emission-line ratios are also compatible with ionization by old stellar populations. The "seagull" shape of the [N n]/Ha vs. [O iii]/H/3 diagram is due to sélection effects. Tins resuit implies a deep review of the rate of nuclear activity in galaxies in the local Universe.

Chapter 7 summarises our main results and proposes some future studies. Tliis thesis présents new methods to study the évolution of galaxies by using properties recovered from stellar population spectral synthesis. These methods open the patli, among otliers, to more detailed studies about the connection between AGN and star formation.

Evoluçâo qufmica e historia de formaçâo estelar no Universo local Resumo Analisamos as galâxias do Sloan Digital Sky Survey corn o nosso côdigo de sfntese espectral starlight. O starlight acha a combinaçâo de populaçôes estelares simples de diferentes idades e metalicidades que melhor modela o espectro de uma galâxia. Ele permite derivar vârios parâmetros associados a uma galâxia, como a massa em estrelas, a historia de formaçâo estelar e a evoluçâo qufmica. A partir do espectro residual puramente nebular (subtraindo do espectro observado o modelado), medimos as linlias de emissâo, das quais derivamos propriedades do gâs dentro das galâxias.

Estudamos a evoluçâo das galâxias com formaçâo estelar. Constatamos que as galâxias de maior massa formaram suas estrelas e seus metais mais rapidamente.

A evoluçâo da metalicidade das estrelas é estudada diretamente. Calibramos também a taxa de formaçâo estelar atual medida pela sfntese com a medida pela luminosidade de Ha.

Derivamos a relaçâo massa estelar metalicidade estelar (MrZ*) em diferentes redshifts. Esta é a primeira vez que a relaçâo M*-Z* é calculada para o mesmo conjunto de galâxias. Observamos que a metalicidade estelar observada tem uma evoluçâo compatfvel com uni modelo simples de evoluçâo qufmica de caixa fecliada.

Para as galâxias classificadas como LINERs, encontramos que a luminosidade observada em Ha é compatfvel com o numéro de fôtons ionizantes emitidos pelas populaçôes estelares velhas dessas galâxias. Este resultado implica em uma profunda revisâo da taxa de atividade nuclear nas galâxias do Universo local.

Palavras-chave: galâxias: evoluçâo -galâxias: estatfstica -galâxias: conteüdo estelar -galâxias: ativas.

Resumo (versâo longa)

Até pouco tempo o ternio quantidade astronômica restringia-se a descrever a enormidade de estrelas e galâxias no céu inacessfveis à observaçâo astronômica. O cenârio mudou lia ültima década, e os astrônomos hoje gabam-se da quantidade as tronômica de observaçôes disponfveis. Essa revoluçâo aconteceu graças aos grandes levantamentos de dados, que varreram pedaços do céu durante meses e meses. Uni desses megaprojetos foi o Sloan Digital Sky Survey (sdss), cujos mais de uni milhâo de espectros de galâxias, quasares e estrelas empurraram os métodos de estudos astronômicos para uma nova era.

O nosso programa de sfntese espectral semi-empfrica de populaçôes estela res STARLIGHT é fruto dessa nova era da astronomia. O STARLIGHT aclia a com binaçâo de populaçôes estelares simples (SSPs) de diferentes idades e metalicidades que melhor modela o espectro de uma galâxia. A partir desse modelo, derivamos vârios paramétras associados a uma galâxia, como a massa em estrelas, a historia de formaçâo estelar e a evoluçâo qunnica. Tudo isso é feito de forma automâtica, apropriada para analisar amostras énormes como a do sdss.

Além disso, a partir do espectro residual puramente nebular (subtraindo do espectro observado o modelado coin o STARLIGHT), medimos *is linhas de emissâo. Com isso, obtemos parâmetros associados ao gâs de mua galâxia, tais como abundâncias qufmicas e avermelliamento. As linhas de emissâo também fornecem informaçôes importantes sobre as fontes de ionizaçâo em uma galâxia a partir de diagramas de diagnôstico de razôes de linhas. No diagrama [N Ii]/Ha vs. [O iii]/H/?, por exemplo, as galâxias se distribuem em duas "asas" : aquelas cujo gâs é ionizado por uma fonte mais dura, como nücleos ativos (AGNs), residem na asa direita, e aquelas c.ujo gâs é ionizado por estrelas jovens, na asa esquerda.

Nesta tese, apresento alguns estudos sobre galâxias do SDSS que ampliam o panorama sobre evoluçâo de galâxias no Universo local. O Capitulo 1 é uma pequena introduçâo sobre evoluçâo de galâxias e ao contexto cientffico dessa tese. O Capitulo 2 révisa os dados do SDSS que utilizamos nesses estudos e os nossos algoritmos de sfntese de populaçôes estelares STARLIGHT e de medida de linhas de emissâo.

No Capitulo 3 obtemos a liistôria de formaçâo estelar e a evoluçâo qunnica para uma amostra de mais de 80000 galâxias com formaçâo estelar (SF) do SDSS.

Dividimos as galâxias em faixas de acordo corn a posiçâo no diagrama [N n]/Ha vs.

[O m]/H/3, que é um indicador a metalicidade nebular (Zneb), e também em faixas de massa estelar. Verificamos que as galâxias de maior metalicidade nebular (e de rnaior massa estelar) formaram suas estrelas mais râpido e pohuram seu gâs com metais mais rapidamente. As galâxias de menor metalicidade (e de menor massa) sâo mais lentas na formaçâo de estrelas e 11a, produçâo de metais. Esta foi a primeira vez em (pie a evoluçâo da metalicidade estelar foi estuda diretam,ente.

O Capitulo 4 traz uma visâo mais detalhada das galâxias SF do SDSS. Veri ficamos a correlaçâo entre Zneb e vârios parâmetros ffsicos e observacionais, como a massa estelar, a idade e metalicidade estelar média, o brilho da galâxia e a quant idade de poeira. Calibramos também a taxa de formaçâo estelar atual medida pela sfntese com a medida pela luminosidade de Ha. Essa calibraçâo pode ser usada em estudos futuros sobre galâxias hospedeiras de AGNs, nas quais Ha nâo pode ser usada como indicador de formaçâo estelar. Também aprofundamos o estudo das historias de formaçâo estelar dessas galâxias, olhando com mais detallies a distribuiçâo de historias para galâxias dentre de um grupo de Zneb, e estudamos possfveis causas de vieses.

O Capitulo 5 trata da evoluçâo da relaçâo entre massa estelar e metalicidade estelar (M*-Z*) ao longo do tempo. A partir das historias de massa e metalicidade estelar das galâxias do SDSS, derivamos a relaçâo M* Z* para diferentes redshifts. Esta foi a primeira vez em que a evoluçâo da relaçâo M*-Z* foi derivada para o m,esmo conjunto de galâxias. Verificamos que M* Z* fica mais inclinada e cobre valores menores de massa e metalicidade à medida em que voltamos 110 tempo. Testamos um modelo simples de caixa fechada e verificamos que a metalicidade estelar tem uma evoluçâo consistente com esse modelo. Isso significa que a relaçâo M*-Z* para galâxias de massa maior do que ~1010 M© pode ser explicada em primeira aproximaçâo somente pela diferença em historias de formaçâo estelar, sein que seja necessârio levar Hâ menos de uni século, ninguém sabia se existiam outras galâxias além da Via Lâctea. Considerando o curtissimo tempo que tiveram para montar o quebra-cabeça, os astrônomos hoje jâ tem um desenlio coerente, ainda (pie incompleto, dos bilhôes de anos de evoluçâo das galâxias no Universo.

O Universo surgiu hâ cerca de 14 bilhôes de anos com o Big Bang. Uma consequência desse infcio intenso é que o espaço continua se expandindo até hoje.

Por isso, as galâxias estâo se afastando umas das outras, salvo quando estâo nniito prôximas e a força gravitacional vence a expansâo do espaço.

Quando o Universo era jovem e os primeiros âtomos de hidrogênio e liélio se formaram, logo surgiram as primeiras estrelas. Mais adiante, formaram-se as primeiras galâxias. As maiores galâxias de hoje tiveram uma juventude conturbada: formaram militas estrelas nniito rapidamente. Essas estrelas produziram elementos mais pesados que H e He (referidos em Astronomia como metais, em contraste com a definiçâo usual de metais em Fisica ou Quhnica) e logo morreram. A perda de matéria e a explosâo em supernova dessas estrelas poluiu rapidamente o gâs nessas galâxias.

Consequentemente, as geraçôes seguintes de estrelas, que se formaram a partir gâs reciclado, nasceram ricas em metais. As galâxias menores tiveram um passado mais tranquilo, formando estrelas e produzindo metais bem mais lentamente. Naquela época, existiam ainda militas galâxias que hospedavam quasares, buracos negros que engolem matéria a taxas extravagantes.

Hoje, no Universo coin 14 bilhôes de anos, o cenario se inverteu. As galâxias de maior massa, mais vermelhas, mais bojudas, e cujo gâs é mais enriquecido em metais formam menos estrelas por unidade de massa. As galâxias de menor massa, mais azuis, de meio interestelar mais pobre, e com discos maiores sâo as que têm maior atividade de formaçâo estelar. Essa diferença na histôria de formaçâo de estrelas *No original: "Galaxies, like éléphants, hâve long memories". A historia de como a astronomia chegou a esse cenârio é um pouco tortuosa. Alguns apresentavam evidências de que essas nebulosas estavam dentro da Via Lâctea;

outras evidências favoreciam a interpretaçâo de que elas eram "universos-ilhas", ou galâxias como a nossa. Trimble (1995) faz uma otima recapitulaçâo desse debate.

A questâo so foi resolvida quando Hubble (1925) Mas a astronomia possui mais de uma maneira de estudar o passado. Ao mesmo tempo, esforços também foram dirigidos para estudar galaxias em funçâo das estrelas que as constituem. Antes mesmo de baver bons modelos para a evoluçâo estelar, Whipple (1935) calculou uni espectro sintético de uma galâxia somando estrelas de diferentes tipos espectrais. Baade (1944) propos que galaxias deveriam ser constituidas por populaçôes estelares parecidas com as de aglomerados.

Como a evoluçâo de estrelas ainda era incerta, esses estudos acabaram nâo rendendo muitos frutos. A ideia, no entanto, era muito boa: estudar as galaxias a partir das suas estrelas. As estrelas maiores de uma geraçâo morrem cedo, enquanto as estrelas menores sobrevivem para contar como era o ambiente na galâxia na época no seu nascimento. Portanto, uma populaçâo estelar velha jâ viu todas as suas estrelas azuis de alta massa morrerem, e, portanto, tem uni espectro muito diferente de uma populaçâo jovem. Dessa forma, no espectro integrado das galaxias esta impresso o registro fôssil das idades e metalicidades das populaçôes estelares que ali nasceram.

Todas as informaçôes estavam la, mas demoraram uni pouco para serein organizadas.

1.1.4

Os primeiros modelos modernos de evoluçâo de galaxias

Tanto o estudo da evoluçâo de galaxias pela mâquina do tempo (observando galaxias em estâgios de evoluçôes diferentes em diferentes redshifts) quanto pelos registros fosseis (decompondo o espectro de galaxias em populaçôes estelares de diferentes ida des) nâo iriam muito longe se nâo houvesse modelos teôricos coerentes para a evoluçâo de galaxias. O trabalho séminal Tinsley (1968) foi uni dos primeiros a sistematizar a construçâo de modelos evolutivos de galaxias, obtendo paramétras sintéticos observâveis como cores, conteudo do gâs e abundância qunnica de galaxias. Tinsley assumiu que apenas a distribuiçâo de idades das estrelas variava, fixando a idade es- A partir de entâo, a evoluçâo de galâxias começou a ser mais seriamente estudada. Vârios estudos previam ou encontravam correlaçôes entre parâmetros ob servâveis e ffsicos. Lequeux et al. (1979), por exemplo, observaram a relaçâo entre a massa em estrelas e a metalicidade do meio interestelar em galâxias, como previsto teoricamente por outros autores (ver referências em Lequeux et al. 1979).

A forma da relaçâo luminosidade ou massa-metalicidade (M~Z) ainda é objeto de controvérsia até hoje. Se a simples correlaçâo pode ser prevista pela evoluçâo quûnica considerando a galâxia uma caixa fechada, a forma exata da relaçâo M-Z pode ser influenciada pela entrada ou a safda de gâs ou por outros fenômenos ffsicos.

Outros fenômenos foram descobcrtos para élucidai-a evoluçâo de galâxias, como o cenârio de downsizing (por exemplo Broadhurst, Ellis, & Shanks 1988;Cowie et al. 1996;Heavens et al. 2004;Cid Fernandes et al. 2005 

Evoluçâo de galaxias

das estrelas no bojo de uma galaxia e a massa do buraco negro central (Ferrarese & Merritt 2000;Gebhardt et al. 2000).

Avançamos déniais na historia e nas perguntas atnais, no entanto. Voltemos um pouco para quando os modelos evolutivos estavam nascendo. O STECMAP/STECKMAP5 trata a recuperaçâo das SSPs a partir do es pectro integrado de uma galâxia como um problema de inversâo de matrizes. Para resolver as complicaçôes de degenerescências (como entre a idade e metalicidade) e as limitaçôes impostas pelo razâo sinal/rufdo do espectro, é usada uma técnica de regularizaçâo que consiste em suavizar a resoluçâo temporal da historia de formaçâo estelar recuperada. Esse processo é feito automaticameute pelo côdigo com base no objeto 4Est,a frase jâ foi atribufda a vârios astrônomos, como Andrea Dupree [START_REF] Dupree | The Future of Cool-St.ar Astrophysics: 12th Cambridge Workshop on Cool Stars, Stellar Systems, and the Sun[END_REF]Dolla 2006), Gary Ferland (Dultzin-Hacyan, Marziani, & Sulentic 2000) e Joe Miller [START_REF] Black | Higli Résolution Infrared Spectroscopy in Astronorny[END_REF], mas em gérai é simplesmente tornada como "sabedoria popular" [START_REF] Brickhouse | Astronomical Data Analysis Software and Systems VIII[END_REF][START_REF] Pallavicini | [END_REF][START_REF] Schlegel | The restless universe: understanding X-ray astronomy in the âge of Chandra and Newton[END_REF]).

5http: / / astro.u-strasbg.fr/Obs/G ALAXIES / stecmap_eng.html Figura de Maiolino et al. (2008).

automaticamente as linhas de emissâo (Stasinska et al. 2006, S06). Na Figura 1. Por outro lado, estudos a redshifts mais altos também revelam comporta-/ mentos muito intéressantes de galâxias. E necessâria a comparaçâo constante entre os resultados obtidos com a "mâquina do tempo" e coin os "registros fôsseis'7 para montar cada vez com mais precisâo o filme da evoluçâo das galâxias. Por exemplo, forain encontradas evidências que a maior parte de formaçào de estrelas no Universo aconteceu no passado (Lilly et al. 1996), com uni pico ao redor do redshift z ~2 (Madau et al. 1998). A evoluçâo da relaçâo massa-metalicidade também é estudada a redshifts intermediârios e altos (por exemplo, Savaglio et al. (2005); Maiolino et al. (2008)). Na Figura 1.4 mostramos os resultados de Maiolino et al. (2008): a relaçâo M-Z desloca-se para baixo e muda de inclinaçâo com o redshift.

Ainda hâ muito explorai-sobre a evoluçâo de galâxias, esta tese tenta avançar algumas dessa novas fronteiras. 

Dados

Nesta tese, utilizamos os espectros das galâxias do DR5 do SDSS, excluindo as observaçôes duplicadas (~1.6% dos objetos). A nossa amostra-mâe contém 573141 objetos.

Corrigimos os espectros por extinçâo galâctica coin os rriapas de Schlegel, Finkbeiner, & Davis (1998) e lei de extinçâo de Cardelli. Clayton, & Mathis (1989) usando Ry = 3.1. Reamostramos os dados para obter AA = 1 Â ira faixa espectral de 3400 a 8900 Â.

Quando apropriado, usamos os seguintes paramétras cosmolôgicos: H0 = 70 km s-1 Mpc-1, Qm -6.3 e Da = 0.7. A nossa base contém 150 SSPs de BC03 calculadas para uma funçâo inicial de massa de Chabrier (2003), caminhos evolutivos de "Padova 1994" (Alongi et al. 1993;Bressan et al. 1993;Fagotto et al. 1994a,b;Girardi et al. 1996) Esse indicador de abundância, 110 entanto, nâo estâ livre de complicaçôes.

Sintese espectral de populaçôes estelares

Ele é sensfvel à presença de gâs ionizado difuso e às variaçôes 11a razâo N/O para diferentes cenârios de evoluçâo qufmica (Chiappini, Romano, & Matteucci 2003).

Além disso, ele nâo é sensfvel a variaçôes de O/H para metalicidades baixas (ver figura 3b de Stasinska 2006). De todo modo, como argumentareinos na Seçâo 4.2 a partir de testes com outras calibraçôes de metalicidade, O3N2 é o indice mais conveniente matematicamente para esse trabalho, além de ser razoâvel para o problema astroffsico em quest âo.

Os valores de Zneb para a nossa amostra de galaxias S F vâo de 0.2 a 1.6 Z©.

Para agrupar galaxias do mesmo tipo e intrinsecamente similares, definimos faixas em logZIieb centradas em -0.580 (faixa A), -0.385 (B), -0.255 (C), -0.125 (D), 0.005 (E) e 0.135 (F). Todas as faixas têm 0.13 dex de largura, exceto para de menor metalicidade, que foi definida como duas vezes mais larga para incluir mais objetos.

O painel b da figura 1 de C07 mostra a disposiçâo das faixas 110 BPT.

Apesar de Zneb ser uni parâmetro conveniente, a metalicidade nebular é um produto da evoluçâo de galaxias. Do ponto de vista astroffsico, a massa de uma galâxia mostra-se um fator déterminante para a sua evoluçâo (por exemplo, Tinsley 1968). Repetimos portanto o nosso estudo para faixas de massa centradas em log M*/M© = 7. 5, 8.5, 9.5, 10.5 e 11.5 e com 1 década de largura.

3.3

Derivando historias de formaçao e evoluçâo qufmica 

Resultados

A Figura 3.1a mostra as curvas da funçâo cumulativa de massa para a amostra de galaxias S F dividida em faixas de metalicidade nebular. As curvas de histôrias de formaçâo estelar exibem uma organizaçâo clara. As galaxias que hoje têm maior Zneb formaram a maior parte das suas estrelas hâ bastante tempo. As galaxias de menor Zneb demoraram significativamente mais para formar estrelas.

A Figura 3.1b mostra a evoluçâo da metalicidade estelar média. Novamente vemos a organizaçâo das curvas em faixas de Zneb: as galaxias cujo meio interestelar é mais pobre em metais tiveram uma evoluçâo da metalicidade estelar mais lenta.

Uma maneira de quantificar a diferença nas histôrias de massa é a partir de , definido como quanto tempo atrâs uma galâxia formou 75% da sua massa em estrelas (ou seja, encontrando ??*(£* = r^f) = 0.75). Da mesma forma, definimos r^5 como quanto tempo atrâs uma galâxia atingiu 75% da sua metalicidade atual. Os Vale destacar que estamos comparando informaçôes obtidas independentemente: as histôrias de massa e metais das galaxias forain derivadas da anâlise de populaçôes estelares com STARLIGHT; a metalicidade nebular foi obtida a partir das linhas de emissâo. A compatibilidade entre esses resultados é outro triunfo dos métodos de sintese espectral.

3.4.

Resultados

A figura 3 de C07 mostra a historia da massa e dos metais para as galâxias SF divididas em faixas de massa estelar atual. Novamente encontramos a evidência de downsizing tanto lia formaçâo de estrelas quando no enriquecimento qmmico de galâxias, isto é, galâxias de maior massa formaram estrelas e polufram seu gâs muito mais rapidamente do que galâxias menores e menos metâlicas.

Um dos aspectos mais importantes deste estudo foi que se pôde, pela primeira vez, estudar diretamente a evoluçâo da metalicidade estelar. Trabalhos anteriores a partir de métodos fosseis estudaram somente a evoluçâo da massa em estrelas (por exemplo, Heavens et al. 2004), e nâo da metalicidade. O programa de linlias de emissâo também foi modificado para resolver o problema do resfduo sistemâtico ao redor de H fi (Figura 4.1). A partir da inspeçâo dos nossos espectros modelados, verificamos que a populaçâo estelar em ~200 À ao redor dessa linha é superestimada. Isso implica que a linha de emissâo de H/3 fica dentro de uma absorçâo larga no espectro residual. Acreditamos que isso seja uni efeito de calibraçâo na biblioteca de estrelas STELIB (ver a figura 21 de Martins et al. 2005).

Dessa forma, se supomos que o contmuo ao redor de H[3 é zéro, como faziamos nas versôes anteriores do programa de medida de linhas, estamos, de fato, perdendo uni pedaço da linha e subestimando o seu fluxo. Para esse estudo, o programa foi modifi cado de forma a levar em consideraçâo o contmuo local ao redor de H fi. As medidas do fluxo sâo em média 2% maiores do que as anteriores, mas para objetos com largura équivalente de H/3 < 2 À o diferença chega em média a 7%. A extinçâo nebular também aumenta com Zne^. Isso esta ligado, entre outros 4.4.

Taxas de formaçâo estelar em funçâo do tempo fatores, à maior formaçâo de grâos de poeira quando a metalicidade nebular alimenta.

Também existe uma relaçâo muito boa entre a metalicidade estelar média e Zneb, corno ficou évidente nos estudos da metalicidade estelar ern funçâo do tempo de C07 (Capftulo 3).

Os painéis f, g e h sâo uni resumo da histôria das galâxias corn formaçâo estelar. A relaçâo entre idade média e Zneb é de novo consequência do downsizing, isto é, as galâxias cuja luz é dominada por estrelas mais jovens têm o meio interestelar mais pobre, e as dominadas por estrelas veillas têm o meio interestelar mais rico em metais. Existe uma anti-correlaçâo entre Zneb e a razâo entre taxa de formaçâo estelar recente e a passada ( 6), definida pela Equaçâo 4.8 na Seçâo 4.5. Ou seja, as galâxias de menor metalicidade formam hoje mais estrelas do que no passado. O painel h mostra que a luminosidade em Ha, em muitos estudos usada como indicador da taxa de formaçâo estelar recente, também correlaciona-se com Zneb.

Taxas de formaçâo estelar em funçâo do tempo

Jâ mostramos na Seçâo 3.3 uma maneira de comprimir as informaçôes sobre as SSPs que saern do STARLIGHT para estudar historias de formaçâo estelar. Naquele caso, usamos a histôria cumulativa de massa das estrelas.

Outra forma de comprimir os dados é derivar a taxa de formaçâo estelar em funçâo do tempo. Obviamente, nâo podemos confiar individualmente em cada uma das 25 idades das SSPs da base. Como discutido ua seçâo 4.2 de A07, estima-se que métodos fôsseis têm uma resoluçâo temporal entre 0.5 a 1 dex. Isso significa que podemos distinguir uma populaçâo de 108 de outra de 109 anos, mas nâo uma de 108 de outra de 108 1 auos. Ou seja, nâo podemos confiar em cada uma das 25 idades das populaçôes estelares da nossa base. Para diminuir a resoluçâo dos nossos resultados, aplicamos um filtro gaussiano de 1 dex de largura a meia altura em log A no vetor de fraçâo de massa //(A) e de luz #(A), obtendo as suas versôes suavizadas ^(A) e

XS(U).

A taxa de formaçâo estelar em funçâo de A é obtida entâo a partir dessa versâo suavizada das fraçôes de massa:

SFR(A) - dM«(A) dA AAg(C) = M?loge fics(Q AU U A log U ' (4.2)
na quai A descreve as versôes continuas nos nossos vetores de populaçâo, com A log O = 0.1 dex. O fator loge vem da trausformaçâo de dA em Ad (log A):

lu A = log A loge ni ^\ dt* => d(ln A) = - t* d(logA) log e
Podemos définir também a taxa de formaçâo estelar especffica, isto é, em relaçâo à massa atual em estrelas:

SSFR(A) SSFR(A) Ml loge ncs(U) A A log A ' (4.3)
A figura 8 de A07 mostra a historia de formaçâo estelar em 4 formas para cada faixa A-F em Zneb: xs, SFR, SSFR e rf (este ültimo definido anteriormente pela Equaçâo 3.2). No painel a, vemos que as faixas de maior Zneb têm maior contribuiçâo em luz de populaçôes mais veillas, e as de menor Zneb, de populaçôes mais jovens.

O painel b mostra a organizaçâo das faixas em Zneb em SFR. As galaxias de maior metalicidade têm uma taxa de formaçâo estelar mais plana e constante do que as de menor Zneb. As de menor Zneb possuem uma taxa de formaçâo estelar que cresce do passado até o présente, atingindo atualmente (~1U6 anos atrâs) uma taxa de formaçâo estelar 100 maior do que quando foi formada (~lü10 anos atrâs). Em termos de quantidade absoluta de massa em estrelas, as galaxias de maior ©neb ainda formam mais estrelas hoje do que as de menor Zne^.

O painel c mostra a organizaçâo das galaxias em faixas de Zneb para a SSFR.

As galaxias de menor metalicidade têm atualmente uma SSFR maior, o que significa que elas dobrariam em massa de estrelas mais râpido do que as galaxias de maior metalicidade. O painel d mostra a historia cumulativa de massa, jâ discutida na Seçâo 3.4.

Como argumentamos em C07, Zneb é antes uma consequência da evoluçâo quimica, e nâo uni ingrediente fisico que afeta diretamente a evoluçâo de uma galâxia.

Portanto, dividimos também a nossa amostra pela massa M* e pela densidade superficial de massa E*. Nas figuras 9 e 10 de A07 mostramos as diferentes representaçôes da historia de formaçâo estelar em faixas desses dois paramétras. As tendências encontradas corn Zneb sâo mantidas.

E importante observar que esses resultados estâo baseados na média das curvas de SFHs para cada faixa de metalicidade, massa ou densidade. Na figura 7 de A07 mostramos a distribuiçâo da SSFR, para a faixas A e F de Zne^. As distribuées sâo assimétricas e aproximadamente log-normais para todas as idades. Apesar de as distribuées serein largas e se superporem, elas sâo diferentes entre si e a separaçâo entre as faixas A-F é real. Testes Kolmogorov-Smirnov para qualquer par de faixas de Zneb revelam que as distribuées sâo diferentes com 99% de confiança para qualquer U.

Taxa de formaçâo estelar atual

Para medir a taxa de formaçâo estelar atual, uni dos indicadores mais usados é a luminosidade em Ha (por exemplo, Kennicutt 1983). Este método assume que a SFR é constante durante o tempo de vida de estrelas massivas ionizantes, isto é, cerca de 10 milhôes de anos. A taxa de fotons ionizantes converge entâo para: Fizemos também testes aplicando extinçôes diferentes para estrelas joverrs e vellras dentro de uma galâxia. A figura 4 de A07 nrostra que a extinçâo nebular Ayeb é cerca de duas vezes a extinçâo estelar para a nossa amostra. Como as populaçôes joverrs ionizantes devenr estar proximas do gâs ionizado dentro de uma galâxia, esse resultado implica também que as SSPs jovens devenr sofrer uni maior avermelhamento por poeira do que populaçôes velhas.

Qh = SFR A/# (IMF, Z*), ( 4 
Assim, forçanros os modelos coin o STARLIGHT para aplicarenr unra extinçâo 0.34 + 2.28Ay para as SSPs mais jovens do que 10' anos, onde Ay é a extinçâo recuperada para as SSPs vellras e a relaçâo anterior é o ajuste linear robusto para a relaçâo Ayeb-Ay. O que se espera coin esse teste é que o programa récupéré unra fraçâo maior de SSPs jovens, jâ que estamos forçando que elas sejanr mais obscurecidas por poeira. Os resultados sâo mostrados na Figura 4.2. Novanrente, o mâxinro que ocorre é a migraçâo de unra curva de uma faixa Zneb para a regiâo onde estava a curva consecutiva, mas a organizaçâo das SFHs se nranténr.

Sumario

A partir do registro fôssil das populaçôes estelares enr espectros integrados de galâxias, mostramos que galâxias SF de menor metalicidade (ou massa estelar, ou densidade superficial de massa) sâo as que hoje têm maior formaçâo estelar especffica. Isso é evidência do downsizing das galâxias.

Derivanros tambénr a taxa de formaçâo estelar atual a partir da smtese, calibrada a partir da SFR de Ha. Além de fornecer um teste para os nossos resultados da smtese, esse nrétodo pode ser aplicado para récupérai' a SFR atual enr galâxias nas quais nâo é possfvel estimar a SFR a partir de Ha, como as liospedeiras de AGN, onde parte da luminosidade enr Ha é dévida à ionizaçâo pela fonte de radiaçâo central mesmas curvas para a nossa amostra de 82662 objetos dividida em faixas de M*, que inclui todos os tipos de galaxias. Vernos de novo o efeito de downsizing: as galâxias de maior massa tiveram uma formaçâo de estrelas e uma produçâo de metais muito râpida no passado.

Desde Lequeux et al. (1979), que observou a correlaçâo entre luminosidade e a metalicidade nebular, a relaçâo luminosidade (ou massa) e metalicidade é usada para entender a evoluçâo qunnica de galâxias. Portanto, uma forma didâtiea de olhar a figura 1 de V09 é plotando a relaçâo M+-Z+ para um tempo dado. Mostramos essas curvas na Figura 5.1.

Assumindo que a fusâo de galâxias é insignifiante na faixa de redshifts que estudamos (< 1.9), encontramos que relaçâo M* Z* em redshifts mais altos é mais inclinada e inclui mais objetos de baixa massa e metalicidade do que a relaçâo para pequenos redshifts. Vernos também que as galâxias mais massivas passaram por pouca evoluçâo na massa e na metalicidade nos ültimos 9 bilhôes de anos.

Essa é a primeira vez que se mostra a evoluçâo da relaçâo Af*-Z* para o mesmo conjunto de galâxias. Em gérai os estudos em altos redshifts estudam objetos dife rentes e comparam-nos a objetos a baixo redshift (por exemplo, Savaglio et al. 2005;Maiolino et al. 2008). No entanto, tais estudos estâo sujeitos a vieses, principalmente coin respeito à calibraçâo das metalicidades nebulares (ver Stasinska 2008).

Modelo simples para a evoluçâo qufmica

A causa da relaçâo M-Z é hâ muito tempo objeto de limitas especulaçôes. Garnett Modelo simples para a evoluçâo qufmica invadidas por ventos de gâs pristino. Nesse caso, as galâxias menores seriam mais afetadas e teriam a sua concentraçâo de metais mais dilufda pela entrada de gâs. Koppen, Weidner, & Kroupa (2007) chamam a atençâo para possfveis diferenças na funçào inicial de massa integrada, isto é, galâxias mais massivas formariam aglomerados de estrelas mais massivos, que por sua vez contêm mais estrelas massivas capazes de enriquecer o meio interestelar mais rapidamente.

Para t.estar os diferentes cenârios de evoluçâo qufmica, tomamos aqui o modelo mais simples possfvel. A ideia é testai' se as nossas curvas de Z*{t) sâo compatfveis com uni modelo de evoluçâo qufmica simples de caixa fechada, com metalicidade inicial nula, e usando a aproximaçâo de reciclagem instantânea e yield (y) contante.

Para esse modelo, a fraçâo de massa em metais dentro de estrelas para a idade t é dado por (ver equaçâo 35 de Edmunds 1990):

Z*(t) = (i-m (5.1)
onde f(t) é a fraçâo em massa do gâs e y é a razâo entre a massa dos metais expelidos por geraçâo de estrelas e a massa restante em estrelas (o "yield" ). A fraçâo em massa dos gâs é definida como:

m = 1 -(lj-) SFR(O d (5.2)
onde Mt é a massa total da galâxia em estrelas e gâs, R é a fraçâo de gâs retornada devido a ventos estelares e supernovas e tQ é a idade da populaçâo mais vellia na base.

Na Seçâo 4.4, escrevemos a SFR. em funçâo de a fraçâo de massa com que uma SSP contribui para uma galâxia segundo o ajuste com o starlight. Podemos entâo reescrever a Equaçâo 4.2 da seguinte maneira:

SFR(t)At = MX(t), (5.3) 
onde M* é a massa total convertida em estrelas, de forma que Mf = M*/( 1 -R).

Usando as Equaçôes 3.2 e 5.3, f(t) pode ser reescrito como uma funçâo da funçâo cumulativa de massa r)Z(t): A evoluçâo da metalicidade estelar derivada pela sfntese pode ser bem reproduzida pelo modelo simples de caixa fechada para toda a faixa de massa e tipos de galâxias analisados nesse estudo. Isso implica que as historias de formaçâo estelar devem ser o principal parâmetro ffsico que dita a forma da relaçâo M-Z para essas galâxias. Outros fatores, como ventos, podem ter uni efeito secundârio nessa relaçâo. LINERs "reais" e as "aposentadas" sâo diferentes; por exemplo, podemos pro curai' evidências de AGNs para esses objetos, como linhas largas no otico (pela combinaçâo de espectros), variabilidade em raio-X, ou poeira no infravermelho. - Também disponfvel em versâo preprint (arXiv:astro-ph/0610815). INTRODUCTION One of the major challenges of modem astrophysics is to understand the physical processes involved in galaxy formation and évolution.

f{t) = 1 -(1 -/now)rç*M, ( 5 
| 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 r N = 52185 objects 0.1 - 0 - © CSJ £> c -0.1 tSl Oû O -0.2 - -0.3 - i J» L ilia AaJSL_JisC
I i i i i I i i i i I i » i i I i i i i I i i i i I 9 
Significant steps in this direction could be made by tracing the buildup of stellar mass and metallicity as a function of cosmic time.

One way to address this issue is through cosmologically deep surveys which map how galaxy properties change for samples at differ ent redshifts (z). Among these properties, the relation first observed by Lequeux et al. (1979) Beyond the Local Group and its neighbourhood, fossil meth ods must dérivé SFHs from integrated galaxy spectra, using either selected spectral features (Kauffmann et al. 2004;Brinchmann et al. 2004;Gallazzi et al. 2005;Nelan et al. 2005;Thomas et al. 2005;Bernardi et al. 2006) or the full spectrum (Panter, Heavens & Jimenez 2003;Heavens et al. 2004;Cid Fernandes et al. 2005;Mathis, Chariot & Brinchmann 2006;Panter et al. 2006). This approach works with much less information than is available from individual stars in galaxies, so one cannot expect similar levels of accuracy. On the other hand, integrated spectra of galaxies are now This Letter shows that we can actually do much better. To illustrate this we hâve culled a sample of star-forming (SF) galaxies from the SDSS (Section 2), to which we apply our synthesis method and dérivé mass-assembly and Chemical enrichment historiés (Sec tion 3). The évolution of stellar mass and metallicity is then studied by grouping galaxies with similar nebular abundances or mass (Section 4). Details of this analysis and cross-checks on the results reported here are presented in a companion paper (Asari et al., in préparation, hereafter A07). interval is chopped on to six logarithmic bins of width A log Zneb = 0.13 dex, except for the first one which is twice as wide to include more sources. Galaxies inside these same bins will be grouped together in the analysis of star formation and Chemical enrichment historiés presented in the following sections.

N(F) = 1727 <Z_> = 1.22 0 -2 -1.5 log [NII]X6584/Ha -0.5 t-[ i i-i i | i i i i-pi-i i i-|-i-i-i-i-[-r N = 84828 (b) -]-i-i-i-i-|-i-i-i-i-|-i-r N = 175234

METHOD OF ANALYSIS

As thoroughly discussed in the literature on stellar population synthesis, the combination of noise in the data and astrophysical and mathematical degeneracies limits the amount of information about galactic historiés that can be extracted from integrated spectra.

Clearly, the 150 components of the population vector used in our fits are a greatly overdetailed description, so some a posteriori com pression of the results to a lower dimension is needed. Our previous papers (SEAGal I-IV) hâve taken this compression approach to its extreme by reducing the whole r* and Z* distributions encoded in the population vector to their first moments;

N, (log Ma/ = ^tij log t+j, Given this success, it is natural to explore less drastic compressions of the population vector.

/'=i v* <Z*)a/ = ]^Mzm, (2) 
One way to compress the population vector but retain the âge in formation is to marginalize over the Z* j distribution and compute the total mass converted into stars as a function of time:

<(M = *e}' (4) 
where pc-dénotés the fraction of the total mass converted into stars (M*) that is associated with the j,h base âge. The superscript 'c' To track the évolution of the stellar metallicity we compute the total mass in metals inside stars as a function of r* and divide it by 1 Because our main goal is to compare the intrinsic évolution of galaxies, throughout this Letter we consider âges in the context of stellar évolution alone, i.e. as if ail galaxies were observed at z = 0, such that /* is both a stellar âge and a lookback time.

the total mass inside stars at the same time. This yields the timedependent mean stellar metallicity:

Z7(M= lxid*)ZKj ( 5)

The computation of Pj(t*) takes into account that the total mass inside stars changes with time, both due to the SFH and to the timedependence of the returned mass. Accordingly, only populations older than r* enter the définition of Pj(r*), as younger ones had not formed yet, i.e. Pj(t+) = 0 for y < f *.

RESULTS

The explicit time-dependence in équations ( 4) and ( 5) provides a description of the star formation and Chemical historiés of galaxies which goes well beyond that obtained with the mere first moments of the âge and metallicity distributions. Given the numerous de generacies which affect population synthesis of galaxy spectra one might be skeptical as to how useful they actually are. Simulations addressing this issue will be presented elsewhere. In what follows

we report what is obtained in practice.

Galaxy évolution in Zneb-bins

Our strategy is to chop the SF-wing into bins in Zneb and study the évolution of galaxies from the statistics of ??*(/*) and Z+(M in each bin. The fact that Zneb correlates strongly with several physical and observational properties guarantees that galaxies inside each Zneb bin are intrinsically similar, which in turn implies that statistics over such bins is meaningful. that, although most of the mass is assembled at early times (r* > 1 Gyr), galaxies at the tip of the SF-wing hâve larger star formation rates (SFR) per unit mass in the recent past than those at the bottom of the wing. This trend can be quantified in tenns of the ratio of the mean SFR in the last 60 Myr to that over the galaxy's history, i.e.

the so-called Scalo's b parameter (e.g. Brinchmann et al. 2004). We M+ (Zaritsky, Gonzalez & Zabludoff 2004;Tremonti et al. 2004; Lee et al. 2006), this behaviour is ultimately a signature of galaxy downsizing (Cowie et al. 1996).

The Chemical évolution of stars in SF galaxies is shown in 

Galaxy évolution in mass bins

While Fig. 2 illustrâtes the strong link between the stellar mass as sembly and Chemical évolution historiés of galaxies and the présent ionized ISM métal content, Zneb is not the cause but the product of galaxy évolution, so its use as an independent variable, although pedagogically useful, is questionable. If one had to choose a single quantity as the main driver of galaxy évolution it would surely be the mass (e.g. Tinsley 1968;Larson 1974;[START_REF] Pagel | Nucleosynthesis and Chemical Evolution of Galaxies[END_REF].

We hâve thus repeated the analysis above, binning galaxies by M * instead of Zneb. Results are shown in Fig. 3 in M* or Zneb samples somewhat different populations of galaxies.

On the whole, however. this alternative grouping scheme leads to the same general scénario outlined above, with low galaxies being slower in their mass and Chemical build-up than more massive ones. A caveat is that A7* is not a perfect tracer of the depth of the potential well, as it does not account for the mass in gas nor dark matter. In any case, the signature of galaxy downsizing is évident in Our analysis shows that the présent ISjM abundance is intimately connected to the past star formation and Chemical enrichment history of a galaxy. Whereas this is intuitively obvious from a physical point of view, it is not obvious at ail that one should be able to recover this behaviour from population synthesis analysis of integrated galaxy spectra of SDSS-like quality. We emphasize that, unlike previous studies such as that by Bica (1988), our fits im pose no a priori Chemical évolution constraints whatsoever upon the mixture of r* j and Z* y values in the base. Given the absence of constraints and the fact that the fits are completely independent of the émission line data used to compute Zneb, the organization of galaxies with different Zllcb on to the clear and systematically distinct stellar Chemical enrichment patterns seen in Fig. 2 Our main achievements may be summarized as follows. (1) A conventional corrélation anal ysis is performed to study how global properties relate to Zncb, leading to the confirmation of previously known relations, such as those between Zneb and galaxy luminosity, mass, dust con tent, mean stellar metallicity and mean stellar âge. (2) A simple formalism which compresses the results of the synthesis while at the same time yielding time-dependent star formation rates (SFR) and mass assembly historiés is presented. (3) A comparison of the current SFR derived from the population synthesis with that obtained from Ha shows that these independent esti mators agréé very well, with a scatter of a factor of 2. An important corollary of this finding is that we now hâve a way to estimate SFR in galaxies hosting active galactic nuclei, where the Ha method cannot be applied. ( 4) Fully time-dependent SFHs were derived for ail galaxies, and then averaged over six Zneb bins spanning the entire SF wing in the [O m]/H/?-[N n]/Ha diagram. ( 5) We find that SFHs vary systematically along the SF sequence. Though ail SF galaxies formed the bulk of their stellar mass over 1 Gyr ago, low-Zneb Systems evolve at a slower pace and are currently forming stars at a much higher relative rate. Galaxies at the tip of the SF wing hâve current spécifie SFRs about two orders of magnitude larger than the metal-rich galaxies at its bottom. ( 6) At any given time, the distribution of spécifie SFRs for galaxies within a Zneb bin is broad and approximately lognormal. ( 7) The whole study was repeated grouping galaxies within bins of stellar mass and surface mass density, both of which are more fundamental drivers of SFH. Given the existence of strong Zneb -M* -Z* relations, the overall picture described above remains valid. Thus, low-M* (low-E*) Systems are the ones which evolve slower, with current spécifie SFRs much larger than more massive (dense) galaxies. (8) This overall pattern of SFHs as a function of Zneb,M* or E* is robust against changes in sélection criteria, choice of evolutionary synthesis models for the spectral fîts, and differential extinction effects.

Key words: galaxies: évolution -galaxies: statistics -galaxies: stellar content.

INTRODUCTION

The Sloan Digital Sky Survey (SDSS; York et al. 2000), with its homogeneous spectroscopic and photometric data on hundreds of thousands of galaxies has revolutionized our perception of the world of galaxies in the local Universe. The enormous amount of objects allowed one to reveal trends that had not been suspected before. For example, while is was known since the work of Baldwin, Phillips & Terlevich (1981) that objects ionized by massive stars and ac tive galactic nuclei (AGN) live in different zones of emission-line ratios diagrams, the fact that emission-line galaxies are distributed in two well-defined wings (Kauffmann et al. 2003c) The left-hand wing of the BPT diagram can be understood as a sequence in metallicity of normal star-forming (SF) galaxies. The present-day nebular metallicity (Z"eb) of a galaxy is intimately connected with its past star formation history (SFH). The main goal of this paper is to explore this link.The focus of many of the pioneering studies of SFH of SF galaxies was instead the variation of SFH with the Hubble type. [START_REF] Searle | [END_REF], for instance, assumed a simple model for the SFH and calculated UBV colours for simulated galaxies. Comparing simulated and observed colours, they concluded that morphological type alone does not explain the différences in SFH, proposing that the galaxy mass should also be used as a tracer of star formation. Gallagher III, Hunter & Tutukov (1984) introduced a way to study the star formation rates (SFR) in three different epochs of a galaxy's history. In order to achieve such time résolution, manifold indices were used: H i observations, dynamical masses, fî-band and Ha luminosities, and UBV colours. Sandage (1986) applied some of the techniques presented by Gallagher III et al. (1984) to investigate différences in SFH along the Hubble sequence. In the same vein, Kennicutt, Tamblyn & Congdon (1994) derived the SFR for SF ob jects from the Ha luminosity and UBV colours, and found that the SFH différences for galaxies of the same Hubble type hâve a much stronger relation with their dise than with their bulge. Gavazzi et al. (2002) measured the présent and past SFRs of late-type galaxies in nearby clusters from Ha imaging and near-infrared observations, and also derived the global gas content from H i and CO observa tions. Most of these studies had to rely on many different indices and observations in order to measure an instantaneous SFR, or at most a 2-3 âge resolution SFH for SF galaxies. Bica (1988) introduced a method to reconstruct SFHs in greater detail by mixing the properties of a base of star clusters of various âges (r*) and metallicities (Z*). In its original implémentation, this method used a set of 5-8 absorption-line équivalent widths as ob servables, a grid of clusters arranged in 35 combinations of r* and Z*, and a simple parameter space exploration technique limited to paths through the r* -Z* plane constrained by Chemical évolution arguments (see also Schmidt et al. 1991;Bica, Alloin & Schmitt 1994;Cid Femandes et al. 2001). Its application to nuclear spectra of nearby galaxies of different types revealed systematic variations of the SFH along the Hubble sequence.

For over a decade the most attractive feature of Bica's method was its use of observed cluster properties, empirically bypassing the limitations of evolutionary synthesis models, which until recently predicted the évolution of stellar Systems at spectral resolutions much lower than the data. This is no longer a problem. Medium and high spectral resolution stellar libraries, as well as updates in evo lutionary tracks hâve been incorporated into evolutionary synthesis models in the past few years. The current status of these models and their ingrédients is amply discussed in the proceedings of the IAU Symposium 241 (Vazdekis & Peletier 2007).

These advances spurred the development of SFH recovery methods which combine the non-parametric mixture approach of empirical population synthesis with the ambitious goal of fitting galaxy spectra on a pixel-by-pixel basis using a base constructed with this new génération of evolutionary synthesis models. Methods based on spectral indices hâve also benefited from these new models, and produced an impressive collection of results (e.g. Kauffmann et al. 2003a,b,c;Brinchmann et al. 2004;Gallazzi et al. 2005). However, current implémentations of these methods do not reconstruct detailed SFHs, although they do constrain it, providing estimâtes of properties such as mass-to-light ratios, mean stellar âge, fraction of mass formed in recent bursts and ratio of présent to past SFR.

The first SFHs derived from full spectral fits of SDSS galax bly historiés. It is shown that there is a continuity in the évolution properties of galaxies according to their présent properties. Massive galaxies formed most of their stars very early and quickly reached the high stellar metallicities they hâve today, whereas low-mass (metal-poor) galaxies evolve slower. These findings are in agreement with recent studies of the mass assembly of larges samples of galaxies through the fossil record of their stellar populations (Heavens et al. 2004), and of studies of the Z* distribution in small samples of galaxies (e.g. Skillman, Côté & Miller 2003, and référ encés therein), but the generality of the resuit applied to the entire population of SF galaxies was shown for the fïrst time.

In the présent paper, we aim at a more complété view of the properties of SF galaxies, and their variations along the SF sequence in the BPT diagram, improving and expanding upon the results only briefly sketched in SEAGal V. In particular, we discuss in depth time-averaged values of quantités such as the SFR and the SFR per unit mass, as well as their explicit time dependence.

The paper is organized as follows. Section 2 describes our par ent sample and explains our criteria to define normal SF galaxies.

This section also explains how we deal with extinction and how we estimate the nebular metallicity. In Section 3, we discuss the global properties of galaxies along the SF sequence in the BPT diagram. In Section 4, we then proceed to explain our formalism to uncover the explicit time dependence of such quantifies as the SFR. In Section 5, we show that the current SFR as estimated by the most cominonly indicator -the Ha luminosity -compares with that obtained from our stellar population synthesis analysis. In Section 6, we analyse the SFH along the SF sequence, binning galaxies in tenns of their present-day nebular metallicity. We show that, despite the important scatter at any Z^t,, there is a clear tendency for the SFH as a function of Zneb, in that in the most metal-rich galaxies most of the stellar mass assembly occurred very fast and early on, while metal-poor Systems are currently forming stars at much higher relative rates.

We also compute mean SFHs binning the galaxies with respect to the stellar mass and the surface mass density, which are expected to better express the causes of the évolution of galaxies. Section 7 discusses possible sélection effects and othercaveats. Finally, Section 8 summarizes our main results.

DATA

The data analysed in this work were extracted from the SDSS DR5 (Adelman- McCarthy et al. 2007). This release contains data for 582 471 objects spectroscopically classified as galaxies, from which we hâve found~l .6 percent of duplicates, i.e. objects with multiple spectroscopic information in the parent galaxy catalogue.

From the remaining 573 141 objects we hâve selected our par ent sample adopting the following sélection criteria: 14.5 ^mr 17.77 and z ^0.002. The magnitude range cornes from the défini tion of the Main Galaxy Sample, whereas the lower redshift limit is used to avoid inclusion of intragalactic sources. The resulting sam ple contains 476 931 galaxies, which corresponds to about 82 per cent of ail galaxies with spectroscopic data gathered by SDSS and publicly available in the DR5. These limits imply a réduction by ~17 per cent in the sample studied in SEAGal V. 1994' evolutionary tracks (Alongi et al. 1993;Bressan et al. 1993;Fagotto et al. 1994a,b;Girardi et al. 1996), and STELIB library (Le Borgne et al. 2003). The base components comprise 25 âges between t*j = 1 Myrand 18Gyr,and six metallicities, from Z+j =0.005 to 2.5 Zq . Bad pixels, émission fines and the Na D doublet are masked and left out of the fits. The emission-line masks were constructed in a galaxy-by-galaxy basis, following the methodology outlined in SEAGal II and Asari (2006). starlight outputs several physical properties, such as the present-day stellar mass, stellar extinction, mean stellar âges, mean metallicities as well as full time-dependent star formation and Chemical évolution historiés, which will be used in our analysis. Section 4.1 describes aspects of the code relevant to this work.

Fitting half a million spectra represented a massive computational effort, carried out in a network of over 100 computers spread over three continents and controlled by a specially designed PHP code.

This huge data base of spectral fits and related products, as well as starlight itself, are publicly available in a Virtual Observatory environment at www.starlight.ufsc.br (see Cid Femandes et al., in préparation).

Examples of the spectral fits obtained for five SF galaxies are shown in Fig. 1. We hâve ordered the galaxies according to their nebular metallicity (Zneb), as defined in Section 2.4, to illustrate how spectral characteristics change along the Zneb sequence. Metalpoor galaxies (top) show blue spectra and strong émission fines in comparison to the redder spectra and weaker émission fines of galaxies with a metal-rich ISM (bottom).The middle panels in Fig. 1 show the fractional contribution to the total flux at A0 = 4020 Â of SSP of âge t*. These panels show that young stellar populations make a dominant contribution in galaxies with low Zncb, whereas at higher nebular metallicities a richer blend of stellar âges is présent.

22 Emission lines 22.1 General procedure Emission fines were measured fitting Gaussians to the residual spec tra obtained after subtraction of the stellar light using an updated version of the line-fitting code described in SEAGal III. The main transitions used in this study are Hfi, (Oin]A5007, Ha and |Nii]à6584. In the next section these fines are used to define the subsample of SF galaxies which will be studied in this paper.

The spécial case of HŴ

e find that a zero-level residual continuum is adéquate to fit the émission fines, except for H/3. Inspection of the spectral fits shows that the synthetic spectrum is often overestimated in the continuum around Hfi, creating a broad, ~200 Â wide absorption trough in the residual spectrum. This problem, which can hardly be noticed in 

2006).1

Though this is a low-amplitude mismatch (équivalent width ~3 Â spread over ~200 Â), it makes Wfi sit in a région of négative residual flux, so assuming a zero-level continuum when fitting a Gaussian may chop the base of the émission line, leading to an underestimation of its flux. To evaluate the magnitude of this effect we hâve repeated the H fi fits, this time adjusting the continuum level from two side bands (4770-4830 and 4890-4910 Â). On average, the new flux measurements are 2 per cent larger than the ones with the continuum fixed at zéro. The différence increases to typically 4 per cent for objects with W< 5 Â, but signal-to-noise ratio (S/N) >3 in the line, and 7 per cent for WHjg < 2 Â. Noise in the side bands introduces uncertainties in the measurement of the flux, but at least it removes the systematic effect described above, so the new measurements should be considered as more accurate on aver age. We adopt these new H/l measurements throughout this work.

Using the zero-continuum measurements changes some of the quan-1 We thank Drs Enrique Perez,Miguel Cervino and Rosa Gonzalez-Delgado for valuable help on this issue.

titative results reported in this paper minimally, with no impact on our general conclusions. Liang et al. 2006;Mollâ et al. 2006).

Our sample of SF galaxies is composed by objects which are below the line separating normal SF galaxies and AGN hosts proposed in SEAGal III. We hâve imposed a lower limit of 3 in S/N on the four fines in the BPT diagram, and S/N ^10 in the 4730^4780 Â continuum to constitute our main sample (the SF sample). The 82302 galaxies composing the SF sample are shown in Fig. 2 of 20 or better. A total of 17 142 sources satisfy these criteria.

Fig. 3 shows the distributions of observational and physical properties for the samples. Naturally, the SFhq sample is skewed towards doser and brighter galaxies with respect to the SF sample, but in terms of physical properties such as stellar mass, mean âge and nebular metallicity the two samples are similar.

Nebular metallicity estimate

It is well known that the SF wing in the BPT diagram is a sequence in nebular metallicity (SEAGal III and references therein), which we quantify by the oxygen abundance obtained through the 03N2 =

[Om]5007/[N n]6583 index as calibrated by Stasiriska (2006):

log Zneb = log = -0.14 -0.25 log 03N2,

where we hâve adopted (O/H)© = 4.9 x 10~4 (Allende Prieto.

Lambert & Asplund 2001).

We hâve chosen to use the 03N2 indicator to estimate the average oxygen abundance in the ISM of SF galaxies mainly because it is a single-valued indicator and it can be easily related to the position of galaxies in the classical BPT diagram. However. this indicator is affected by the presence of diffuse ionized gas in galaxies and by the fact that the N/O ratio dépends on the way Chemical évolution proceeded (Chiappini, Romano & Matteucci 2003). In addition, for the lowest metallicity galaxies, 03N2 is not sensitive to O/H anymore, as a wide range of metallicities correspond to the same value of 03N2, as can be seen in fig. 3 Other metallicity estimâtes hâve been used for galaxies. For ex ample, Tremonti et al. (2004) obtained the nebular metallicities by comparing the observed line ratios with a large data base of photoionization models. While a priori appealing, this method is not devoid of problems. as shown by Yin et al. (2007). There is a sys tematic offset of -0.28 dex and an rms of 0.09 between our nebular metallicities and theirs. Their method also yields a larger range of values for Zneb. For the SF sample. their calibration covers from Zneb = 0.78 to 2.70 Z© for the 5-95 percentile ranges, whereas our calibration covers from 0.47 to 1.13 Z© for the same percentile ranges.

The calibration by Pettini & Pagel (2004) is more similar to our own. There is a slight offset of -0.04 dex with respect to our cal ibration and the dispersion for the SF sample is 0.03 dex. Their calibration also stretches a little the Zneb ranges: Zneb = 0.45 to 1.36 Z© for the 5-95 percentile ranges.

Although we believe that our calibration is likely more reliable, we hâve performed ail the computations in this paper also with the Tremonti et al. (2004) and Pettini & Pagel (2004) calibrations.

While the results differ in absolute scales, the qualitative conclusions remain identical.

ZnCb bins

As seen above, both physical and mathematical motivations make Our SF sample spans the Zneb = 0.2-1.6 Zq range from the tip of the SF wing to its bottom. In Fig. 2 this interval is chopped into six bins of width AlogZneb = 0.13 dex, exceptforthe one of lowest metallicity which is twice as wide to include more sources. Table 1 lists some properties of galaxies in each of these bins, which are hereafter labelled A-F. Galaxies inside these bins will be grouped together in the analysis of star formation presented in Section 6.

Note that the bias in the détermination of Z"eb from O3N1 at low metallicities has no conséquence for our study, since almost ail the objects from bin A remain in this bin.

2J5 Extinctions

Stellar extinction

As explained in SEAGal I, starlight also returns an estimate of the stellar visual extinction, A*, modelled as due to a foreground dust screen. This is obviously a simplification of a complex problem (Witt, Thronson & Capuano 1992), so that Av should be called a dust atténuation parameter instead of extinction, although we do not make this distinction. Previous papers in this sériés hâve used the Cardelli et al. (1989, hereafter CCM), reddening law, with Rv -3.1. In order to probe different recipes for dust atténuation, we hâve selected 1000 galaxies at random from the SFhq sample and fitted them with four other functions: the starburst atténuation law of Calzetti, Kinney & Storchi-Bergmann (1994), the Small Magellanic Cloud (SMC) and Large Magellanic Cloud (LMC) curves from Gordon et al. (2003) and the À-0-7 law used by Kauffmann et al. (2003a).

We find that the quality of the spectral fits remains practically unchanged with any of these five laws. Averaging over ail galaxies the SMC law yields slightly better x2 values, followed closely by the Calzetti, À"0-7, LMC and CCM, in this order. As expected, these These différences, which are already small, should be considered upper limits, since they are derived from the most extincted objects.

Somewhat larger différences are found when using the SMC and A.-0,7 laws. For instance, compared to the Calzetti law, the SMC law produces mean stellar âges 0.15 dex younger and masses 0.07 dex smaller, again for the most extincted objects.

We hâve opted to use the Calzetti law in our starlight fits and emission-line analysis. The reasons for this choice are twofold: [START_REF] Fernandes | Alternative diagnostic diagrams and the "forgotten" population of weak line galaxies in the sdss[END_REF] the Calzetti law yields physical properties intermediate between the SMC and CCM laws and ( 2) this law was built up on the basis of integrated observations of SF galaxies, similar to the ones studied in this paper. In any case, the experiments reported above show that this choice has little impact upon the results.

starlight also allows for population-dependent extinctions. Al though tests with these same 1000 galaxies sample show that in general one obtains larger A* for young populations, as expected, simulations show that, as also expected, this more realistic modelling of dust effects is plagued by degeneracies which render the results unreliable (see also Panter et al. 2007, SEAGal V). We therefore stick to our simpler but more robust single A* model.

Nebular extinction

The nebular V-band extinction was computed from the Ha/H/3 ratio assuming a Calzetti et al. (1994) law:

1 neb 1V 7.96 log (Ha/Hl)obs (Ha/H/Om, '

where (Ha/H/3)obs and (Ho'/H/J)in, are the observed and intrinsic ratio, respectively. Instead of assuming a constant value, we take into account the metallicity dependence of (Ha/H/J)int, which varies between 2.80 and 2.99 for Z"eb in the 0.1-2.5 Zq range, as found from the photoionization inodels in SEAGal III.2

We obtain the intrinsic ratio as follows. We start by assuming (Ha/H/Ü)int = 2.86, from which we dérivé a first guess for A"cb.

We extinction curve), and in agreement with detailed studies of nearby SF galaxies (Calzetti et al. 1994). We also find that the différence between nebular and stellar extinctions increases systematically as the mean âge of the stellar population increases.

Given the spatial association of the line-emitting gas and the ioni zing populations, these results ultimately imply a breakdown of our simple single-A* modelling. In fact, starlight experiments with population-dependent A* point in the same direction, i.e., the need to allow young populations to suffer more extinction than older ones. To evaluate to which extent this simplification affects the re sults reported in this paper, Section 73 présents experiments where the extinction of f* ^10 Myr components is set according to the empirical relation Ayeb (A*) found above.

Interstellar absorption as traced by the Na D doublet

The most conspicuous spectroscopic feature of the cold ISM in the optical range is the Na D doublet at XÀ5890,5896 Â. For a constant gas to dust ratio, the strength of this feature, which measures the amount of cold gas in front of the stars, should correlate with A*, as found for far-infrared bright starburst galaxies (Heckman et al. 2000). To perform this test for our sample, we first measure the flux of the Na D doublet in the residual spectrum, integrating from 5883 to 5903 Â. We thus remove the stellar component of this feature, which is also présent in stellar atmosphères, particularly late-type stars (Jacoby, Hunter & Christian 1984;Bica, Barbuy & Ortolani 1991). In principle, this is a more précisé procedure than estimating the stellar Na D from its relation to other stellar absorption lines (Heckman et al. 2000;Schwartz & Martin 2004), but since the Na D window was masked in ail fits (precisely because of its possible contamination by ISM absorption), the stellar Na D predicted by the fits rely entirely on other wavelengths, so in practice this is also an approximate correction. The residual flux is then divided by the continuum in this range (defined as the médian synthetic flux in the 5800-5880 plus 5906-5986 Â Windows), yielding the excess équivalent width AlVNaD, which says how much stronger (more négative) the Na D feature is in the data with respect to the models. Regardless of such details, the discovery of this astrophysically expected corrélation strengthens the confidence in our analysis. Furthermore, it opens the interesting prospect of measuring the gas-todust ratio and studying its relation with ail other galaxy properties at hand, from nebular metallicities to SFHs. This goes beyond the scope of the présent paper, so we defer a detailed analysis to a future communication.

CORRELATIONS WITH NEBULAR

METALLICITY

Galaxy properties change substantially from the tip of the SF wing, where small, metal-poor Hll-galaxy-like objects live, to its bottom, populated by massive, luminous galaxies with large bulge-to-disc ratios and rich in metals (Kennicutt 1998). The simplest way to investigate these systematic trends is to correlate various properties with the nebular metallicity (e.g. Brinchmann et al. 2004;Tremonti et al. 2004).

In this section we correlate Zneb with both observed and physical properties extracted from our stellar population fits. This traditional analysis, based on current or time-averaged properties, helps the interprétation of the more detailed study of time-dependent SFHs presented in the next sections. In fact, this is the single purpose of this section. Since most of the results reported in this section are already known or indirectly deducible from previous work, we will just skim through these corrélations.

Fig. 5(a) shows Z,,eb against absolute r-band magnitude.This is the luminosity-nebular metallicity relation, previously studied by many authors, and interpreted in terms of a mass-metallicity relation. but in units of 2 x 108 L©, such that it can also be read as an estimate of the current SFR in Mq yr-1 (Section 5).

different methods to estimate both stellar and nebular abundances.

Even though we express both quantifies in solar units, these two metallicities are derived by such radically different means that, as discussed in SEAGal V, they should not be compared in quantitative terms.

The relation between the mean stellar âge (logr*)L and Znei,, shown in Fig. 5(f), reflects the fact that young stars hâve a larger share of the light output at the tip of the SF wing than at its bottom, where old populations hâve a greater weight. Metal-rich SF galaxies thus hâve a more continuous SF history than metal-poor ones, which are often dominated (in light, but not mass) by the latest génération of stars (e.g. Corbin et al. 2006). This is another way to look at métal 1 icity-age trend, discussed previously in the analysis of Fig. 1. Ultimately, this relation represents a summary of Chemical évolution, in the sense that more evolved Systems hâve a more enriched ISM. In a related vein, Fig. 5(g) shows how the ratio of current to mean past SFR (defined in Section 53) varies along the metallicity sequence of SF galaxies. This indicates that the lower metallicity galaxies are slower in forming stars. When one considers the mass-metallicity relation (Fig. 5b), this is just another way of looking at the downsizing effect (Heavens et al. 2004;Thomas et al. 2005; SEAGal II). Finally, Fig. 5(h) shows the relation between reddening corrected Ha luminosity and ZnCh. The y-axis is given in units such that the values correspond to the current SFR in Mq yr"' (see Section 5). The corrélation, although statistically unquestionable, has a large scatter. This implies that galax ies in the 6-Zneb bins defined in Fig. 2 hâve heavily overlapping current SFRs. Section 6 présents independent confirmation of this fact.

As expected, ail corrélations discussed above are also présent for the SFhq subsample. For most they are in fact somewhat stronger, whereas for sainples defined with less stringent criteria the corréla tion strengths weaken, indicating that noise in the data is responsible for part of the scatter in these relations. Finally, as is widely known and can be deduced from Fig. 5 itself, there are many interrelations between galaxy properties. Our use of Z™* as the 'independent' vari able axis in Fig. 5 is not meant to indicate that Zneb is the underlying cause of the corrélations; it simply reflects our interest in mapping physical properties of galaxies along the SF wing of the seagull in the BPT diagram.

METHODS TO INVESTIGATE STAR FORMATION HISTORIES

The main goal of this paper is to study how the SFH varies among SF galaxies. Most other investigations in this same line used absorption, émission or continuum spectral indices such as the 4000-Â break, the absorption, the K, G and Mg bands, or the Ha luminosity and équivalent widths to characterize the SFH (e.g. Raimann et al. 2000;Kong et al. 2003;Cid Fernandes, Leâo & Lacerda 2003;Kauffmann et al. 2003b;Brinchmann et al. 2004;Westera et al. 2004). Our approach, instead, is to infer the SFH from detailed pixel-by-pixel fits to the full observed spectrum, thus incorporating ail available information.

Whereas our previous work concentrated on the first moments of the âge and metallicity distributions, here we présent some basic formalism towards a robust description of SFHs as a function of time.

From the point of view of methodology, these may be regarded as 'second-order' products. Astrophysically, however, recovering the SFH of galaxies is of prime importance. SEAGal V presented our first results in this direction, including empirically derived timedependent mean stellar metallicities. In this section we expand upon these results, exploring new ways to handle the output of the synthesis, focusing of the SFHs. Not surprisingly, the 150 components of the population vector (jc) are highly degenerate due to noise, and astrophysical plus mathematical degeneracies, as confirmed by simulations in Cid Fernandes et al. (2004) and SEAGal 1. These same simulations, however, proved that compressed versions of the population vector are well recovered by the method.

Different compression approaches exist among spectral synthesis codes. In moped (Heavens, Jimenez & Lavah 2000;Reichardt, Jimenez & Heavens 2001 ;Panter et al. 2003;Heavens et al. 2004; Panter et al. 2007), for instance, compression is done a priori, replacing the full spectrum by a set of numbers associated to each of the Ni, + 1 parameters (the mass fractions and metallicities in several time-bins plus a dust parameter). STECMAP (Ocvirk et al. 2006) performs a compression by requiring the resulting SFH to be relatively smooth. The preference for a smooth solution over a ragged one is effectively a prior, but the algorithm adjusts the degree of smoothing in a data-driven fashion, so we may call it an 'on-the-fly' compression method. The same can be said about vespa (Tojeiro et al. 2007). a new code which combines éléments from these two approaches. starlight is less sophisticated in this respect. Its only built-in compression scheme is that the final stages of the fit (after the Markov Chains reach convergence) are performed with a reduced base comprising the subset of the original Ni, populations which account for ^99 per cent of the light. For our parent sample of 573 141 galaxies the average size of this subset is Nf{ = 24 populations, while for the 82302 galaxies in the SF sample Nf{ = 41. (This différence happens because the full sample has many old, passive Systems, which require relatively few SSPs, whereas SF galaxies hâve more continuous SF régimes, thus requiring more SSPs to be adequately fit.) Compression beyond this level must be carried out a posteriori by the user. As explained in the next section, in this study we in fact compress this information into only four âge bins by smoothing the population vectors.

Previous papers in this sériés hâve taken this a posteriori com pression approach to its limit, condensing the whole âge distribution to a single number, the mean stellar âge:

N"

(log tt)L = ^xj los

7=1

where the subscript L dénotés a light-weighted average. Massweighted averages are readily obtained replacing x by the mass fraction vector p. Similarly, stellar metallicities were only studied in terms of their mass-weighted mean value:

N" = (4) 

7=1

Simulations show that both of these quantities hâve small uncertainties and essentially no bias. Regarding practical applications, these first moments proved useful in the study of several astrophysical relations, some of which hâve just been presented in Section 3 (see Fig. 5). Notwithstanding their simplicity, robustness and usefulness, these averages throw away ail time-dependent information contained in the population vector, thus hindering more detailed studies of galaxy évolution. In what follows we explore novel ways to deal with the population vector which circumvent this limitation.

Star formation rate as a function of time

One alternative to characterize higher moments of the SFH is to bin jc on to âge groups, a strategy that goes back to B ica (1998; see also Schmidt et al. 1991;Cid Femandes et al. 2001).Though useful, this approach introduces the need to define bin limits, and produces a discontinuous description of the SFH.

A method which circumvents these disadvantages is to work with a smoothed version of the population vector. We do this by applying a Gaussian filter in logt*, with a full width at half-maximum (FWHM) of 1 dex. Given that our base spans approximately four orders of magnitude in f*, this heavy smoothing is équivalent to a description in terms of ~4 âge groups, but with the advantage that xs can be sampled continuously in logf*. This approach is analogous to smoothing a noisy high-resolution spectrum to one of lower res olution, but whose large-scale features (colours, in this analogy) are more robust. From the results in SEAGal I, where it was shown that three âge groups are reliably recovered, we expect this smoothing strategy to be a robust one. Furthermore, averaging over a large num ber of objects minimizes the effects of uncertainties in the smoothed SFH for individual galaxies.

Technically, the issue of âge resolution in population synthesis is a complex one. Ocvirk et al. (2006), for instance, find that bursts must be separated by about 0.8 dex in log f* to be well distinguished from one another with their SFH inversion method. Considering that the âge range spanned by our base is 4.2 dex wide, one obtains five 'independent' time-bins. This number is similar to that (six bins) used by Mathis et al. (2006) to descri be the SFH of SDSS galaxies (covering a wider À range than those simulated by Ocvirk et al. but at lower S/N) with a variant of the moped code. Up to 12 timebins were used in other applications of moped. Panter et al. (2007) argue that this may be a little too ambitious for individual galaxies, but uncertainties in this overparametrized description average out in applications to large samples. Tojeiro et al. (2007) hâve a useful discussion on the number of parameters that can be recovered with synthesis methods. By using their vespa code and calculating the number of parameters on the fly for each individual object, they find that typically 2-8 parameters can be robustly recovered for SDSS spectra. Hence, despite the complexity of the issue, there seems to be some general consensus that the âge resolution which can be achieved in practice is somewhere between 0.5 and 1 dex, so our choice of smoothing length is clearly on the conservative side.

A further advantage of this continuous description of the SFH is that it allows a straightforward dérivation of an SFR. Recall that we describe a galaxy's évolution in terms of a succession of instantaneous bursts, so an SFR is not technically definable unless 

where M* is the total mass converted to stars over the galaxy history until ti, = 0. and pcs (ti,) is the fraction of this mass in the bin.3

We can also define the time-dependent spécifie SFR:

SSFR(ri) = 1 d Af°(ri) Ml d/* loge jffiri) t* A log ri (6) 
which measures the pace at which star formation proceeds with respect to the mass already converted into stars. This is a better quantity to use when averaging the SFH over many objects, since it removes the absolute mass-scale dependence of équation ( 5).

3 The superscript c is introduced to distinguish from the mass still locked inside stars (Mi,), which must be corrected for the mass returned to the ISM by stellar évolution. This distinction was not necessary in previous SEAGal papers, which dealt exclusively with M* and its associated mass fraction vector (p). When computing SFRs, however, this différence must be taken sures the rate at which gas is tumed into stars of any metallicity.

(2) The upper âge limit of our base (18 Gyr) is inconsistent with our adopted cosmology, which implies a 13.5-Gyr universe. Given the uncertainties in stellar évolution, cosmology, observations and in the fits themselves, this is a merely formai inconsistency, and, in any case, components older than 13.5 Gyr can always be rebinned We sample rç* in the same logr* = 5.6-10.5 grid used to describe the évolution of the SFR, but here we operate on the original population vector, not the smoothed one. Since most of the mass as sembly happens at large r*, computing rç* with the smoothed SFHs leads to too much loss of resolution. In essence, however, ??*(r*)

and SSFR*(/*) convey the same physical information in different forms.

THE CURRENT SFR

The most widely employed method to measure the 'current' SFR is by means of the Hcr luminosity (Kennicutt 1983(Kennicutt , 1998;;Hopkins et al. 2003). We hâve just devised ways of measuring the timedependent SFR which rely exclusively on the stellar light, from which one can define a current SFR averaging over a suitably defined time interval. Before proceeding to the application of these tools to study the detailed SFHs of galaxies, this section compares these two independent methods to estimate the current SFR.

The purpose of this exercise is threefold. First, it serves as yet an other sanity check on the results of the synthesis. Secondly, it allows us to define in an objective way the ratio of current-to-past average SFR, often referred to as Scalo's b parameter (Scalo 1986), which is a useful way to summarize the SFH of galaxies (e.g. Sandage 1986;Brinchmann et al. 2004). Finally, defining and calibrating a synthesis-based measure of current SFR équivalent to that obtained with Hcr, allows one to estimate the current SFR in galaxies where Hcr is not powered exclusively by young stars. This tums out to be very useful in studies of AGN hosts (Torres-Papaqui et al., in préparation).

Current SFR from Hcr luminosity

For an SFR which is constant over times-scales of the order of the lifetime of massive ionizing stars (rion ~10 Myr), the rate of H-ionizing photons converges to

Qh = SFRA/h(IMF, Z*), (8) 
where A/h is the number of hv > 13.6 eV photons produced by a SSP of unit mass over its life (in practice, over 95 per cent of the ionizing radiation is produced in the first 10 Myr of évolution). We computed A/h by integrating the Qn(t) curves for SSPs using the tables provided by BC03. obtaining A/h = 9.12, 7.08, 6.17, 5.62, 4.47 and 3.16 x ÎO60 photons MA1 for Z* = 0.005, 0.02, 0.2, 0.4, 1 and 2.5Zq, respectively, for a Chabrier (2003) IMF between 0.1 and 100Mo.4

One in every 2.226 ionizing photons results in émission of an Hcr photon, almost independently of nebular conditions (Osterbrock & Ferland 2006). This assumes case B recombination and that no ionizing photon escapes the Hn région nor is absorbed by dust.

Adopting the Chabrier IMF and the Zq value of A/h leads to

OCD 2.226LHa _j f LHa ŜFR-=Àü^=2M°yr (9) 
This calibration is strongly dépendent on the assumed IMF and upper stellar mass limit. Given its reliance on the most massive stars, which comprise a liny fraction of ail the stars formed in a galaxy, SFRn« involves a large IMF-dependent extrapolation, and thus should be used with care.

52 Current SFR from the spectral synthesis

The SFR from spectral synthesis is based on ail stars that contribute to the visible light, and thus should be more représentative of the true SFR. We define a mean 'current' SFR from our time-dependent SFHs using équation (7) to compute the mass converted into stars in the last r years, such that

1 -??c(t) SFR*(t) = Ml - (10) 
r is the mean SFR over this period. Because of the discrète nature of our base, the function SFR.(r ) has a 'sawtooth' appearance, j umpi ng every time t crosses one of the t* j bin borders.

For the reasons discussed in Section 4.1, it is désirable to include components spanning ~1 dex in âge to obtain robust results. Since our base starts at 1 Myr, r ~10 Myr would be a reasonable choice.

This coincides with the minimum time-scale to obtain SFR.(t) es timâtes comparable to those derived from LHa, which are built upon the assumption of constant SFR over r ^~10 Myr. Our base âges in this range are t+ j = 10, 14, 25, 40 and 55 Myr.

Synthesis versus Hcr-based current SFRs

To compare the SFRs given by équations ( 9) and ( 10) we must first choose a spécifie value for r. We do this by correlating the SFR per unit area obtained with these two estimators, and seeking the value of r which yields the best corrélation. Surface densifies were used to remove the d2 factors common to both SFRs, thus avoiding distance-induced corrélations. Data for the SFhq sample was used in this calibration. Also, since L\\u refers to the émission from within the 3-arcsec SDSS fibres, M* was not extrapolated to the whole galaxy in this comparison. The rms différence between these two SFR estimators is 03 dex, corresponding to a factor of 2. We consider this an excellent agreement, given that these estimators are based on entirely different premises and independent data, and taking into account the uncertainties inhérent to both estimators. It is also reassuring that r turns out to be comparable to rion ~10 Myr, which is (by construction)

the smallest time-scale for SFRiia to be meaningful. That the scatter between SFRI!ü. and SFR* is typically just a factor of 2 can be attributed to the fact that we are dealing with integrated galaxy data, thus averaging over SF régions of different âges and emulating a globally constant SFR, which works in the direction of compati bilizing the hypothèses underlying équations ( 9) and ( 10).

With these results, we define the ratio of 'current' to mean past SFR as SFRrfr = 24.5 Myr)

SFRrfr = rG) (H)
where rG is the âge of the oldest stars in a galaxy. In practice, since the overwhelming majority of galaxies contain components as old as our base allows, the denominator is simply M* di vided by the âge of the Universe, such that b is ultimately a measure of the current spécifie SFR. This définition was used in Section 3, where it was shown that b decreases by an order of magnitude in the médian from the top to the bottom of the SF wing.

THE STAR FORMATION HISTORIES OF SF

GALAXIES

Spectral synthesis methods such as the one employed in this work hâve historically been seen with a good deal of skepticism, best epitomized by [START_REF] Searle | Stellar Popula tions[END_REF], who, when talking about the spectral synthesis of integrated stellar populations, said that 'too much has been claimed, and too few hâve been persuaded'. Persuading the reader that one can nowadays recover a decent sketch of the timedependent SFR in a galaxy from its spectrum requires convincing results. In this section we apply the new tools to describe SFHs presented above (équations 5-7) to SF galaxies in the SDSS. As shown below, the temporal dimension leads to a more detailed view of SF galaxies than that obtained with mean âges or current SFR estimâtes.

Distributions of star formation historiés

Our general strategy to explore the statistics of the sample is to group galaxies according to certain similarity criteria and dérivé mean SFHs for each group. Since ail the results presented from In what follows we will présent only mean SFHs, obtained group ing galaxies according to a subset of the available physical parameters. Whilst there is clearly more to be learned from the SFH dis tributions discussed above, this is a useful first approach to explore the intricate relations between galaxy properties and their SFHs.

62 Trends along the SF wing

We start our statistical study of galaxy SFHs grouping galaxies in the six Zneb bins defined in Section 2. As shown in Fig. 2, Zneb traces the location of a galaxy along the SF wing in the BPT diagram. this convergence seems to be at odds with the fact that galaxies with low and high Z,ieb differ by about one order of magnitude in the médian LHu (Table 1), and thus should differ by a similar factor in the recent SFR. In fact, there is no contradiction, since what needs to be considered when comparing the synthesis-based SFR with that derived from Ha is the mean SFR over scales of at least 10 Myr, and these are clearly smaller for low-ZnCb galaxies than for those of higherZnet,. As shown in Fig. 6, Ha and synthesis-based SFRs agréé very well. Ultimately, the apparent coincidence of mean SFR curves of different Z"eb is due to the fact that the relation between recent SFR and ZnCb is a relatively weak and scattered one (Fig. 5h), such that along the whole SF wing one may find galaxies that transform a similar amount of gas into stars per year. 7.5.8.5.9.5, 10.5 and 11.5, which contain 542,4057,26700,47 153,3808 galaxies, respectively. nebular abundance is not a cause, but a conséquence of galaxy évo lution. In this section we leave aside our focus on the BFT diagram and group galaxies according to properties more directly associatedto physical drivers of galaxy évolution. Two natural candidates are the mass (M+) and surface mass density (S*). Like Zncb, both A/* and E* can be considered the end product of a SFH, yet they are clearly more direct tracers of depth of the potential well and degree of gas compression, two key parameters affecting physical mechanisms which regulate galaxy évolution (Schmidt 1959;Tinsley 1980;Kennicutt 1998).

Fig. 9 shows our different représentations of the SFH of SF galax ies for five 1 dex wide mass bins centred at logA/+/MQ = 7. 5 -11.5. Given that Af* and are related, the overall evolutionary picture emerging from this plot is similar to the one obtained bin ning galaxies in Z"cb, i .e. massive galaxies assemble their stars faster than low-mass galaxies. On the whole, Fig. 9 provides a compelling visualization of galaxy downsizing.

The most noticeable différence with respect to Z"eb-binned results is on the absolute SFR curves (compare Figs 8b and9b). This différence is rooted in the fact that galaxies of similar Z,lcb span a much wider range of SFRs than galaxies of similar M*. This can be illustrated focusing on recent times, and inspecting the LH« -Zneb relation (Fig. 5h), with the understanding that LH" can be read as the current SFR (Fig. 6). Despite the statistically strong corrélation (Rs = 0.49), the typical 5-95 percentile range in LHa for a given Results grouping galaxies according to S* are presented in Fig. 10. Since S* and A/* correlate very strongly (Rs = 0.73 in our sample), the results are similar to those obtained grouping galaxies by their stellar mass. Kauffmann et al. (2003b), based on an anal ysis of two SFH-sensitive spectroscopic indices [namely D"(4000)

and HA^], propose that Z* is more directly connected to SFHs than Af*. This is not obviously so comparing Figs 9 and 10. A more detailed, multivariate analysis is needed to evaluate which is the primary driver of SFHs.

SELECTION EFFECTS AND MODELLING

CAVEATS

This section deals with the effects of sample sélection, synthesis ingrédients and model assumptions on our results.

Sélection effects

We now study to which extent the mean SFHs derived in the last section are affected by the way we hâve defined SF galaxies. We Our BPT-based sélection of SF galaxies used the dividing line proposed in SEAGal III, which is more restrictive than the empirical line proposed by Kauffmann et al. (2003c). We define the SFW sample as the 111 026 galaxies classified as SF according to the Another concem is the inclination effect. The spectra of edgeon objects are biased by the metal-poor outer parts of the galaxies, leading to an underestimation of Z^. This may lead us to place a galaxy in a lower Z"eh bin than it would if it were seen face on, possibly affecting the mean SFH in that bin. To investigate this effect we hâve defined a subsample of nearly face-on galaxies, SF°( 6842 objects), selecting by the inclination parameter, b/a ^0.9. Qualitatively, one expects that forcing a uniform A* fit to a galaxy where the young stars suffers more extinction than the others should lead to an overestimation of the âge of the young population. This This expectation is fully confirmed by these new fits. Fig. 12 shows a comparison between three Zneb bins (B, D and F) for the old and new fits of the SFhq sam pie. One sees that the new average SFR and SSFR curves are shifted by ~03 dex upwards in the f* ^107 yr range with respect to the ones obtained with a single extinction.The rearrangements in the population vector tend to be in the sense of shiftingsome lightfrom old populations tothe ^107-yrcomponents.

Not surprisingly, the properties which change most are those directly related to the strength of the young population, such as current SFR, which increases by 03 dex on average, and the mean stellar These experiments are obviously a simplification of the problem of dust distribution in galaxies, yet they suggest that the choice of the extinction modelling can hâve non-negligible effects on the derived SFH curves. On the whole, however, the qualitative pattern of SFHs as a function of Z"cb or other variables stays the same. As found i n the sample sélection studies, quantitative changes are at best équivalent to moving from one bin to the next, so that our general conclusion does not change.

SUMMARY

In this paper we hâve studied physical properties of 82302 nor mal SF galaxies from the SDSS DR5, by using results from our 5 We thank the referee for suggesting this approach. Before reviewing our main results, we highlight some aspects of this study which hâve relatively little impact upon our general con clusions, but represent significant refinements in our methodology. (ii) SF galaxies were selected according to the theoretical criterion proposed in SEAGal III, which minimizes contamination by AGN émission.

(iii) Nebular extinctions and metallicities were derived selfconsistently, allowing for the metallicity dependence of the Balmer décrément. Five different reddening laws were explored, but found to produce equally good spectral fits and relatively small différences in derived physical properties.

(iv) We hâve confirmed the strong corrélation between A"7eb and A* found in SEAGal I. We hâve also identified a strong corrélation between the strength of the ISM component of the Na D absorption doublet and the amount of dust derived from the synthesis.

(v) Different recipes for nebular metallicity estimâtes were tried.

Some of them proved not to be adéquate for this study, either because of the lack of spectral data (e.g. measures of [Arm]À7135

and [O IIIJ4363 émission lines were available for few objects), or because such calibrations were only valid in the low-Zneb régime, thus encompassing a very small fraction of objects from our sample.

Therefore, throughout our analysis we use the 03N2 index and the calibration by Stasinska (2006) to measure the nebular metallicity.

Although this is not a reliable calibrator at the lowest metallicities, it is good enough for our analysis in Zneb bins. Furthermore, it has the nice virtue of being directly related with the position of the objects in the BPT diagram.

We now summarize results related to the main goal of this paper, namely to investi gâte the SFH of galaxies along the SF wing in the BPT diagram. In practice, this means studying how SFHs change as a function of nebular metallicity, even though Zncb is more a product than a cause of galaxy évolution.

(i) We started our study with a traditional analysis, correlating

Znet, with several physical and observed properties. This analysis confions results obtained directly or indirectly in the past by other works, such as relations between the nebular metallicity and galaxy luminosity, mass, dust content, mean stellar metallicity and mean stellar âge.

(ii) Formalism towards a time-dependent analysis was then presented. Simple ways to compress the output of our stellar population synthesis code were proposed. These are based either on a posteri ori smoothing of the âge distribution, which allows the dérivation of time-dependent SFRs, or a cumulative mass assembly history.

(iii) As a first application of this time-dependent description of SFHs we computed the current SFR obtained from our spectral fits.

The resulting values of SFR* agréé very well to more traditional estimâtes based on the luminosity of Ha.The scatter between SFR* and SFRhu is just a factor of 2, despite the différences in the underlying assumptions and sensitivity to the IMF. This resuit strengthens confidence in our method, and, more importantly, opens the possibility of measuring current SFRs in galaxies hosting AGN, where Ha method does not apply.

(iv) Fully time-dependent SFHs were then derived grouping galaxies into six ZnCb bins spanning the entire SF wing of the BPT diagram. Mean SFHs for each of these bins were presented in four different représentations: (i) the smoothed population vector, je7( 6),

(ii) the SFRs SFR(/*), (iii) spécifie SFRs SSFR(/*) and (iv) massassembly historiés, /?*(t*).

(v) We found that SFHs vary systematically along the SF sé quence. Though ail galaxies assembled the bulk of their stellar mass over 1 Gyr ago, low-Z^b Systems evolve at a slower pace. Galaxies at the tip of the SF wing hâve current spécifie SFRs about two orders of magnitude larger than the metal-rich galaxies at the bottom of the BPT diagram.

(vi) At any given time, the distribution of SSFRs for galaxies within a Z"eb bin is quite broad and approximately lognormal.

(vii) We performed the same SFH study grouping galaxies by their stellar mass and surface mass density. Given the existence of Zncb-M*-E* relations, the overall picture is obtained grouping by Zncb-Thus, low M* (low E*) Systems are the ones which evolve slower, with current SSFRs much larger than more massive (dense) galaxies.

(viii) Finally, we hâve analysed a number of sélection and modelling effects that might bias our results, and show that while they may affect the derived SFHs quantitatively, the organization of SFHs Since the study by Lequeux et al. (1979), who found a luminositymetallicity relation for irregular galaxies, many papers havereported the existence of a luminosity-metallicity or mass-metallicity (M-Z) relation for ail kinds of galaxies: The more massive galaxies are also the ones with more metal-rich interstellar medium (ISM).

Considérable work has also been devoted to explaining the M-Z relation. For instance, Tremonti et al. (2004) find that outflows play an important rôle in shaping the M-Z relation, as the weaker gravita tion potential well of less massive galaxies makes them more prone to lose enriched gas via galactic winds or supernovae explosions.

Finlator & Davé (2008) daim that inflows of pristine gas could also be a Sound explanation, as the same amount of gas falling into galaxies would hâve a greater impact on the Chemical composition of less massive ones. Kôppen, Weidner & Kroupa (2007) advocate the rôle of an integrated initial mass function: massive galaxies comprise more massive clusters which may contain more massive stars that enrich the ISM faster.

The M-Z relation derived from émission lines has been found to change with redshift (Savaglio et al. 2005;Maiolino et al. 2008). The comparison between M-Z at low and high redshifts is, however, not straightforward. Different populations of galaxies live at different redshifts, and this alone can affect the dérivation of the nebular metallicity, even if one were to apply the same calibrations for both *E-mail: natalia@astro.ufsc.br tSemi-Empirical Analysis of Galaxies samples (see Stasinska 2008). Nevertheless, at leasl qualitatively, The use of Z* also allows us to explore more massive galaxies, for which the nebular metallicity estimâtes are not possible, such as AGN hosts and galaxies which hâve stopped forming stars. The R sample is composed of galaxies with little evidence of either nuclear or SF activity, which includes passive galaxies, galaxies with very weak émission fines and even some which would be traditionally classified as LINERs (see Stasinska et al. 2008). Therefore, we impose the following cuts for the R sample: S/N < 3 in the 4 BPT We hâve also divided our sample and subsamples into six present-day stellar mass bins centred in logAf"/MQ = 10.0 (A), 10.3 (B), 10.6 (C), 10.9 (D), 11.2 (E) and 11.5 (F), each one 0.30 dex wide. In the same vein as our previous works (Asari et al. 2007;Cid Fernandes et al. 2007), defining bins allows us to alleviate the problem of the mass completeness of the sample. We consider narrow mass bins in order to hâve very similar galaxies inside each bin. Although the starting sample contained a small proportion of galaxies with masses smaller than 101oMq, we do not consider them in the current analysis. Indeed, due to the redshift eut we imposed, only the most luminous low-mass objects are présent and they are not représentative of the bulk of galaxies of their mass range. 

tj>,

Another way to look at the metalficity évolution would be to consider the average metalficity of ail the stars born at the same epoch, for each value of t. This, however, is a much less robust quantity, and we found that SSP bases at hand nowadays do not allow such detailed définition.

Fig. 1 shows the médian values of rp and Z. versus t for mass bins A-F of our sample. As previous studies hâve already shown (e.g. Cid Fernandes et al. 2007;Panter et al. 2008, for SF galaxies only -here we extend the study for ail galaxies), there is a clear downsizing effect: the less massive galaxies are slower in converting their mass into stars and in building up their métal content. We now investigate the behaviour of the M*-Z* relations shown in Fig. 2 with the help of a simple model of Chemical évolution of galaxies. We assume a closed box with initial metallicity equal to zéro, and use an instantaneous recycling approximation and a constant yield. The mass-weighted average mass fraction of metals at lookback time t is then given by (Edmunds 1990, équation 35) z.it) = y-y fit) Injl/fjt))

(1-/(0)

whereft) is the gas mass fraction and y is the ratio of the mass of metals released per stellar génération to the total stellar mass locked in remnants. The gas mass fraction is given by m

= 1 - j' °SFR(f') dt', (4) 
where MT is the total mass of the galaxy (stars and gas), SFR is the star formation rate, R is the returned mass fraction due to stellar winds and supemovae and t0 is the âge of the oldest SSP in the base.

The star formation rate SFR(r) can be obtained from the results of the starlight synthesis:

SFR(f)df = MZnl(t), (5) 
where M[ is the total mass converted into stars {Mct = M*/( 1 -R)).

Using équations (1) and {5), fit) becomes simply fit) = 1 -(1 -/now)r?*(0, [START_REF] Fernandes | Uncovering the Chemical enrichment and mass-assembly historiés of star-forming galaxies[END_REF] where/now is the present-day gas mass fraction. Knowing the value of y, one can use équations ( 3) and ( 6 We therefore adopt y = 0.03 and apply our model to try to reproduce the observed metallicity évolution for ail the mass bins. shows that this simple model can reproduce the observed metallicity évolution reasonably well, given the crudeness of our approach. The discrepancies for âges above 109 yr are at least partly due to the known problems faced by spectral synthesis in this range (Gomes 2005;Koleva et al. 2008). In particular, the jumps in both ??* and Z* at ~1 Gyr are artefacts related to the spectral libraries used in BC03 (Cid Fernandes et al. 2008). New evolutionary synthesis models with more accurate stellar spectra should improve the recovery of SFHs in this range.

It is worth noting that/now is at least qualitatively in agreement with expectations: it is larger for SF galaxies than for R galaxies, and decreases with increasing M*. As a matter of fact, for ail the mass bins of the SF sample /now is compatible with the values derived from the Schmidt-Kennicutt law (bin E is compatible by construction). S galaxies also show little residual gas mass. The Chemical évolution of S galaxies is (within our resolution) similar to the évolution of non-S galaxies.

To sum up, for present-day stellar masses down to 1010 Mq and within the current limitations of astropaleontology, the closed-box model explains the present-day M-Z relation and its évolution quite well if one uses the SFR(r) obtained directly from the spectral synthesis of each galaxy. This is not to say that infall or outflow does not play a rôle. What we want to stress is that it is the SFH which is the main driver of the M-Z relation and its évolution.

SUMMARY AND DISCUSSION

Until recently, the M-Z relation in galaxies had been mostly studied using abundances derived from émission lines. Exploiting the stellar metallicities obtained with our spectral synthesis code starlight, sample. The number of galaxies in each panel is indicated at the bottom. Panels were left empty if they contained less than 250 galaxies.

Our main results are the following. We hâve found that the M-Z relation steepens and spans a wider range in both mass and métallicity at higher redshifts. The more massive galaxies show very little évolution since a lookback time of 9 Gyr. This means that they evolved very fast in a distant past, supporting the downsizing scénario. This is in agreement with other studies of the build-up of metals in S DSS galaxies as revealed by fossil records, e.g. Cid Fernandes et al. (2007) and Panter et al. (2008).

We hâve modelled the observed time évolution of the mean stel lar metallicity using a closed-box, instantaneous recycling Chemical évolution model, for galaxies of different types and masses. We find that this model is compatible with the observations. This suggests that, in the mass range studied here (from logM*/MQ = 9.85 to 11.65), the M-Z relation is mainly driven by the SFH and not by inflows or outfiows. By comparing the nebular and stellar metallicities of SDSS galaxies, Gallazzi et al. (2005) had argued that galaxies are not well approximated by closed-box models and that winds may be important. Our study suggests that those processes are not dominant, although they may play a rôle in the scaîter of the M-Z relation in our mass range.

Low-mass galaxies are not présent in our sample because of our sélection in redshift. They would require a dedicated study to be compared to the resuit found by Tremonti et al. (2004) and Garnett (2002) that such galaxies do suffer strong outfiows. As a matter of fact, Garnett finds that outfiows are important for galaxies with fi-band magnitude > -18. Only six out of the 82 662 objects in our sample hâve g-band magnitude > -18. Using the Schmidt-Kennicutt law to compute the effective yields for each galaxy (i.e.

the yield calculated from the observed nebular metallicities assuming a closed-box model), Tremonti et al. (2004) concluded that galactic winds are ubiquitous and very effective in removing metals from galaxies. This however concerns only galaxies with masses up to 10'°Mq, which are not présent in our sample. Our work emphasizes that there is a large range of galaxy masses -almost two décades -where the closed-box model seems to give a valid représentation of the Chemical évolution of galaxies.

Our models are parametrized by the present-day gas mass frac tion,/,^. We hâve found that the values of/n0w that allow us to reproduce the observed Z*(r) historiés are larger for less massive galaxies, in agreement with expectations. For galaxies of similar masses, they are larger for SF galaxies than for R galaxies. The values of/now we find for SF galaxies are compatible with those ob tained from the Schmidt-Kennicutt law. This Sound resuit supports the validity of our interprétation.

Within our uncertainties, Seyfert galaxies hâve as little gas as retired galaxies. INTRODUCTION

In the Baldwin, Phillips & Terlevich (1981, hereafter BPT) diagram which is used to isolate star forming from active galaxies, the galaxies from the Sloan Digital Sky Survey (SDSS; York et al. 2000) occupy a well-defined région, evoking the wings of a flying seagull. The left wing consists of star-forming (SF) galaxies while the right wing is attributed to galaxies with an active nucléus (Kauffmann et al. 2003;Stasinska et al. 2006). The right wing has been subdivided into an upper and a lower branch, called 'Seyfert'

and 'LINER' branch, respectively (Kewley et al. 2006). This dé ( 1994) found that, in early-type galaxies that hâve stopped form ing stars, hot post-asymptotic giant branch (post-AGB) stars and white dwarfs provide enough ionizing photons to account for the observed Hcr équivalent widths and can explain the LINER-like emission-line ratios observed in such galaxies (see also Sodré & Stasinska 1999). Taniguchi, Shioya & Murayama (2000) followed up this idea, proposing that some LINERs could be post-starburst nuclei powered by planetary nebulae central stars.

In Binette et al. ( 1994), the comparison with observations was very limited, due to the scarcity of adéquate data at that time. With the SDSS, we hâve a homogeneous data base of over half a million galaxy spectra. In addition, the techniques to extract émission Unes after modelling the stellar continuum using population synthesis (Kauffmann et al. 2003;Cid Femandes et al. 2005) now allow emission-line measurement even in objects where these Unes are difficult to detect in the raw spectra.

The computations by Binette et al. (1994) and Taniguchi et al. (2000) were based on ab initio stellar population models. Here, we directly use the populations inferred from detailed fits of the observed galaxy spectra to compute the Lyman continuum radiation and estimate its impact on the émission Unes. The radiation from the L30 G. Stasinskci et al.

post-AGB and white dwarf stars présent in old stellar populations is much harder that the one from young stars, so that galaxies that are not forming stars presently will contain hotter emission-line régions, and lie above the pure S F sequence in the BPT diagram. In the remaining of this Letter, we qualify these galaxies as

We explore the expected properties of retired galaxies in terrns of their line luminosities and their location in emission-line diagnostic diagrams.

2 DATA BASE

Sample and data processing

This work analyses data extracted from SDSS Data Release 5 (Adelman-McCarthy et al. 2007). Our parent sample is defined as the 573 141 objects spectroscopically classified as galaxies and with no duplicates. We also adopt the following criteria: 14.5 < mr < 17.77 (from the définition of the Main Galaxy Sample), a minimum signal-to-noise ratio (S/N) of 10 at ~4750 Â (to ensure a reliable stellar population analysis), z > 0.002 (to avoid intragalactic sources) and total z-band light inside the fiber >20 per cent (to reduce aperture effects).

The data are processed as in Cid Femandes et al. (2005) and Mateus et al. (2006). The stellar populations composing a galaxy are inferred through a pixel-by-pixel modelling of its continuum with our code starlight, using a base of simple stellar popula tions computed with the evolutionary population synthesis code of Bruzual & Chariot (2003, hereafter BC03). The émission fines are then measured by fitting Gaussians to the residual spectrum. We use the same base of 150 stellar populations with âges 1 Myr < /* < 18Gyr, and metallicities 0.005 < Z./Zq < 2.5 as Cid Femandes et al. (2007). Illustrative fits are presented in Asari et al. (2007). Two versions of the ionizing spectrum are considered for each bin:

One including ail the stellar populations inferred from the synthesis (case F, for 'full') and another with the contribution of populations 1 Retired is to be opposed to active, with référencé to star formation. We avoid using the term passive since this might suggest 'without émission fines' (Miller et al. 2003). Ironically, as shown in this Letter, a fraction of so-called active galaxies, by référencé to nuclear activity (e.g. Kauffmann & Heckman 2005), could be genuine retired galaxies. the upper envelope of pure SF galaxies according to Stasinska et al. (2006).

younger than 107 5 yr set to zéro (case O, for 'old'). The reason is that, due to errors in the data and uncertainties inhérent to population synthesis, spectral fits of galaxies without current star formation may attribute a small amount of light to young populations.

For instance, along the LINER branch for jr > 7, we obtain light fractions xy of ~1 per cent at 4020 Â for populations younger than We now ask: what is the impact of the ionizing radiation produced by the stars responsible for the optical continuum upon the émission fines in galaxies?

The number of ionizing photons

The first question to examine is whether the ionizing photons from the evolved stellar populations are enough to account for the observed Ha luminosities in the seagull's wings. For each galaxy, we compute <2hi» the number of stellar photons with energies above 13.6 eV arising from the populations uncovered with starlight (we use the BC03 models also in the Lyman continuum). We then estimate L(Ha)exp, the Ha luminosity expected if ail these ionizing photons are absorbed by the gas présent in the galaxies, and com pare it to L(Ha)obs> the observed value corrected for extinction using We see that, in the LINER branch at large radii, the old stellar populations contribute to the ionizing radiation at least as much as the young ones. Given that xy is very small and thus very uncertain (see Section 3.1), old populations could even be dominant. For the Seyfert branch, on the other hand, it is the young stars which provide most of the stellar ionizing photons. For the SF branches, the old populations play no rôle at ail. The dispersion in L(Ha)exp/ L{Ha)obs as function of ir can be judged by comparing the curves corresponding to the quartiles (left-and right-hand panels of Fig. 2)

with those corresponding to the médian (central panel).

Overall, as seen in Fig. 2, while the stellar populations cor responding to case F can largely explain the upper SF branch in terms of total number of ionizing photons, they explain only ~25 per cent of the Seyfert and LINER branches. The central panel of Fig. 2 shows a déficit by a factor of 1.5 to 4 for médian values of L(Ha;)eXp/L(Ha)obs. In addition, as mentioned in Section 3.1, the young stellar populations uncovered by starlight for the LINER branch are not reliable and case O models could be more appropriate, further increasing the discrepancy. However, our plots are plagued with many uncertainties, as will be discussed in Section 5.

The hardness of the ionizing radiation field

The hardness of the ionizing radiation field can be judged by com paring the values of (2uei and (2Hen (the number of photons above 24.6 and 54.4 eV, respectively), to <2hi-Fig. 3 shows Quei/Q^i In photoionized nebulae, the emission-line ratios are basically determined by three parameters: the hardness of the ionizing radiation, the nebular metallicity and the ionization parameter U (defined as (2iii/(47t/?2ftc) where R is nebular radius, n is the gas density and c is the speed of light). We présent photoionization models using the ionizing radiation from the stellar populations corresponding to provided that log U is between -3 and -4. Models with even higher Z can be found in this same région. With the softer radiation field corresponding to smaller values of ir, one produces models that can cover the inner part of the LINER branch. We conclude that the radiation from old stellar population in metal-rich galaxies can easily account for the emission-line ratios observed in the LINER branch. The other diagrams shown in Fig. 5 are also in reasonable agreement with this proposition.

We now tum to the Seyfert branch. Here, the ionizing radia tion field provided by the evolved stellar populations is significantly softer than in the LINER branch, as seen from Figs 3 and4. Therefore, one can infer from Fig. 5 that the emission-line ratios for the upper Seyfert branch cannot be explained by stellar radiation al one.

UNCERTAINTIES

As has just been shown, our models indicate that the old stellar populations responsible for the optical continua of the considered galaxies can explain the observed emission-line ratios of the LINER branch. However, the observed Hcr luminosities are more difficult to explain. Our case O models reproduce L(Hû')obs within a factor of 2 for about 25 per cent of the LINER galaxies with ir > 1? Does this mean that about one quarter of the LINER galaxies are ionized by hot post-AGB and white dwarf stars, the rest being powered by some other sources? In view of the many uncertainties involved in this study, such a conclusion would be prématuré.

First, there are already uncertainties in the mere process of pop ulation synthesis fitting of galaxy spectra (see Cid Femandes et al. 2005). Secondly, the applied extinction correction may not be appropriate.

Perhaps the most important point is that the modelling of the Lyman continuum in old stellar populations is very uncertain. Unfortunately, ready-to-use evolutionary population synthesis codes that compute the ionizing radiation from hot post-AGB and white dwarf stars are scarce. We inade experiments with the BC03 code and code pegase (Fioc & Rocca-Volmerange 1997). With BC03, for an instantaneous starburst, we find that <2iii dépends on the metallicity: it increases by about 0.3 dex Zq to 0.02 Zq at âges larger than 10* yr, and by over one order of magnitude at around 105 * * 8 yr.

Therefore, any mistakenly assigned metallicity during the contin-2 For /r < 7, the contribution of young stellar populations found by starlight is likely real, and L(Ha)obs is easily reproduced by case F.

uum fitting process will induce some error on Qu,. With pegase, the values of Qh, for solar metallicity are larger by 0.2-1 dex than with BC03. What is the reason for such a différence? Different physical ingrédients (stellar evolutionary tracks and atmosphères)?

Or inaccuracies in the numerical treatments? It seems to us that the major effect cornes from the initial-final mass relation of white dwarfs. The analytic form of this relation has been revised recently (Catalan et al. 2008), which will certainly lead to the changes in the Qu\ prédictions. However, the observational dispersion is very large (see their fig. 2), so that the prédictions will remain uncer tain until one understands better the drivers of the initial-final mass relation(s). et al. (2006), we find that photoionization models built with realistic stellar populations are able to cover nearly the entire BPT plane.

We now discuss why some régions are avoided by real galaxies.

SF galaxies form a narrow wing because, as shown by Dopita et al. (2006), in régions of star formation the metallicity Controls both the ionizing radiation field and the ionization parameter.

But why is the right wing also narrow? Our models for retired galaxies cover the entire space between the two wings of the seagull as well as below the wings. One factor which plays a rôle in shaping the seagull is that, in order to appear in the BPT diagram, galaxies must be sufficiently bright in the r band to satisfy the S DSS sélection criteria for spectroscopy, and must hâve the four involved émission

fines measured with sufficient S/N. As a matter of fact, the lower border of the right wing is traced by the galaxies with the lowest [O ni] équivalent widths in our sarnple.

Metal-poor retired galaxies (models A and B in Fig. 5) should be found below the SF wing and between the wings, if they exist and can be detected. As shown in Cid Femandes et al. (2007), low-metallicity SF galaxies are forming stars efficiently (the 'downsizing' phenomenon), so that our retired low-Z galaxies (models A and B in Fig. 5) may just be 'ahead of their time'. Even if such galaxies existed, they would be faint both in their émission fines and stellar continuum. To illustrate this, we predict the future évo lution of typical low-metallicity SF galaxies under the assumption that they stop forming stars now. Applying the models of BC03

to their present-day stellar populations, we find that 0.1 Gyr from now their radiative output will hâve dropped by a factor of ~300-1000 at hv > 13.6 eV, and in 1 Gyr their r-band flux will hâve faded by ~0.4-0.8 mag (Fig. 6). From the observed mx distribution of the metal-poor SF galaxies in the SDSS, we estimate that over 90 per cent of them should fade beyond the mx < 17.77 limit by the time they retire. Hence, metal-poor retired galaxies either do not 

  de la formation stellaire et l'évolution chimique. Tout cela est fait dans une manière automatisée, appropriée à analyser des échantillons énormes comme le SDSS. De plus, à partir du spectre résiduel purement nébulaire (en soustrayant le spectre observé du spectre modélisé), on mesure les raies d'émission. Ceci permet d'obtenir les propriétés du gaz dans les galaxies, comme la composition chimique et le rougissement. Les raies d'émission fournissent également d'importantes informations sur les sources d'ionisation dans une galaxie, à partir de diagrammes basés sur des rapports de raies. Par exemple, dans le diagramme [N n]/Ha vs. [O m]/H/3, les galaxies sont distribuées en deux "ailes de mouette" sur le diagramme : celles dont le gaz est ionisé par une source de photons durs, par exemple les noyaux actifs (AGNs), résident dans l'aile droite, celles dont le gaz est ionisé par des étoiles jeunes se trouvent dans l'aile gauche.

  différents. C'était la première fois que l'évolution de la relation M*-Z* a été calculée pour les mêmes galaxies. Nous trouvons que la relation M*-Z* est plus inclinée et s'étend à niasses et métallicités plus faibles lorsque l'on remonte dans le temps. Nous avons présenté un modèle simple d'évolution chimique en boîte fermée et constaté que la métallicité stellaire observée a une évolution compatible avec ce modèle. Cela signifie que la relation M*-Z* pour les galaxies de masse supérieure à r>1010 Me peut être expliquée, en première approximation, uniquement par des différences d'histoires de formation stellaire, sans qu'il soit nécessaire d'invoquer des vents. Le Chapitre 6 se concentre sur les LINERS, galaxies classées comme hôtes de AGN, mais qui pourraient en fait être ionisées par une population stellaire vieillissante. Pour certaines galaxies classées comme LINERs dans le SDSS, nous calculons que la luminosité observée en Ha est compatible avec le nombre de photons ionisants émis par les populations stellaires d'âge supérieur à 107 5 ans dans ces galaxies. En outre, les modèles de photoionisation utilisant le champ de rayonnement obtenu pour des populations stellaires vieillissantes couvrent presque complètement les diagrammes de diagnostic tels que [N Ii]/Ha vs. [O Ill]/H/?, ce qui signifie que les rapports des raies d'émission observés dans ces galaxies classées comme LINERs sont également compatibles avec l'ionisation par des populations stellaires vieillissantes. La forme de "mouette" du diagramme [N Ii]/Ha vs. [O m]/H/3 est le résultat d'un effet de sélection. Ce résultat implique une profonde révision du taux d'activité nucléaire dans les galaxies de l'Univers local.

  Zneb também foi feita com mais cuidado do ({ne em C07. Zneb é obtida pelo indice O3N2 (Equaçâo 3.1) e, portanto, dépende das medidas desavermelhadas das linhas [O m] A5007 e [N ii]À6584. A extinçâo é calculada a partir da razâo Iia//H/3)0bs e (Ha/H/3)int sâo as razôes observada e intrinseca, respectivamente. O valor de (Ha/H/3)int varia entre 2.80 e 2.99 para Zneb -0.1-2.5 Z© de acordo comos modelos de fotoionizaçâo de S06. Em resumo, Zneb dépende de Ay*h e vice-versa. Para achar entâo a melhor medida para Zneb e Ay b, escrevi uni programa que partia da hipôtese inicial (Ha/H/?)int = 2.86 e fazia quantas iteraçôes fossern necessârias até que os valores da metalicidade e da extinçâo ficassem estâveis. Em gérai, 2 ou 3 iteraçôes eram suficientes. Além disso, fizeinos testes com outras calibraçôes para Zneb. O indicador [Ar m] 7135/ [O m] A5007, por exemplo, nâo sofre dos mesrnos problemas de calibraçâo para baixas metalicidades que o 03N2. Infelizmente, a linha de [Ar m] é medida para apenas poucos objetos do SDSS, e a nossa anâlise iria sofrer com o numéro muito menor de objetos em cada faixa de metalicidade. Também testamos a calibraçâo de Pilyugin & Thuan (2005), baseada em [O m]A5007/H/? e [O n]3727/H/? e os métodos diretos a partir da temperatura eletrônica de Izotov et al. (2006), derivados da razâo de linhas [OIII]4363/[OIII]5007. Ambos os métodos adicionam unia fraçâo nnnima de objetos à amostra. de galaxias SF em seis faixas A F de Zneb, tais como definidas na Seçâo 3.2. A tabela 1 de A07 caracteriza cada faixa em funçâo da distribuiçâo da metalicidade nebular, massa estelar, idade média, luminosidade em Ha e razâo entre taxa de formaçâo estelar recente e a passada.A figura 5 de A07 traz a correlaçâo de Zneb com vârios parâmetros observacionais e derivados da smtese. Os painéis a e b revelam as relaçôes luminosidademetalicidade e massa-metalicidade, respectivamente. O painel c mostra que Zneb também tem uma forte correlaçâo coin a densidade superficial de massa, definida como E* = M*/27rr50 z, onde r\Qz é o raio petrosiano que engloba metade da luz na banda z. Essa relaçâo nâo é surpreendente dada a correlaçâo forte entre E* e a massa estelar.

. 4 )

 4 onde Nh é o numéro de fêtons com energia hvhi maior do que 13.6 eV produzidos por uma SSP de uma massa solar ao longo de sua vida. Nh dépende tanto da funçâo inicial de massa quanto da metalicidade da populaçâo estelar. Para uma SSP de BC03 de metalicidade solar e IMF deChabrier (2003) entre 0.1 e 100 M0, Nh = 4.47 x 1060 fêtons M©-1.4.6.Efeitos de seleçâo e incertezas do modeloAssumindo o caso de recombinaçâo B e que todos os fotons sâo absorvidos, 1 a cada 2.226 fotons ionizantes résulta ern uni foton em Ha (Osterbrock &taxa de formaçâo atual a partir dos resultados da sintese, temos que integrar a equaçâo 3.2 para obter a SFR. nos ültimos r anos: achar r calibrando a SFRÏVI1 a partir da SFRHa-Pelo argumento acima, esse numéro deve ser consistente coin 107 anos e uma incerteza de d=0.5 dex, que é a resoluçâo temporal mâxima que deve ser possivel atingir coin a nossa sintese espectral para o SDSS. Para isso, o melhor método é correlacionar a SFR por unidade de area (Esfr), evitando que a distancia ou o tamanho da galâxia afetem a calibraçào(Kennicutt 1990). Na figura 6 de A07, mostramos essa calibraçào. O painel a mostra o coeficiente de correlaçâo de Spearman em funçâo de r para Esfr (linha cheia) e para SFR, (pontilhada). A flécha indica o ponto que a correlaçâo é mais forte em Esfr: t -24.5 Myr. Os painéis bec mostram as correlaçôes entre SFRsyn e SFR,Ha, e entre Esfr (Ha) e EsFR(syn) usando o valor para r defini do pela calibraçào. Podemos entâo définir a razâo entre a SFR atual e passada como: SFR*(t = 24.5Myr) b = SFR7(t = tg) ' (48) onde tg é a idade das estrelas mais velhas em uma galâxia. Esse parâmetro é conhecido como parâmetro b de Scalo (Scalo 1986). Na Seçâo 4.3 havia sido mostrado que o parâmetro b tem uma forte correlaçâo com a metalicidade Zneb • 4.6 Efeitos de seleçâo e incertezas do modelo Repetimos a nossa anâlise para varias seleçôes de amostra. Verificamos que nenhuma seleçâo afet.a as tendências que encontramos nas SFHs; no mâximo a curva de uma faixa em metalicidade assume o lugar da curva de uma faixa adjacente, mas todas as tendências se mantêm. A figura 11 de A07 traz os resultados (lesses testes. Os painéis a e b revelam que a amostra de maior qualidade SFhq (S/N > 6 nas quatro linhas do BPT e S/N > 20 no continuo) nâo difere muito da amostra usada ao longo do artigo. Os

Figura 4 . 2 :

 42 Figura 4.2: Historias de formaçâo estelar em diferentes representaçôes para a amostra SFhq. Cores diferentes representam as faixas em metalicidade nebular A (preto) a F(vermelho). Comparamos os resultados da sfntese ajustados com apenas uma extinçâo estelar (linhas pontilliadas) e com extinçâo diferencial (linhas cheias). Essa figura é quase idêntica à figura 12 de A07, mas todas as faixas estâo desenhadas para ressaltar que o efeito desses novos modelos é, no mâximo, aproximar uma curva da curva de uma faixa adjacente.

  Figura 5.1: A relaçâo massa-metalicidade para diferentes redshifts. As linlias clieias marcam a mediana, e as tracejadas, 25 e 75% da distribuiçâo. Essa figura expande a informaçâo condensada no painel à direita da figura 2 de V09.

(

  2002) eTremonti et al. (2004), por exemplo, encontraram que a perda de gâs por ventos détermina a forma da relaçâo M-Z. Em galâxias pequenas (de magnitude na banda B > -18 para Garnett e massa < 1010 M© paraTremonti et al.), o potencial gravitacional nâo é suficiente para evitar a perda do gâs enriquecido por ventos ou explosôes de supernova.Finlator & Davé (2008), por outro lado, afirmam que simulaçôes cosmolôgicas hidrodinâmicas favorecem o cenârio em que as galâxias sâo5.3. 

Figura 7. 3 : 4 :

 34 Figura 7.3: Resultados preliminares da calibraçâo de abundâiicias a partir de linhas de emissâo fortes. Nesses testes forain usados 115 objetos de baixa metalicidade do Data Release 7 do sdss (Abazajian et al. 2009). Nos eixos das abscissas estâo as inedidas feitas por Izotov (em preparaçâo) pelo método direto a partir da temperatura eletrônica de derivada de [O m]À4363/[0 m]A5007. Nos eixos das ordenadas estâo as medidas recuperadas pelo nosso método de calibraçâo. A linlia em cinza marca a identidade. O nosso rnétodo ainda nâo récupéra com boa confiança a abundância de oxigênio ou as razôes de linhas [O m]À4363/[0 m]À5007 e [Ar m]A7135/H/3.

Figura 7 . 5 : 6 :

 756 Figura 7.5: Variaçâo da extinçâo estelar para galâxias SF no diagrama de massa estelar contra metalicidade nebular (calibrada de acordo com a Equaçâo 3.1). Cada caixa mostra o histograma de Ay, cuja escala é dada pelo pequeno plot em vermelho no canto inferior direito; o numéro dentro das caixas indica a mediana da distribuiçâo.Sào mostradas apenas as caixas coin mais de 200 galâxias. A extinçâo estelar tende a aumentar ao longo da rclaçào M*-Zneb-

  Figura 7.9: Como a Figura 7.5, mas para a velocidade de escape, calculada a partir do raio da galâxia e da massa estelar. A velocidade de escape tende a aumentar ao longo da relaçâo M*-Zneb-

  Figura 7.11: Como a Figura 7.5, mas para a historia de formaçâo est.elar dada por SSFR(t). Os numéros em cada caixa indicam a mediana da densidade de massa (em log M© kpc~2). A formaçâo estelar recente diminui ao longo da relaçâo M*-Zneb-

  between heavy-element nebular abundance and galaxy mass, or its extension, the luminosity-metallicity rela tion (e.g.Skillman, Kennicutt & Hodge 1989;Zaritsky, Kennicutt & Huchra 1994), are being extensively used to probe the métal en richment along cosmic history. Clear signs of évolution are being revealed by studies of these relations at both intermediate(Savaglio et al. 2005;Lamareille et al. 2006; Mouhcine et al. 2006) and high z(Shapley et al. 2005;Maier et al. 2006;Erb et al. 2006), which generally find significant offsets in these relations with respect to what is observed in the local Universe.Altematively, one may study galaxy évolution through fossil methods, which reconstruct the star formation history (SFH) of a *E-mail: cid@astro.ufsc.br galaxy from the properties of its stars. The most detailed studies of this sort are those based on nearby galaxies (including our own), where colour-magnitude diagrams and even spectroscopy of individual stars are feasible (e.g.Smecker-Hane et al. 1996;Hernandez, Gilmore & Valls-Gabaud 2000;Dolphin 2002;Rizzi et al. 2003; Skillman et al. 2003;Monaco et al. 2005;Koch et al. 2006). Comparing such observations with prédictions of evolutionary synthesis models then allows one to infer how stars were bom as a function of time, as well as to constrain the Chemical enrichment historiés of the galaxies (e.g.Lanfranchi, Matteucci & Cescutti 2006;Carigi, Colin & Peimbert 2006). By mapping how stars become richer in metals as a galaxy âges, these studies provide useful constraints for Chemical évolution models.

Figure 1 .

 1 Figure 1. (a) 175 234 SDSS galaxies in the B PT diagram. The solid line (Stasirtska et al. 2006) divides star-forming galaxies from those hosting AGN. The dashed line shows the dividing line used byKauffmann et al. (2003). (b) The SF sample, chopped into six bins of nebular abundance. The number of galaxies in each bin is given in the right, along with the corresponding mean Zneb values (in solar units). Galaxies close to bin borders are not plotted for clarity.

  completed detailed pixel-by-pixel spectral fits of 573 141 galaxies from the Main Galaxy Sample of the SDSS Data Release 5 (Adelman-McCarthy et al. 2007) with the synthesis code starlight described in SEAGal I and II. starlight décomposés an observed spectrum in terms of a sum of simple stellar popula tions covering a grid of 25 âges from r* y = 1 Myr to 18 Gyr and six metallicities Z* y = 0.0001 to Ô.05 (i.e. from 0.005 to 2.5 Zq). Each of these A * = 150 instantaneous bursts is represented by a spectrum extracted from the evolutionary synthesis models of Bruzual & Chariot (2003) for a Chabrier (2003) initial mass function, Padova 1994 tracks and the STELIB library (Le Borgne et al. 2003). Illustrative fits are presented in SEAGal I and A07. Emission lines were measured from the starlight subtracted spec tra. Fig. l(a) shows the Baldwin, Phillips & Terlevich (1981, here after BPT) diagnostic diagram for 175 234 galaxies for which H/3, [Oiii]X5007, Hcr and [NnJA.6584 are ail detected with signal-tonoise-ratio S/N ^3. Reddening-corrected line fluxes are used in this plot and throughout our analysis. The distribution of points in this diagram resembles a flying seagull, with SF galaxies occupying the left wing and galaxies hosting an active galactic nu cléus (AGN) on the right. We define SF galaxies as those under the divisory line proposed in SEAGal III (solid line in Fig. la). Further requiring a minimum continuum S/N of 10 at ~4750 Â leaves a set of 84 828 objects, hereafter the SF-sample, shown in Fig. l(b). Results for alternative sample définitions (involving different S/N thresholds, other AGN/SF séparation criteria and cuts on galaxy inclination) are presented in A07. Note that by selecting SF galaxies we indirectly select spirals over elliptical galaxies. The SF-wing is essentially a sequence in nebular metallicity (SEAGal III and référencés therein), which we quantify by the oxygen abundance obtained through the strong-line calibration of Stasinkska et al. H)q = 4.9 x I0"4 (Allende Prieto, Lambert & Asplund 2001). We believe this is the most reliable calibration available in the literature, as it is based on temperature-based abundances in the metal-poor case and on the ([On] + [Oni])/H/3 calibration of Pilyugin (2000) in the metal-rich régime, using a large sample of giant H n régions. L18 R. Cid Fernandes et al. Z neb is both a physically motivated and mathematically convenient coordinate to map galaxy positions along the left wing in the B PT diagram. From the tip of the SF-wing in Fig. l(a) to its bottom, our SF sample spans the Zneb ~0.2-1.6 Zq range. In Fig. l(b) this

  i=i where pj is the fraction of the total stellar mass (A7*) which is in the ;'th population. Both (log/*) and (Z*) proved to be mathematically robust and astrophysically valuable summaries of the full population vector. In SEAGal I, II and A07 these averages are shown to cleanly reveal important relations, such as those between (log/*) and M* (associated with the galaxy downsizing phenomenon), (Z*) and M* (the mass-metallicity relation), and (Z*) and Zneb (which reveals a connection between stellar and nebular Chemical enrichment levels).

  stands for 'converted into stars', and is introduced to distinguish pcj and /Vf* from pj and M*, which refer to the mass still locked inside stars, i.e. discounting the mass returned to the interstellar medium (ISM) by stellar évolution. rç*(^*) is thus a monotonie function which grows from 0 to 1 as /* runs from the time that the galaxy formation starts to the présent, and maps what fraction of M* was assembled up to a given lookback time.1

Fig. 2 (Figure 2 .

 22 Fig. 2(a) shows the mean mass-assembly functions obtained for each of the six Zneb bins illustrated in Fig. l(b). The systematics is évident to the eye. Galaxies with low Zneb are slower in assembling their stars than those with high Zneb-This behaviour reflects the fact

  find b decreases by a factor of ~20 from the lower to the higher Zneb bins. (SFRs are discussed in A07, where it is also shown that our synthesis based SFRs agréés very well with those derived from the Ha luminosity.) Alternatively, one may quantify this behaviour by computing the time at which the mean rç*(r*) function reaches 0.75, i.e. the time at which the mass involved in star formation reaches 3/4 of the total. We find that 175 increases from ~0.8 Gyr at Zneb ~0.31 (bin A in Fig. lb), to 2 Gyr at Zneb ~0.77 Zq (bin D), and 8 Gyr for the highest Zneb bin. Given that Zneb correlates strongly with

Fig. 2 (

 2 Fig. 2(b). Again, the main message is clear: galaxies with a metalpoor ISM are slow in their stellar Chemical enrichment. In the lower Zneb bins, Z,(/*) only reaches the current stellar metallicity level (~1/3Zq) as recently as f* ~100 Myr. Galaxies with larger Zneb, on the other hand, hâve flatter and systematically higher Z.(r") curves. These massive Systems hâve essentially completed their mass assembly and Chemical évolution long ago, so that recent starforming activity has a negligible impact upon Z,(r*).The fact that bins of increasing Zneb hâve increasingly higher présent Z, values confirms that nebular and stellar metallicities are related (SEAGal I;Gallazzi et al. 2005; A07). We nevertheless warn the reader not to overinterpret the quantitative relation between these two quantities, given that they are derived by completely different and hardly comparable means. In this Letter we use [Om]/[N n] as an indicator of Zneb (équation 1 ), because of its évident relation to the position of the points in the BPT diagram. However, as emphasized byStasinkska et al. (2006), strong line methods for deriving nebular abundances hâve been calibrated using giant Hn régions and their relevance for integrated spectra of galaxies has still to be examined, because of the rôle of metallicity gradients and of the diffuse ionized medium in galaxies. In addition, there is so far no study comparing nebular abundance calibrations with stellar abundances. Therefore, our Zneb values are not necessarily on the same scale as stellar abundances. Insofar as this Letter is concerned, this is not a problem because, irrespective of absolute values, it is the existence of a Zneb -Z, relation which is responsible for the ordering of the different Chemical évolution curves in Fig.2(b).

Figure 3 .

 3 Figure 3. As Fig.2, but binning SF galaxies by their stellar mass, using five 1-dex-wide bins, centred at (from bottom to top) log M*/Mq -7.5, 8.5, 9.5, 10.5 and 11.5, which contain 583, 4750, 29056. 46 847 and 3550 galaxies, respectively.

Fig. 3 :

 3 Fig.3: massive galaxies today formed most of their stars and metals a long time ago, whereas Iow-mass galaxies are still forming stars and metals actively, in line with recent theoretical results(De Rossi, Tissera & Scannapieco 2006).The signal of Chemical évolution is présent in ail curves in Figs 2(b) and 3(b). Curiously, however, this évolution does not seem to start from very low Z* for massive, high-Zneb galaxies (e.g. bins E and F). We interpret this apparent lack of old, metal-poor popula tions as due to a combination of lack of âge resolution at high r* and an intrinsically fast évolution ocurring at early times. In low-mass Systems (e.g. bin A), Z,(r") starts its A/* ~2 Gyr long rise period at r* ~2-3 Gyr, such that A/*/r* is of the order of unity. The same process happening at r* > 10 Gyr would yield Af*/f* < 0.2, and consequently be much harder to résolve. As shown above, by /* = 10 Gyr massive galaxies had already assembled most of their mass, so their Chemical évolution must also hâve occurred early on. This mixture of populations of widely different metallicities but similar âges results in Z* values at r* > 10 Gyr which are dominated by the metal-rich stars already présent at those times. Such an early and rapid formation period should lead to enhancement of éléments synthesized in Type II supernovae(Worthey, Faber & Gonzalez 1992; Thomas et al. 2005). Though we cannot map elemental abundances with our method, we do detect a signal of increasing a enhancement towards the bottom of the BPT diagram by means of systematically increasing residuals in spectral bands due to a éléments such as Mg (A07), indirectly supporting the scénario outlined above.

  in the famous [Om]5007/H/3 versus [N n]6583/Ha diagnostic diagram (hereafter, the BPT diagram) came as a surprise.

5 (

 5 photoionization models to refine the criteria to distinguish between normal SF galaxies and AGN hosts.Mateus et al. (2007, hereafter SEAGal IV) deals with environment effects, studied in terms of the relations between mean âge, current SFR, density, luminosity and mass.Only in CidFernandes et al. (2007, hereafter SEAGal V), we turned our attention to the detailed time-dependent information provided by the synthesis. We hâve used the entire Data Release 5(DR5 Adelman-McCarthy et al. 2007) to extract the population of SF galaxies, and study their Chemical enrichment and mass assem-

2. 1

 1 STARLIGHT fits After correcting for Galactic extinction (with the maps of Schlegel, Finkbeiner & Davis 1998 and the reddening law of Cardelli, Clayton & Mathis 1989, using Rv =3.1), the spectra were shifted to the rest frame, resampled to AX = 1 Â between 3400 and 8900 Â, and processed through the starlight spectral synthesis code described in SEAGal I and II. starlight décomposés an observed spectrum in terms of a sum of simple stellar populations (SSPs), each of which contributes a fraction xj to the flux at a chosen normalization wavelength (À0 = 4020 Â). As in SEAGal II, we use a base of A* = 150 SSPs extracted from the models of Bruzual & Chariot (2003, BC03), computed for a Chabrier (2003) initial mass function (IMF), 'Padova

Fig. 1 ,Figure 1 .

 11 Fig.1, becomes évident when averaging many residual spectra (seeSEAGal V and Panter et al. 2007), and tends to be more pronounced for older objects. The comparison between STELIB stars and theoretical models presented byMartins et al. (2005) gives a due to the origin of this problem. A close inspection of fig.21in their pa per shows that the STELIB spectrum has an excess of flux on both sides of Hfi when compared to the model spectrum. This suggests that the 'Hfi trough' is related to calibration issues in the STELIB library in this spectral range. This was confirmed by starlight experiments which showed that the problem disappears using the SSP

23

  Définition of the star-forming sample Since the pioneering work of Baldwin et al. (1981), emission-line objects are classified in terms of their location in diagrams involving pairs of line ratios. As explained in SEAGal III, the [N n]A6584/Hcr versus [O in |À5007/H/I diagram (the BPT diagram) is the most useful for this purpose, mainly due to the partially secondary nature of N (i.e. the increase of N/O as O/H increases, e.g.

Figure 2 .

 2 Figure 2. The SF sample in the BPT plane, chopped into bins of nebular abundance Zneb = (O/H)/(O/H)© • Ail lines hâve been corrected by reddening (see Section 2.5). These same bins are used throughout this paper. The number of galaxies in each bin is given on the left-hand side. On the righthand side, we show the corresponding mean Zneb and log mean M* values (in solar units). Galaxies close to bin borders are not plotted for clarity.

Figure 3 .

 3 Figure 3. Normalized histograms of observed and physical properties for the SF (solid) and SFht| (dotted) samples.

  différences increase with the amount of dust, as measured by the derived A* values or by Ha/H/L Yet, Kolmogorov-Smimov (KS) tests showed that in no case the distributions of x2 values differ significantly. This implies that the choice of reddening law cannot be made on the basis of fit quality. A wider spectral coverage would be needed for a definitive empirical test.When using different recipes for the dust atténuation, the synthesis algorithm has to make up for the small variations from one curve to another by changing the population vector and the value of A*. To quantify these changes we compare results obtained with the Calzetti and CCM curves, and consider only the most extincted objects. Compared to the results for a CCM law, with the Calzetti law the mean stellar âge decreases by 0.09 dex on the médian (qualitatively in agreement with the results reported in fig.6of Panter et al. 2007), the mean stellar metallicity increases by 0.05 dex, A* increases by 0.07 mag and stellar masses decrease by 0.02 dex.

Figure 4 .

 4 Figure 4. (a) Relation between stellar and nebular extinctions for the SF sample. Lines indicate the 5,50 and 95 percentiles, (b) Equivalent width of the ISM component of the Na D doublet (as measured from residual spectra) against the stellar extinction.

Fig. 4{b )

 4{b Fig. 4{b) shows the relation between AWNraD and A*. The plot shows that these two independently derived quantifies correlate

Fig. 5 (

 5 Fig. 5(b) shows our version of the A/*-Zneb relation. Because of the expected bias in our Zneb estimate at the lowest metallicities (see Section 2.4), we expect the real mass-metallicity relation to be flatter at low M* than seen in this plot. As shown by Kauffmann et al. (2003b), stellar mass and stel lar surface mass density (S*) are very strongly related. It is thus no surprise to find that Z* also correlates with Zneb, as shown in Fig. 5(c). Our définition of Z* is the same as adopted by Kauff mann et al. (2003b), namely Z* = M*/2rrr?0^, where r^is the half light Petrosian radius in the z band.

Fig. 5 (

 5 Fig. 5(d) shows how nebular extinction increases systematically with Zneb One factor which surely contributes to this relation is the rise in dust grain formation with increasing gas metallicity, but other factors may corne into play as well(Stasinska et al. 2004).

Fig. 5 (Figure 5 .

 55 Fig. 5(e) shows how nebular and stellar metallicities correlate. This important relation, first presented in SEAGal I (for differ ent sample, SSP base and Zncb scale), shows that stellar and ISM Chemical enrichment levels scale with each other, as one would expect on the basis of simple Chemical évolution scénarios. The large scatter in Fig. 5(e) is mostly intrinsic (as we verified comparing the relation obtained for data of different qualifies), in qual itative agreement with the idea that stellar and nebular metallici ties reflect different evolutionary phases and react differently to the several processes which regulate the Chemical évolution of galax ies. A similar relation was obtained by Gallazzi et al. (2005) using

4. 1

 1 Compression methodsAs reviewed in Section 2.1, starlight décomposés an observed spectrum in terms of a sum of SSPs, estimating the Xj(j = 1, ..., A*) fractional contribution of each population to the flux at À0 = 4020 Â. For this work we used a base of (V* = 150 SSPs from BC03, spanning 25 âges between t+j -1 Myr and 18 Gyr, and six metallicities (0.005 < Z*,; ^2.5Zq). Exemplary fits were shown in Fig.1.

=

  one associâtes a duration to each burst. The SFR (/*) function is constructed by sampling the smoothed mass fraction vector fxcs in a quasi-continuous grid from logf* = 5.6 to 10.5 in steps of Alogf*

  into account. From the BC03 models for a Chabrier (2003) IMF, a 1010yr-old population has ~2 M+, i.e. only half of its initial mass remains inside stars nowadays.Three clarifying remarks are in order. (1) Ail équations above are marginalized over Z*, i.e. SFR(/,) = SFR(r*, Z,) mea-

  to a cosmologically consistent time-grid if needed. (3) Finally, since our main goal is to compare the intrinsic évolution of galaxies in different parts of the SF wing in the BPT diagram, throughout this paper we will consider âges and look-back times in the context of stellar évolution alone. In other words we will not translate r* to a cosmological look-back time frame, which would require adjusting the r* scale by adding the z-dependent look-back time of each galaxy. 43 Mass assembly historiés Another way to look at the population vector is to compute the total mass converted into stars as a function of time: rilit.) = Y P* (7) which is a cumulative function that grows from 0 to 1, starting at the largest /*, tracking what fraction of M* was converted to stars up to a given look-back time.

Fig. 6 Figure 6 .

 66 Fig. 6 shows the results of this exercise. Panel (a) shows the run of the Spearman coefficient (Rs) for different values of r, with the

Section 6 .

 6 Figure 7. Left-hand part: Distributions of SFHs for Zneb bins A (black) and F (red), as defined in Fig. 2. For each bin, we show the mean SSFR (solid line), the médian (dashed fine) and the 5 and 95 percentiles of the distributions (dotted fine). Righl-hand part: Normalized distributions of log SSFR for r* = 25 Myr. Gaussians are superimposed to illustrate that the distributions are lognormal.

Fig. 8

 8 Fig. 8 shows the derived SFHs for the A-F bins in four different représentations, from top to bottom: ,vs, SFR, SSFR and as a function of stellar âge /*. Each line represents a /*-by-/* average over ail galaxies in the bin. The plots show that young populations are présent in a proportion which increases systematically as Zneb decreases. This is évident, for instance, in the xs(t+) panel, which shows how Zneb-related âge distributions combine to produce the corrélation between (log/*)L and Zneb depicted in Fig. 5(f). The SFR(/*) curves (panel b) show that SF galaxies of different Zncb differ more in their past SFR, low Zneb having SFRs about 100 times lower than those of high Zneb afew Gyr ago. In the more recent past, ail curves converge to SFRs of a few Mq yr~'. At first sight,

Figure 8 .

 8 Figure 8. Mean SFHs for the six different Zneb bins defined in Fig. 2, in four different représentations: (a) smoothed population vector, Â7(/*), (b) SFR(/*), (c) SSFR(/*) and (d) /?£(/*)

  Figure10. As Fig.8, but binning SF galaxies by their stellar surface densities, with five 0.5 dex wide bins centred at log E* = 7.25, 7.75, 8.25, 8.75 and 9.25M© kpc-2, containing 1477, 12554, 37 177, 27742, 3046 galaxies, respectively.

Figure 11 .

 11 Figure 11. Average mass assembly historiés (/)*, top panels) and spécifie SFR (SSFR, bottom) historiés for Z™* bins B to F, colour-coded as in Fig. 2. Solid fines show the curves for different sample définitions. For comparison, the dotted fines in ail panels show the curves for bins B (lower curves; magenta) and F (upper curves; red) of the full SF sample.

  Kauffmann et al. (2003c) line. Figs 11 (c) and (d) show that the SFHs for bins in this sample are nearly indistinguishable from those obtained with the SEAGal classification criterion.

Figs

  Figs 11 (e) and (f) show that the SFHs derived with this sample are practically the same as for the full sample.Aperture effects are a common source of concem in studies of SDSS spectra (e.g.Gômez et al. 2003, SEAGal I). To investigate how such effects impact upon our SFHs we defined two samples: the SFJ sample, which comprises 58 153 SF galaxies with z ^0.05 (as opposed to 0.002 for the full SF sample), and the SPP sample, comprising only the 1096 objects with more than half of their z-band luminosity inside the fibre. Both criteria exclude proportionately more the population of low-Zncb, low-M* galaxies (distant galaxies of this category are not présent in the SDSS, due to limiting mag nitude). Accordingly, the change in SFHs is only noticeable for the lowest Z"eb bins, as shown in Figs 1 l(g)-(j). In particular, Fig.11 (j)shows that the mean SFH for bin C in the SFap sample matches that of bin B in the full sample. Of ail selection-induced changes

  older and thus redder young population would compensate the mismatch in A*. Allowing for A*(r < 10^yr) larger than the A* of the t > 107 yr populations makes it possible to fit the same spectrum with younger and dustier populations. Hence, the recent SFRs should increase.

  âge, which decreases by A(log r*)L ~0.1 dex.The changes in other global properties such as A* and M* are much smaller than this.

Figure 12 .

 12 Figure 12. Comparison between mean SFHs for the SF1,q sample modelled with only one stellar extinction (dotted lines) and with differential extinction (solid lines). For clarity, only Zncb bins B, D and F are drawn, and colourcoding is the same as in Fig. 2. Each panel shows a different représentation of SFHs: (a) smoothed population vector, I7(?*), (b) SFR(r*),(c) SSFR(f*) and (d) rç£(r*)-

  (i) We hâve detected a systematically overestimated continuum level around H/3, whose origin we tentatively attribute to deficiencies in the STELIB calibration in this range. Gaussian fits to the H/3 émission which disregard this offset tend to underestimate the line flux by 4 per cent on average, and ~7 per cent in the case of weaker lines. These are relatively small, yet systematic effects, which propagate to estimâtes of nebular extinction, metallicities and galaxy classification.

  the observed results are in agreement with what is expected for the évolution of galaxies. This brings us to another concern in M-Z studies: the bias in the dérivation of the nebular metallicity from émission lines. The évolution of the M-Z relation is usually studied in terms of the abundance of oxygen in the ISM gas. As it is known (see e.g.Kewley & Ellison 2008), the measurement of nebular abundances is very dépendent on the method and calibrations used.With our stellar population synthesis code starlight, we can dé termine the stellar metallicities, Z*, and the total masses in stars, M,, at different epochs for any given galaxy(Cid Fernandes et al. 2007).We can thus follow the évolution of the M-Z relation for the same set ofgalaxies at different redshifts.Panter et al. (2008) used similar techniques and discussed the nebular metallicities with respect to stellar metallicities. In this study, we rely on stellar metallicities alone because, although they hâve their own problems, they are free from the biases affecting nebular abundance déterminations.

  Since starlight allows the détermination of the star formation historiés (SFHs), we can go one step further and construct a simple Chemical évolution model using these SFHs and see under what conditions it reproduces the observed time évolution of Z*. Savaglio et al. (2005) made a similar study, but using a theoretical SFH with e-folding time, related to galaxy mass, and for nebular metallicity instead of Z,. Here, we use the SFH obtained directly for each galaxy. DATA Our parent sample is composed of ail objects spectroscopically classified as galaxies from Fifth Data Release of the Sloan Digital Sky Survey (SDSS) (Adelman-McCarthy et al. 2007). We exclude duplicate observations, and impose the following sélection criteria: 14.5 < m, < 17.77 (from the Main Galaxy Sample définition), z-band aperture covering factor >20 per cent (to reduce aperture effects and avoid intragalactic sources), S/N > 10 at 4750 Â (to provide reasonable stellar population fits) and a narrow range in redshift |z -0.11 < 0.015 (to be able to transform stellar âges into redshift in a simple way). The central redshift of 0.1 was chosen to maximize the number of objects and thus ensure reliable statistics over the whole range of âges considered in the analysis. We adopt a Ho = 70 km s-1 Mpc"1, QM = 0.3 and QA = 0.7 cosmology where appropriate. There are 82662 objects in the resulting sample.We use the same data processing as outlined inCid Fernandes et al. (2005) andMateus et al. (2006). Our code starlight recovers the stellar population content of a galaxy by fitting a pixel-by-pixel model to the spectral continuum. This model is a linear combination of 150 simple stellar populations (SSP) extracted fromBruzual & Chariot (2003, hereafter BC03) with âges 1 Myr < r* < 18 Gyr, and metallicities 0.005 < Z/Zq < 2.5, as in CidFernandes et al. (2007). Emission fines are then measured in the residual spectra, which reduce the contamination by stellar absorption features.We divide our parent sample into three subsamples: star-forming (SF), retired (R) and Seyfert (S) galaxies. For both the SF and S samples, we select objects with S/N > 3 in the[Oiii]à5007, H/!, [N n]À6584 and Ha émission fines, i.e. ail the fines involved in the Baldwin, Phillips & Terlevich (1981, BPT) diagram. The SF sample is composed by the objects below the fine proposed by Stasiriska et al. (2006) to separate SF galaxies from AGN hosts. The S sample is defined as the galaxies with log([N n]/Ha) > -0.3 and above the fine defined by Schlickmann (2008) in the BPT diagram to separate S galaxies from low-ionization nuclear émission région (LINER)-like galaxies: log([0 m]/H/3) = 0.901og([Nii]/Ha) + 0.48.

  fines and [On]k3727 ('passive'); or S/N < 3 in [O m] and/orH/1, but S/N >3 in [Nu] and Ha, and log([Nn]/Ha) > -0.15 ('weak émis sion fines'); or S/N > 3 in the 4 BPT fines, log([N u]/Ha) > -0.15, below the fine defined by Schlickmann (2008), and Ha émission predicted by old stellar populations equal or greater than the observed value ('retired mimicking LINERs').

Figure 1 .

 1 Figure 1. Historiés of our parent sample in stellar mass bins. Thicker fines are used for more massive bins. Panels show the médian curve for time évolution of stellar mass évolution (left; see équation 1) and metalficity évolution (right).

Fig. 2 Figure 2 .

 22 Fig. 2 shows the M-Z relation. The left-hand panel shows the present-day M-Z relation for ail the galaxies in our sample. The right-hand panel depicts the évolution of the M-Z relation for the galaxies in mass bins A-F (black fines): Z*(r) versus M*(/), where M"(t) is the stellar mass integrated from the epoch of galaxy formation until a lookback time t, assuming that the stars always pertain to the galaxies where they were born. The resulting M-Z relations for different redshifts are plotted in bold fines. One can see that the M-Z relation evolves. As lookback time increases, the M-Z relation steepens and covers a larger range of values. As investigated by Panter et al. (2008), using stellar modelling codes other than that of BC03 changes the values of M"{t) and Z*(r), but the general behaviour seen in the présent Letter as well as in Panter et al. (2008) remains. One should note that

  ) with t?*(f) given by starlight to find what value of /now is needed to reproduce the observed évolution of the mean stellar metallicity Z*(f)-Unfortunately, y is not well known. The measured stellar metallicities are mainly sensitive to the abundance of iron, and the Fe/O ratio is expected to vary with time. For our simple approach, we assume that ail the metals vary in lockstep and simply speak of the 'métal yield'. We détermine an empirical value of y by applying our model to the highest mass bin of the SF sample, using the Schmidt-Kennicutt law (Kennicutt 1998)1 to dérivé the value of/n0w from the extinction-corrected Ha luminosity. This procedure should give a reasonable value of the métal yield, since it is for the highest mass bin that the closed-box model is expected to be the most relevant.We obtain y = 0.03. This value is in satisfactory agreement with the value (0.04) obtained for aChabrier (2003) stellar initial mass function by multiplying the oxygen yield fromWoosley & Weaver (1995) by the Solar metal-to-oxygen mass fraction.

Fig. 3

 3 Fig. 3 shows Z*(r) as a function of t for bins A-F of our SF, R and S galaxy samples. The black lines represent the médian and quartiles of the observed distributions. The red lines are the model results for the values of /now indicated at the top of each panel. The models were adjusted by eye to reproduce well at least the last 1 Gyr of the observed Z*(f). For the SF sample, we also indicate the médian value of/now obtained from the Schmidt-Kennicutt law. This figure

Figure 3 .

 3 Figure 3. Chemical évolution in our samples of SF, R and S galaxies in thc 0.30-dex-wide stellar mass bins. centred in logM,/MQ = 10.0 (A), 10.3 (B), 10.6 (C), 10.9 (D), 11.2 (E) and 11.5 (F). Each panel shows the médian and quardles of the évolution of Z* as found by starlight (black lines) and as obtained with the simple closed-box model (red lines). The value of/n0w needed to reproduce the médian present-day Z. is indicated at the top right of each panel. For the SF sample, we also indicate the médian of/now as measured by the Schmidt-Kennicutt law (/sk)-The yield y was adjusted to rcproduce bin E of the SF

  nomination is given with reference to the typical emission-line ra tios of active galactic nuclei (LINER stands for 'Low-Ionization Nuclear Emission Régions' in the original paper byHeckman 1980). However, it is by no means obvious that ail galaxies of the right wing are powered by an energetically dominant active nucléus, i.e. they are genuine Seyferts or LINERs.Binette et al. *E-mail: grazyna.stasinska@obspm.fr fSemi-Empirical Analysis of Galaxies.

2. 2

 2 Chopping the seagull Given the distribution of the galaxies in the BPT plane ([O m]/H/3 versus [N n]/Ha), it is convenient to use polar coordinates (r, 0), with the centre defined at the point of inflection of the médian curve for the distribution of [O m]/H/3 as a function of [N n]/Ha : (log[N ii]/ Ha = -0.43; log[Om]/H/) = -0.49). The SF wing corresponds to 0 ~117°for the low-metallicity branch, and ~-94°for the highmetallicity branch, while the right wing corresponds to 0 = 45-90°, with 9 ~77°for Seyferts and ~59°for LINERs. Emissionline properties, stellar populations and the corresponding ionizing spectra were obtained after chopping the seagull in 40 angular and 18 radial bins, as shown in Fig. 1, and analysed for different bins in 9, as a function of the radial bin index, it.3 THE STARS IN THE SEAGULL'S WINGS 3.1 Two extreme cases to consider

Figure 1 .

 1 Figure 1. The chopped seagull and the définition of the SF branches (light blue and dark blue), Seyfeit branch (green) and LINER branch (red). The data points represent ail the galaxies defined in Section 2.1 that hâve S/N > 3 in the four fines involved in the plot (131 287 objects). The black curve is

  107 s yr. This corresponds to mass fractions of less than 10~5 and is clearly noise. Yet, given the five orders of magnitude différence between ionizing fluxes of young and old populations, even such optically insignificant fractions can hâve a substantial impact on the shape and intensity of the radiation field in the Lyman continuum of hydrogen. Setting xy = 0 circumvents this problem, while at the same time emphasizing the effects of post-AGB stars alone.Along the Seyfert branch, on the other hand, xy averages to 7 per cent, implying that Seyferts hâve younger stars than LINERs, in full agreement with both SDSS (e.g.Kauffmann et al. 2003) and independent studies (e.g. GonzalezDelgado et al. 2004). Ail this is clearly illustrated in fig.3of CidFemandes et al. (2008), where we show how star formation historiés vary across the BPT diagram.

  Ha/H/f and theCardelli. Clayton & Mathis (1989) extinction law for RV = 3A. In each bin of the chopped seagull, we define the mé dian and quartiles of the L(Ha)exp/L(Ha)0bs distributions. In Fig.2, we plot their values as a function of i,, with style and colour coding as indicated in the caption.

Figure 2 .

 2 Figure 2. Variations of L(Hûî)eXp/L(Ho;)obs along ;r for the upper SF branch (light blue), lower SF branch (dark blue), Seyfert branch (green) and LINER branch (red). Full Unes: case F; dotted lines: case O. The middle panel shows the médian value of L(Ha)eXp/L(Ha)obs, while the left-and right-hand panels show the 25 and 75 quartiles.

Figure 3 .

 3 Figure 3. Variation of Q\\eJQ\u along ix. Conventions as in Fig. 2.

Figure 4 .

 4 Figure 4. Variation of <2hch/<2hi along iT. Conventions as in Fig. 2.

Fig. 4

 4 Fig.4is analogous to Fig.3for ÔHen/ÔHi and shows a very different behaviour of the LINER branch with respect to the other branches, since the effect of hot post-AGB stars and white dwarfs is here dominant. Note that the values of Quell do not take into account the X-ray radiation produced by hot stars as well as X-ray binaries, both of which could hâve a non-negligible contribution to Qhco/Qhi-

Figs 3 and 4

 4 Figs 3 and 4 for case F. Not much is known about the gas distribution in the galaxies of the right wing. For simplicity, we hâve assumed a thin shell geometry, with density n = 500 cm~3. Using the code photo, we hâve computed models with different values of U and of the nebular metallicity Z (defined as the oxygen abundance in units of 4.9 x 10~\ the solar value ffornAllende Prieto, Lambert & Asplund 2001). The abundances of the other éléments follow the same prescriptions as in[START_REF] Stasiriska | [END_REF].We first discuss the LINER branch. In Fig.5, we show the lo cation of several model sequences in classical line ratio diagrams, superimposed on the data points. Each sequence has Z varying from 0.03 to 6Zq, and is defined by U as indicated in the figure caption. The ionizing radiation for ail the models is given by the stel lar population of the galaxy having the médian value of ÔiWÔhi in the bin marked in orange in Fig.1. Note that the results are almost the same when changing the radiation field from it = 7 upwards in the LINER branch. We see in Fig.5that models with metallicities twice solar cover the tip of the LINER branch in the BPT diagram,

Figure 5 .

 5 Figure 5. Our sarnple galaxies in four classical emission-line ratio diagrams: [Om]/H/3 versus: [On]/H/i (a), [Nn]/Ha (b), [Snl/Hcr (c) and [Oi]/H/S (d). Model sequences for the LINER branch are shown for differ ent values of log U: -2.3 (red), -2.7 (green), -3 (blue), -3.3 (cyan), -3.7(purple), -4 (yellow) and -4.4 (orange). The black line is the model sé quence for SF galaxies fromStasinska et al. (2006). The metallicities Z/Zq are marked with letters as follows: 0.2 (A), 0.5 (B), 1 (C), 2 (D) and 5 (E).

6

  WH Y WE SEE A SEAGULL IN THE BPT DIAGRAM By merging the results of Sections 3 and 4 with those of Stasinska

Figure 6 .

 6 Figure 6. Evolution of (a) the r-band flux, (b) <2hi and (c) Qu&IQhi for the galaxy corresponding to the médian of stellar âge in bin it = 8 of the upper

  em conta ventos. No Capitulo 6, chamamos a atençâo para as LINERs, galâxias classificadas como hospedeiras de AGNs, mas que na verdade podein ser ionizadas por estrelas veillas. Para algumas galâxias classificadas como LINERs no sdss, calculamos que a luminosidade de Ha observada é compativel com o numéro de fotons ionizantes emitidos pelas SSPs de idade maior do que 107 5 anos dessas galâxias. Além disso, modelos de fotionizaçào utilizando o campo de radiaçâo obtido para populaçôes estelares ve illas cobrem quase completamente alguns diagramas de diagnostico, como o [N n]/Ha vs. [O iii]/H/3, confirmando que também as razôes de linhas sâo compativeis com o cenârio de ionizaçâo por estrelas veillas. Argumentamos como a forma da gaivota do diagrama [N Ii]/Ha vs. [O iii]/H(3 é resultado de efeitos de seleçâo. Isso implica lima profunda revisâo da taxa de atividade nuclear nas galâxias do Universo local.

		CapituCo
		Introducâo
		/
		As galâxias, como os elefantes, têm boa memôria.1
		Binney & May (1986)
	1.1	Evoluçâo de galâxias
	1.1.1	Um resumo do cenario atual

O Capitulo 7 sintetiza os principais resultados e propôe alguns trabalhos a serein seguidos no futuro. Apresentamos nessa tese novos métodos para estudar a evoluçâo de galâxias a partir da sintese de populaçôes estelares. Nossos métodos abrem portas, dentre outras, a estudos mais a fundo sobre a as relaçôes entre os AGNs e formaçâo estelar.

1 .

 1 Introduçâodas galâxias de diferentes massas é o que se chaîna de downsizing, aludindo ao fato de que lioje as galâxias menores sào mais eficazes em formai' estrelas e de que no passado os papéis estavam invertidos. Hâ também galâxias que quase nâo têm mais gâs dispomvel e pararam de formar estrelas. E, dentre as galâxias que ainda formam estrelas e as que jâ se aposentaram, existem aquelas que possuem uni buraco negro ativo menos conspfcuo do que um quasar. De acordo coin a ionizaçâo do gâs produzida pela atividade central, elas sâo classificadas como Seyferts ou LINERs.

  Para comparaçâo, definimos também uma amostra cujos dados sâo de rnaior qualidade: S/N > 6 nas quatre linlias do BPT e S/N > 20 no contmuo. A amostra de contrôle SFhq tem uni total de 17142 objetos.

	Capitulo Jf particularidades dos dados desse estudo sâo descrit.as a seguir.
		Aplicamos os seguintes cortes adicionais à amostra: magnitude aparente na
	banda r entre 14.5 < mr < 17.77 e redshift > 0.002. O corte em mr vem da definiçâo
	da Main Galaxy Sample do SDSS e o limite em redshift élimina da amostra as fontes
	A historia das galâxias com formaçâo estelar intragalâtic.as.
		Definimos a amostra de galaxias SF como aquelas abaixo da linha divisoria
	no BPT definida por S06, com S/N > 3 em [O m]À5007, H/3, [N il] À6584 e Ha;, e
	S/N > 10 na regiâo de 4730-4780 Â do contmuo espectral. O numéro final de objetos
	da amostra SF é de 82302 galaxias.
	Apresento aqui uma versâo concisa do estudo por Asari et al. (2007) (de agora em
	diante A07), cujo artigo completo esta no Apêndice B. Este trabalho é uma extensâo
	do apresentado no Capitulo 3. Vârios cuidados adicionais forain tomados, como,
	por exemplo, com os residuos sistemâticos ao redor de H/3, a determinaçâo auto-
	consistente da metalicidade e da extinçâo nebular, e testes com outros indicadores de
	Zneb-	
		Caracterizamos a nossa amostra a partir de uni estudo convencional da cor-
	relaçâo de parâmetros fisicos e observacionais com Zneb. Mostramos que existe uma
	ôtima concordância entre a taxa de formaçâo estelar atual (SFR) calculada a partir
	da luminosidade em Ha e da sintese de populaçôes estelares. Esse resultado pode
	ser usado em estudos futures de formaçâo estelar em galaxias hospedeiras de nücleo
	ativo, para as quais o método mais tradicional de Ha nâo pode ser aplicado. Tanrbém
	derivamos a taxa de formaçâo estelar em funçâo do tempo para as galâxias SF divi-
	didas em faixas de metalicidade nebular (Zneb), massa estelar (M*) e densidade de
	massa (E*). Encontramos novamente que galâxias de menor massa, menos densas
	e mais pobres em metais evoluiram mais lentamente e hoje têm uma maior SFR es-
	pecffica. Mostramos que essa tendência se mantém mudando os critérios de seleçâo da
	amostra, a base de populaçôes estelares e levando em consideraçâo efeitos de extinçâo
	diferencial.
		Este artigo foi escrito em grande parte por mini, Roberto Cid Fernandes Jr. X [Â]
	e Grazyna Stasiiiska. As minhas majores contribuiçôes forain nas seguintes inves-
	tigaçôes: desenvolvimento do método auto-consistente para câlculo da extinçâo e
	metalicidade nebular e teste de outros indices para a calibraçâo da metalicidade (am-Figura 4.1: Exemple» de espectro residual para mua combinaçâo de galaxias SF. A
	linha tracejada marca o residuo nulo e pontilhada o centro da linlia H/3. Note o bos com orientaçâo da Grazyna); testes com diferentes leis de extinçâo, correlaçôes de parâmetros com Zneb, SFR em funçâo do tempo e distribuiçâo dessas funçôes, tes resfduo largo ao redor de H/3, estendendo-se por quase 200 À.
	tes dos critérios de seleçâo e extinçâo diferencial (em colaboraçâo com Roberto Cid);
	comparaçâo entre a SFR. atual da sintese e de Ha (em colaboraçâo corn Roberto Cid e
	Juan Pablo Torres-Papaqui). O estudo de resfduos em H/3 foi feito principalmente por
	Roberto Cid e Abilio Mateus, e os testes com as novas bases de populaçôes estelares
	por Jean Michel Gomes.
	4.1	Anâlise de dados
	Os dados utilizados liesse estudo sâo do DR5 do sdss. A amostra, os modelos com
	o STARLIGHT e a medida de linhas de emissâo estâo detalhadas no Capitulo 2. As

Para os ajustes com o STARLIGHT, preferimos usai* a lei de extinçâo de

Calzett.i, Kinney, & Storchi-Bergmann (1994, CAL)

, em vez da lei

de Cardelli, Clayton, & Matliis (1989, CCM) 

dos nossos outres estudos. Fizemos testes com 1000 galaxias aleatôrias com quatre leis de extinçâo diferentes, dentre as quais CAL e CCM. Nâo encontramos diferenças significativas nem na qualidade, nem nos resultados da sfntese espectral para nenhuma das leis (ver seçâo 2.5.1 de A07). Por isso, decidimos utilizar a lei de CAL, que foi desenvolvida especificamente para observaçôes integradas de galaxias SF, muito parecidas com as que utilizamos nesse estudo.

  também o efeito de abertura. No SDSS a abert.ura da fibra é a mesura para todas as galâxias, independente do seu redshift ou tanranlio. Testamos entâo duas novas amostras: selecionando galâxias mais afastadas (z > 0.05) e galâxias que tivessem pelo menos metade da luz lia banda z englobada pela fibra. Os painéis g-j mostram que essas duas seleçôes também nâo interferem nos nossos resultados.

4.7. Sumario painéis c e d ilustram o fato de que a escolha da linha de S06 para selecionar as galâxias SF no BPT nâo interfère nos nossos resultados: utilizando a linha de Kauffmann et al. (2003a) a diferença nas SFHs é quase imperceptfvel.

A inclinaçâo de uma galâxia afeta a estimativa de Zneb, porque para as galâxias mais inclinadas observamos preferencialmente as partes mais externas e, portanto, de menor metalicidade. No entanto, eliminadas as galâxias muito inclinadas na amostra, as SFHs nâo mudam muito, como mostram os painéis e e f. Estudamos Experimentos coin outras bases de SSPs, como as coin a biblioteca MILES

(Sânchez-Blâzquez et al. 2006

) em vez da STELIB, revelaram que existem algumas idades em que surgern diferenças importantes, como ao redor de 1 billiâo de anos

(Gomes 2009)

. No entanto, a separaçâo e a ordenaçâo das SFHs nâo nrudou, e o nosso resultado continua robusto nrudando a base de populaçôes estelares.

  Muitos dos espaços vazios do BPT sâo produto de efeitos de seleçâo. Galâxias aposent.adas de baixa metalicidade, por exemplo, mesmo se existissem, estariam abaixo do mvcl de detecçâo do SDSS (ver explicaçâo sobre a figura 6 ern SOS). Isso élimina os modelos entre as duas asas e abaixo da asa das SF. Usamos nesse estudo a metalicidade estelar, que esta livre dos vieses que afetam a metalicidade nebular. Encoiitrainos que, para redshifts maiores, a relaçâo M*-Z* é mais inclinada e se estica até massas e metalicidades menores.

	6.2.1 6.3	Capitula 6 A asa direita da gaivota explicada Numéro de fotons ionizantes CapituCo
		em galâxias.
	Como mostramos na Figura 6.1, em comparaçâo com as velhas, estrelas jovens domi-Jâ sabe-se bem por que a asa das SF é estreita: ali estâo interligados a metalicidade,
	nam a luminosidade em Ha. Por isso, ao longo da nossa anâlise consideramos dois o campo de radiaçâo e o parâmetro de ionizaçâo (Dopita et al. 2006). A asa direita, Esse resultado, novamente, é outra faceta do downsizing de galâxias. Mostramos que 3. Medida de abundâncias de elementos quimicos a partir de linhas de
	As LINERs ionizadas por estrelas casos: o espectro ionizante devido a todas as populaçôes estelares (caso F, do inglês full) e apenas devido a populaçôes estelares mais velhas do que 107 5 anos (caso O, de por out.ro lado, nâo tem razôes fîsicas para ser uma sequência estreita. De acordo com os modelos de fotoionizaçâo, que cobrem todo o diagrama BPT exceto a regiâo das podcmos modelar Z*(t) a partir de uni modelo simples de caixa fechada para objetos emissâo fortes. Começamos a desenvolver uma malha de modelos de fotoio-Conclusôes e Perspectivas de massa > 109'85M©. Isso implica que a forma relaçâo massa metalicidade deve ser nizaçâo que servira de base para calculai' a abundância de elementos quimicos
	old). Seyferts mais extremas, as galâxias aposent.adas deveriam aparecer por toda a parte. mais influenciada pelas diferenças em liistorias de formaçâo estelar das galâxias do em galâxias a partir de suas linhas de emissâo fortes. A nossa idéia é atuali-
	Calculamos entâo o numéro de fotons ionizantes Qhi produzido pelo espectro que por ventos. zar o método desenvolvido por McGaugh (1991), criando uma malha bem mais
	ionizante médio de cada faixa do BPT. Pelos mesmos argumentos da Seçâo 4.5, pode-Por fini, estudamos as galâxias da asa direita "das AGNs" do diagrama de fina de modelos e sem assumir uma relaçâo entre as abundâncias N/O e O/H.
	Este capitulo t.raz a anâlise e os principais resultados apresentados por Stasinska et al. mos derivar a luminosidade em Ha supondo que todos os fotons sâo absorvidos. Essa 7.1 Sumario diagnôstico BPT. Mostramos que a luminosidade em Ha e a razâo entre linhas de Encontramos diversas dificuldades inesperadas, e estamos investigando métodos
	(2008) (de agora em diante S08), cujo artigo completo esta dispomvel no Apêndice é luminosidade mâxima em Ha esperada, L(Ha)exp-emissâo de vârias galâxias classificadas como LINERs podem ser explicadas pela io-melhores para o câlculo e se as malhas (pie escolhemos sâo as mais adequadas
	D. Nesse estudo, voltamos o foco para a asa direita do diagrarna BPT. As galâxias Para cada uma das sequências de galâxias (SF, Seyfert, LINERs) definidas na Aplicamos o nosso algoritmo de sintese espectral starlight a 573141 galâxias do nizaçâo dura de estrelas veillas (pôs-AGB e anâs brancas, por exemplo). Portanto, para o nosso problema ffsico. A Figura 7.3 mostra resultados preliminares desse
	figura 1 de SOS, comparamos a L(Ha) esperada com a observada a partir da razâo método para galâxias de baixa metalicidade do SDSS.
	delas sâo aposentadas.
	6.4	Sumârio
	5.4	Sumârio
	6.2.2	6.1 Modelos de fotoionizaçâo Anâlise de dados
	Derivamos a evoluçâo em funçâo do redshift da relaçâo M+-Z+ a partir de registres Desenvolvi uni programa que calcula o continuo nebular liesses casos e apliquei
	fosseis para o mesmo conjunto de galâxias. A partir de um modelo de evoluçâo Usamos os dados do DR5 do SDSS analisados com o STARLIGHT e o nosso programa 6.2 A fonte de ionizaçao em galaxias aposentadas Ha pode ser explicado sô com populaçôes velhas. E quanto às outras linhas de ao Data Release 7 do SDSS. A Figura 7.1 traz um exemplo de ajuste sem
	quhnica simples, argumentamos que as historias de formaçâo estelar sâo o parâmetro de medidas de linhas de emissâo como descrito no Capitulo 2. Aplicamos os seguintes emissâo? Para responder essa pergunt.a, fizemos modelos de fotoionizaçâo com o e com correçâo pelo continuo nebular. Verificamos que essa é uma melhoria
	ffsico principal para a forma da relaçâo massa-metalicidade para galâxias com massa cortes adicionais à nossa amostra: 14.5 < mr < 17.77 (da definiçâo da Main Galaxy Certas tendências astroffsicas sô foram reveladas pela quantidade abismal de dados côdigo PHOTO. A radiaçâo ionizante para todos os modelos foi considerada como das suas estrelas hâ muito tempo e de forma râpida. O enriquecimento qunnico lies-técnica muito importante para galâxias de baixa metalicidade, e estudaremos
	maior do que 109,85 MQ. Esses resultados mostram unia tendência clara em Z*, mas Sample do SDSS), S/N < 10 ao redor de 4750 A (para garantir um bom ajuste com dos ültimos tempos. A forma de gaivota do diagrama BPT, por exemplo, sô apareceu tra a evoluçâo com a idade do numéro de fotons com energia > 13.6 eV (Qhi), isto sendo o espectro tfpico das das populaçôes velhas (caso F) para a faixa ir ~7 da sas ültimas também ocorreu mais depressa. Esses resultados estâo de acordo com o mais adiante as mudanças da propriedades recuperadas com a sintese para esses
	ainda nâo se sabe a confiança com que Z* pode ser recuperado pela sfntese com os a sintese), 2 > 0.002 (para excluir fontes intragalâticas) e fraçâo de luz na banda z claramente a partir de estudos do SDSS (por exemplo, Kauffmann et al. 2003a). A asa é, capazes de ionizar o hidrogênio, para uma populaçâo estelar de metalicidade solar sequência das LINERs (marcado em laranja na figura 1 de S08). Tomamos como o cenârio de doumsizing de galâxias. objetos.
	modelos de SSPs atuais. dentro da fibra > 20% (para diminuir efeit.os de abertura). direita revelon-se também mais larga do que a asa das SF. Em estudos anteriores aos de BC03. Quando essa SSP atinge 108 anos, ela ainda produz uma quantidade sig-espectro tfpico aquele correspondente à galâxia cuja dureza do campo ionizante dada Fernandes et al. 2005), e mudanças leves entre populaçôes estelares podem fazer essa A partir dos resultados da sintese, derivamos também a taxa de formaçâo
	Consideramos apenas os objetos que podem ser plotados com segurança no com os dados do SDSS (por exemplo, Ho et al. 1997) jâ se separava que essa regiâo nificativa de fotons ionizantes. O painel à esquerda da Figura 6.1 mostra a evoluçâo por Qne\/Qm fosse a mediana da distribuiçâo dessa variâvel na faixa. Mais detalhes porcentagem de galaxias aposentadas flutuar. estelar recente para as galâxias SF, calibrada a partir da luminosidade em Ha. Esse 2. Detecçâo de linhas de emissâo fracas em galâxias SF. Outra investigaçâo
	diagrarna BPT. Aplicando o corte de S/N > 3 nas quatro linhas do diagrarna, a nossa do BPT em duas famflias de objetos: Seyferts (na parte de cima) e low-ionization da razâo QneJQn\, onde Qne] é o numéro de fotons capaz de ionizar uma vez o hélio dos modelos encontram-se em S08. Por ültimo, e talvez mais importante, os modelos de populaçôes estelares para exercicio mostrou que as medidas independentes para a SFR derivadas pela sintese iniciada durante o doutorado é a de linhas de emissâo fracas em espectros com-
	amostra final conta com 131287 objetos. nuclear émission régions (LINERs, na parte de baixo). (> 24.6 V). Quanto rnaior Quei/Qni mais fotons de alta energia sâo fornecidos; ou Os modelos de fotoionizaçâo para diferentes metalicidades e paramétras de a parte do contfnuo < 13.6 eV é incerta. Para a base de SSPs de BC03 que usamos, a e pelas linhas de emissâo sâo compativeis. Essa calibraçâo ainda pode ser usada em binados de galâxias SF. Queremos combinai' os espectros de galâxias de propri
	As linhas de emissâo caracteristicas das LINERs na faixa visfvel do espectro, seja, essa razâo quantifica a dureza da fonte de ionizaçâo. As estrelas mais velhas, ionizaçâo U sâo mostrados na figura 5 de S08. Os modelos cobrem a maior parte dos uma idade de 108 anos, hâ unia variaçâo de 0.3 a 1.0 dex em QHi entre as populaçôes de galâxias hospedeiras de AGN, nas quais Ha nâo deve ser usada como indicador de edades semelhantes, como, por exemplo, pela massa em estrelas ou pelas suas

nessa asa normalmente sâo classificadas em Seyferts ou LINERs, uma nomenclatura que sugere um buraco negro central ativo

[START_REF] Ho | [END_REF]

. No entanto, argumentamos que alguns dos objetos cujo espect.ro de emissâo é semelhante ao de LINERs podern ser galâxias que pararam de formar estrelas e que nâo têm atividade nuclear. A partir dos resultados da sintese de populaçôes estelares e de modelos de fotoionizaçâo, mostramos que essas galâxias aposentadas sâo compatfveis com a ionizaçâo por estrelas queutes veillas (do pos-ramo assintotico das gigantes e anâs brancas). A redaçâo desse artigo foi feita em grande parte por Grazyna Stasinska e Roberto Cid Fernandes Jr. Nesse estudo, a previsâo de Ha a partir das populaçôes estelares foi desenvolvida por Roberto Cid e Jean Michel Cornes; eu verifiquei e refinei o procedimento de câlculo. A divisâo do diagrarna BPT em coordenadas polares e a anâlise da previsâo da luminosidade em Ha e da dureza da radiaçâo ao longo das asas foi feita por mim e Grazyna St.asiiiska. O estudo das incertezas e da forma da asa direita forain uni trabalho conjunto feito por mim, Grazyna e Roberto Cid. Os modelos de fotoionizaçâo forain feitos pela Grazyna. A fini de estudar tendências ao longo do BPT, dividimos o diagrarna em coordenadas polares (figura 1 de SOS). O centro do eixo de coordenadas foi escolhido para coincidir com o ponto de inflexào da mediana da distribuiçâo de log [O iii]/H/3 em funçâo de log [N Ii]/Ha para a nossa amostra. O diagrarna foi dividido em 40 partes no ângulo 9 e 18 partes no raio r. A asa correspondente às galâxias SF de menor

log t* [yr] log t* [yr] Figura 6.1: Evoluçâo com o tempo do numéro de fôtons ionizantes Qhi e da dureza da radiaçào, quantificada pela razâo Quei/Qm-Essas curvas foram produzidas para uma populaçâo estelar de BC03 com metalicidade solar formada em um burst instantâneo. Note que as populaçôes velhas (t > 108 anos) ainda produzem fôtons capazes de ionizar o hidrogênio, mas em uma quantidade de 4 a 5 ordens de grandeza menor do que quando jovem. Por outro lado, a dureza da ionizaçâo alimenta com a idade. Ambos efeitos, diminuiçâo do numéro de fôtons capazes de produzir Ha e aumento da dureza da radiaçào, podem produzir um espectro de emissâo parecido com o das LINERs. metalicidade esta em 6 ~117°(marcada em azul claro), a de maior metalicidade em ~-94°(azul escuro), as Seyferts em ~77°(verde) e as LINERs em ~59°(vermelho). Analisamos cada uma dessas faixas em funçâo da posiçâo radial (ir = 0 a 17). no entanto, nâo sâo exclusivas de galaxias com atividade nuclear, valendo às LINERs a alcunha de "caixinha de surpresas". Outros fenômenos ffsicos podem fornecer uma fonte de ionizaçâo dura e produzir espectros parecidos com os de LINERs, como ondas de choque no gâs e explosôes de supernova em regiôes de formaçâo estelar ao redor do Figura 6.2: Variaçâo com a posiçâo radial ir da razâo entre luminosidade de Ha esperada e a observada. Esta é uma simplificaçâo da figura 2 de S08: sâo graficadas apenas os quartis e medianas para a sequêneia das LINERs e caso O (isto é, levando em conta apenas a produçâo de Ha por populaçôes mais velhas do que 10 ' 5 anos). nücleo de uma galâxia (Heckman 1987). Portanto, sem evidências de atividade nuclear em outros comprimentos de onda, uma galâxia na regiâo das LINERs no diagrama BPT nâo necessariamente é hospedeira de AGNs, ao contrario do que sugere o terrno "asa direita das AGNs". Consideremos uma outra fonte de ionizaçâo que produz espectros do tipo LI-NER. Trinchieri h di Serego Alighieri (1991), Binette et al. (1994) e Taniguchi, Shioya, h Murayama (2000) encontraram que estrelas queutes na fase pôs-ramo assintotico das gigantes (pôs-AGB, do inglês asymptotic giant, brandi) e anâs brancas podem pro ver uma fonte de ionizaçâo compatfvel com as linhas de emissâo em galâxia elfpticas com pouco gâs e sem formaçâo estelar atual. O painel à esquerda da Figura 6.1 mosportanto, produzem uni espectro de radiaçâo mais duro do que estrelas jovens. A pergunta que nos propomos responder é se as LINERs do SDSS podem ser explicadas apenas por populaçôes estelares. Para isso, estudamos o espectro ionizante das populaçôes estelares recuperadas com o STARLIGHT. L(Ha)exp/L(Ha)obs para os casos F e O. A figura 2 de S08 mostra esses resultados. Vamos nos concentrai em duas sequências: a de SF caso F (em azul escuro, linhas continuas) e a de LINERs caso O (em vermelho, linhas pontilhadas). Para as SF, nâo temos nenhum problema em explicar a emissâo em Ha usando todas as populaçôes estelares. Isso é mais uni teste de sanidade para verificar se as populaçôes recuperadas com o STARLIGHT sâo razoâveis, e é outra forma de visualizar a compatibilidade entre a SFR atual derivada por Ha e pela sfntese (ver Seçâo 4.5). A Figura 6.2 mostra uma versâo simplificada da figura 2 de S08, apenas para o caso O das LINERs. Para essa sequência, as populaçôes estelares velhas chegarn a explicar a emissâo em Ha para até 25% dos objetos (L(Ha)exp/Z/(Ha)0bs > 0.5 para as faixas ir > 7). E ôbvio que essa porcentagem é incerta por diversos motivos (ver Seçâo 6.2.3), mas ainda assim révéla que uma porçâo dos objetos chamados LINERs podern ser galâxias a,posentadas -isto é, que pararam de formar estrelas -cujo gâs é ionizado por estrelas velhas e nâo por uni nücleo ativo. diagramas de diagnôstico, incluindo a regiâo onde ficam as LINERs. Portanto, todas as razôes de linhas desses diagramas podem ser reproduzidas a partir das SSPs velhas, a saber [O hi]A5007/H/?, [N n]A6584Ha, [O n]À3727/H/?, [S n]AA6717+6731/Ha e [O i]A6300/Ha. log t+ [yr] Figura 6.3: Como o painel à esquerda da Figura 6.2, mas para diferentes populaçôes estelares de metalicidade solar: BCU3, em vermelho; PEGASE coin caminhos evolutivos de Bressan et al. (1993) em azul; e PEGASE com caminhos evolutivos de Schaller et al. (1992) e Charbonnel et al. (1996) em verde. Ao redor de 108 anos a diferença entre os modelos chega até a 1 ordem de magnitude e cai para 0.2-0.3 dex em idades mais avançadas. 6.2.3 Incertezas Como vimos, os nossos modelos de fotoionizaçâo com populaçôes veillas explicam bem as razôes de linlias das LINERs. No entanto, para apenas 25% das nossas LINERs conseguimos explicar a luminosidade em estrelas vellias. Uma maneira precipitada de interprétai' esse numéro é dizer que uni quarto das LINERs no universo local sào, na verdade, galaxias aposentadas. No entanto, essa quantidade é apenas um indicativo de que devem existir galaxias aposentadas passando-se por LINERs na nossa amostra, e nâo deve ser tomada como um numéro estatisticamente significativo. Em primeiro lugar, a nossa amostra deveria ser compléta em volume para ser representativa do que acontece no Universo local. Em segundo lugar, a propria sintese tem incertezas associadas (Cid menor e maior metalicidade. Considerando bases diferentes, como a BC03 e a PEGASE (Fioc & Rocca-Volmerange 1997), as diferenças vâo também de 0.2 a 1.0 dex para a metalicidade solar, como mostra a Figura 6.3. Desconfiamos (pie essas diferenças sejam dévidas principalmente à relaç.âo massa inicial final em anâs brancas. Também argumentamos que, abaixo da gaivota, a sequência das SF deveria continuai-, sendo preenchida por galâxias de maior metalicidade. O fato é que essas galâxias sâo tâo ricas em metais que a temperatura eletrônica fica muit.o baixa e, a partir de uni certo limite, [O m] A5007 nâo pode mais ser excitado. Na asa direita, tanto nücleos ativos quanto estrelas velhas produzem uma radiaçâo dura o suficiente para évitai' esse fenômeno. Achamos galâxias de baixo [O m] na nossa amostra, e elas sâo compativeis com a sequência de massa-met.alicidade das SF (figura 7 de S08). Outras galâxias est.âo ainda faltando no BPT, como as de baixo S/N em H/?, por exemplo. Essas galâxias sâo estudadas junto com outras de linhas de emissâo fracas por Cid Fernandes et al. (2010), onde mostramos que a esmagadora maioria A classificaçâo de galâxias como LINERs com base somente em razôes de linhas de emissâo ôpticas cria, na verdade, uni grupo heterogêneo de objetos sujeitos a diferentes cenârios iïsicos. Mostramos que é possfvel explicar uma fraçâo significativa de LINERs do SDSS como sendo ionizadas por populaçôes velhas apenas, sem a necessidade de nücleos ativos. Este foi o primeiro estudo a considerar esse cenârio no SDSS e a calcular a ionizaçâo a partir das SSPs recuperadas a partir de métodos fôsseis (em contraste com as populaçôes estelares escolhidas por Trinchieri & di Serego Alighieri 1991, Binette et al. 1994 e Taniguchi et al. 2000). Em consequência disso, os estudos que consideram todas as LINERs como hospedeiras de AGN devem ser revistos. Data Release 5 do SDSS, utilizando uma base de 150 populaçôes estelares de diferentes idades e metalicidades de BCQ3. Dessa anâlise derivamos varias propriedades fïsicas desses objetos, tais como a massa em estrelas, a atenuaçâo por poeira, a idade e a metalicidade média das populaçôes estelares. Subtraindo a contribuiçâo estelar do espectro obscrvado, obtivemos o espectro puramente nebular. Medimos entâo no espectro residual as linhas de emissâo coin uni algoritmo automâtico. A partir dessas medidas, obtemos informaçôes sobre a abundância de elementos quimicos no gâs, a extinçâo, e até mesrno a temperatura, a taxa de formaçâo estelar atual e a fraçâo de gâs em relaçâo à massa total da galâxia para alguns objetos. Ainda, podemos usai* diagramas de diagnôstico, como o [N il]/Ha. vs. [O m]/H/3 para classificar galâxias em diferentes tipos, como SF, Seyfert ou LINERs. Aproveitamos essa quantidade fenomenal de dados para estudar as proprieda des e a evoluçâo intnnseca das galâxias com formaçâo estelar. A partir das populaçôes estelares, derivamos a taxa de formaçâo estelar em funçâo do tempo para galâxias S F marginalizado a metalicidade. Também derivamos a metalicidade estelar média em funçâo do tempo. Encontramos (pie as galâxias cujo meio interestelar é menos enriquecido e de menor massa estelar atual forain mais lentas em formar estrelas, e lioje têm uma taxa de formaçâo estelar especifica mais alta. O enriquecimento qunnico das estrelas também se deu de forma mais graduai nesses objetos. Por outro lado, as galâxias de maior metalicidade nebular e de maior massa formaram grande parte formaçâo estelar porque essa linha também é produzida pela ionizaçâo pela fonte central, e nâo apenas por estrelas jovens e queutes. Estudamos ainda a evoluçâo côsmica das galâxias do SDSS. Obtivemos a evoluçâo da relaçâo massa-metalicidade dos redshift 0.1 a 1.9 para galâxias de todos os tipos do SDSS. algumas LINERs do SDSS nâo devein ter atividade nuclear. Estudos sobre AGNs devem levar em conta essa populaçâo de galâxias aposentadas que parecem LINERs, a fim de évitai-conclusôes precipitadas sobre a evoluçâo de galâxias e a relaçâo entre a atividade nuclear e formaçâo estelar no Universo local. 7.2 Trabalhos em curso e futuros

Listo a seguir alguns dos projetos em curso que aprofundam as técnicas e os estudos da evoluçâo de galâxias a partir de métodos fôsseis. 1. Continuo nebular. Uma das ferramentas que usamos em muitos dos nossos estudos é a sfntese de populaçâo estelares com o algoritmo starlight. Esse algoritmo assume que o espectro de uma galâxia é uma mistura de populaçôes estelares simples e acha a mellior combinaçâo dessas populaçôes. Isso é bastante razoâvel para a maioria das galâxias do SDSS. No entanto, para galâxias de metalicidade muito baixa, sabemos que o gâs pode contribuir para o continuo espectral em uma proporçâo nâo desprezivel quando comparado as estrelas. linhas de emissâo fortes. Essa é uma maneira de aumentar a razâo sinal-ruido e detectar linhas fracas que hcavarn escondidas no ruido dos espectros individuais. Jâ hz alguns testes preliminares e consegui, com esse método, detectar a linha de [N il] À5755 em galâxias SF (uni exemplo é mostrado na Figura 7.2), que pode ser usada para calculai diretamente a abundância de elementos quimicos 4. A populaçâo de galâxias "esquecidas". Estudamos as propriedades ffsicas de galâxias de linhas de emissâo fracas. Esse tipo de objeto é geralmente negligenciado em estudos de galâxias coin linhas de emissâo, mas mostramos que eles sâo uma populaçâo muito numerosa e importante para entender o ciclo de vida das galâxias locais. Jâ submetemos o primeiro artigo sobre esse tema (Cid Fernandes et al. 2010) e estamos escrevendo outro com ênfase lia natureza ffsica dessas galâxias. A Figura 7.4 mostra como as LINERs e as Seyferts diferem em termos de massa, £ = L(Ha)0^s/[L(Ha)e^p(t > 106 * 8yr)], dispersâo de velocidades e histôrias de formaçâo estelar. Queremos também verificar se as

  Ajuste com o STARLIGHT de uma galâxia SF de baixa metalicidade. Os painéis superiores mostram os resultados sein subtrair o contmuo nebular, e os inferiores com a subtraçâo. Nos painéis à esquerda, o espectro utilizado no ajuste esta em preto, o modelo em vermelho e as regiôes mascaradas em verde. O contmuo nebular, subtrafdo do espectro observado no painel inferior, é mostrado em laranja.

			7.2.	Trabalhos em curso e futuros
			log [O III]4363/[0 III]5007 [109]
	[ n o r m a l i z e d ]		
	log |0 IIJ3727/Hp [109]	log [O III]5007/H(3 [109]		log [Ar III]7135/HJ3 [109]
	X [A]	X [A]	log t*[yr]
	6. Outras faixas espectrais. Queremos estender os nossos estudos para outras comprimentos de onda, para obter mais inforrnaçôes sobre o gâs, as populaçôes estelares, os AGNs, jatos, poeira, e variabilidade na luminosidade das galâxias em redshifts baixos. Para isso, iremos procurai' objetos no SDSS que estejam Figura 7.2: Espectro residual para uma combinaçâo de galâxias SF. A linha pontilhada niarca o centro da linha de emissao [N n]À5755, que pode ser usada para medir Figura 7.1: Os painéis à direita mostram a fraçâo de luz (em %) recuperada para cada populaçâo diretamente a abundância nebular de elementos qunnicos.
	em outras catâlogos, como o G ALEX ( ultraviolet a). Além disso, as galâxias em estelar. 0 ajuste para o espectro do (pial foi subtrafdo o contmuo nebular récupéra
	redshifts mais altos têm o seu espectro optico deslocado para o infravermelho; populaçôes estelares menos jovens.
	para estudâ-las, devemos saber lidar com os problemas e vantagens que dessa
	faixa espectral. Para isso, precisaremos modificar o STARLIGHT para ajustai'
	bandas espectrais e espectros amostrados inomogeneamente.

5. Tendências ao longo da relaçâo massa-metalicidade. As Figuras 7.5 a 7.12 mostram como vârios parâmetros ffsicos variam com a relaçâo massa estelar-metalicidade nebular. Esses resultados preliminares forain os que primeiro motivaram o estudo da relaçâo AI+-Z+. Um dos grandes desafios desses estudos é entender quais relaçôes sâo realmente ffsicas e quais os paramétras que causam essas correlaçôes.

  of Stasinska (2006). From that figure, the O/H given by équation (1) lies towards the upper end of the possi

	ble values of O/H. We considered using the [Arm]7135/[0m]5007
	(Ar303) index which, as argued by Stasinska (2006), does not suf-
	fer from the problems mentioned for the 03N2 index. It turns out
	that, in objects where the [Arm] line could be measured. Ar303 and
	03N2 are extremely well correlated (with a Spearman corrélation
	coefficient of Rs = 0.58). Unfortunately, the quality of the SDSS
	spectra did not allow us to measure the [Arm] line intensity with
	sufficient accuracy in a large number of objects, and principally in
	the zone where it would hâve been helpful to break the 03N2 versus
	O/H degeneracy. Using the Pilyugin & Thuan (2005) metallicity
	calibration based on [Oiii]5007/H/3 and [On]3727/H/J adds only
	a tiny fraction of galaxies. The same applies when using the O/H
	values obtained by Izotov et al. (2006) from direct methods using
	the électron température derived from [O iii]4363/[0 m]5007. When
	comparing our measures with theirs, there is a systematic offset of
	0.2 dex and a rms of 0.13 dex in Zneb for the 177 objects we hâve in
	common. This effect is greater the lower Z,,eb is. We thus decided to
	use 03N2 as a nebular metallicity indicator ail along the SF galaxy
	sequence, keeping in mind that équation (1) will tend to attribute
	a metallicity Zneb of about 0.2 Z© for the galaxies with the lowest
	observed 03N2 in our sample.

Table 1 .

 1 Statistics of properties in bins A-F.

			Bin A	Bin B	BinC	Bin D	Bin E	Bin F
	logZneb min(Zo) -0.710 -0.450 -0.320 -0.190 -0.060	0.070
	logZneb p05	-0.608	-0.435	-0.311	-0.180	-0.053	0.071
	logZneb p50	-0.494	-0.364	-0.242	-0.111	-0.001	0.081
	log Zncb p95	-0.454	-0.324	-0.194	-0.064	0.054	0.112
	logZneb max	-0.450	-0.320	-0.190	-0.060	0.070	0.200
	log yW* p05 (M©)	7.146	7.938	8.693	9.277	9.800	10.089
	log	p50	8.319	8.922	9.452	9.958	10.460	10.678
	log M* p95	9.313	9.634	10.083	10.646	11.073	11.154
	(logr*)L p05 (yr)	6.755	7.428	7.695	7.915	8.127	8.167
	(logr*)L p50	7.762	8.155	8.385	8.567	8.726	8.724
	(log Ml p95	8.556	8.900	9.139	9.255	9.276	9.255
	logZ.Hr/ p05 (L©)	5.167	5.139	5.492	5.988	6.551	6.956
	logLn" p50	6.475	6.475	6.750	7.149	7.544	7.816
	log Z.] I(/ p95	7.834	7.736	7.961	8.172	8.426	8.559
	log b p05	-0.157	-0.392	-0.698	-0.947	-0.973	-0.890
	log b p50	0.769	0.436	0.211	-0.006	-0.213	-0.199
	log b p95	1.527	1.192	0.905	0.644	0.366	0.272

of Zneb using this index defines a straight line of unit slope in the B PT diagram.

  Of ail selection-induced changes discussed here, this is the largest one; yet, ail it does is to shift the SFHs from one group of galaxies to the adjoining one.To summarize, sélection criteria may influence the derived mean SFHs in quantitative ternis, but do not modify the relative pattern of mean SFHs of galaxies in different Zneb bins. The same applies to grouping galaxies according to properties other than Z"cb. The general trends in the SFH as a function of global galaxy properties obtained in this work are therefore robust against variations in the sample sélection criteria. comesout AA* ~0.4maglargerwith the MILES models. Extensive tests with these new bases, including different prescriptions of stellar évolution as well as different spectral libraries, are underway, but these first results show that significant changes can be expected. The 17 142 galaxies in the SFhq subsample were refitted with this more realistic modified recipe for extinction effects.

	Reassuringly, however, these same experiments show that the
	pattern of SFHs as a function of Zneb,M* and E* reported in this
	paper does not change with these new models. In any case, these
	initial results provide an éloquent reminder of how dépendent semi-
	empirical SFH studies are on the ingrédients used in the fits.
	73 Experiments with differential extinction
	We saw in Section 2.5.2 that the nebular and stellar extinctions are
	strongly correlated, but with A"eb ~twice A*. This indicates that the
	uniform stellar extinction used in our fits is not adéquate to model
	SF régions, which should be subjected to a similar extinction than
	the line-emitting gas, i.e. A*(f* < 107 yr) ~A£eb. It is therefore fit
	to ask whether and how such a differential extinction affects our
	general results.
	Given the difficulties in recovering reliable population-dependent
	A* values from spectral synthesis in the optical range alone, we
	address this question by postulating that populations younger than
	107 yr are extincted by 034 + 2.28 A*, i.e. the empirical A£,cb (A*)
	relation found in Section 2.5.2, with A* now denoting the extinction
	to > 107 yr stars.3

  as a function of Zneb,M*, E* remains the same. Experiments with new evolutionary synthesis models and differential extinction fits were reported and found to lead to substantially different SFHs, yet preserving this same overall pattern.

		ftpendice
		ACKNOWLEDGM ENTS Artigo: The évolution of the
		We are greatly in debt with several colleagues and institutions around
	the globe who hâve contributed to this project by allowing ac-cess to their computers. The starlight project is supported by the mass-metallicity relation in SDSS galaxies
		Brazilian agencies CNPq, CAPES, FAPESP, by the France-Brazil
		CAPES/Cofecub program and by Observatoire de Paris. Funding for the SDSS and SDSS-II has been provided by the Al uncovered by astropaleontology
		fred P. Sloan Foundation, the Participating Institutions, the National
		Science Foundation, the US Department of Energy, the National
		Aeronautics and Space Administration, the Japanese Monbuka-
		gakusho, the Max Planck Society, and the Higher Education Funding
		Council for England. The SDSS web site is http://www.sdss.org/.
	The SDSS is managed by the Astrophysical Research Consortium Artigo por Vale Asari et al. (2009) (DOI: 10.1111/j.1745-3933.2009.00664.x). Também
	for the Participating Institutions. The Participating Institutions are dispomvel em versào preprint (arXiv:0903.4167).
		the American Muséum of Natural History, Astrophysical Institute
		Potsdam, University of Basel, University of Cambridge, Case West
		Institute for Astronomy (MPIA), the Max-Planck-Institute for As
		trophysics (MPA), New Mexico State University, Ohio State Univer
		sity, University of Pittsburgh, University of Portsmouth, Princeton
	INTRODUCTION	University, the United States Naval Observatory and the University
		of Washington.

ern Reserve University, University of Chicago. Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-

Por completeza, estudamos ainda outras estimativas de metalicidade, embora elas nâo sejani tâo üteis para o estudo do problema da degenerescência entre O3N2 e O/H para baixas metalicidades. Uni dos métodos utilizados foi o delineado porTremonti et al. (2004), que obteve metalicidades nebulares por comparaçâo de razôes de linhas observadas com uma grande base de dados de modelos de fotoio nizaçâo. Embora pareça uni método razoâvel,Yin et al. (2007) mostram que existem alguns problemas com essa abordagem. Nos encontramos uma diferença sistemâtica de -0.28 dex e uni rnis de 0.09 entre as nossas medidas para Znet e as deles. A calibraçâo dePettini & Pagel (2004) é semelhante à nossa, jâ que utilizam as razôes [O m]/H/3 e [N il]/Ha para a medida de O/H. Existe uma pequena diferença média de -0.04 dex com respeito à nossa calibraçâo, e uma dispersâo de 0.03 dex para a amostra SF. Dessa forma, embora consideremos a nossa calibraçâo mais confiâvcl e mais simples para os propôsitos deste estudo, repetimos os nossos câlculos com as calibraçôesde Tremonti et al. (2004) ePettini & Pagel (2004). Os resultados obtidos diferiram em escala absolut.a, mas as conclusôes qualitativas continuaram as mesmas.

© 2006 The Authors. Journal compilation © 2006 RAS, MNRAS 375, L16-L20

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*E-mail: natalia@astro.ufsc.br fSemi-Empirical Analysis of Galaxies.© 2007 The Authors. Journal compilation © 2007 RAS

© 2007 The Authors. Journal compilation © 2007 RAS, MNRAS 381, 263-279

Note that the models take into account collisional excitation of Balmer lincs, so that at low metallicities the intrinsic Ha/H/J is different from the pure case B recombination value. © 2007 The Authors. Journal compilation © 2007 RAS, MNRAS

381, 263-279

© 2007The Authors. Journal compilation © 2007 RAS, MNRAS 381, 263-279

A/h is 1.66 times smallerfor a Salpeter IMF within the same mass Iimits. © 2007 The Authors. Journal compilation © 2007 RAS, MNRAS 381, 263-279

Experiments with different modelsOne should also ask to which extent our results are robust against changes in the base of evolutionary synthesis models. While answering this question requires an in-depth study far beyond the scope of this paper, we believe this lias a much larger impact on SFHs than sélection effects.Panter et al. (2007) reported results of moped experiments using spectral models from different sources(Bruzual A. & Chariot 1993; Fioc & Rocca-Volmerange 1997; BC03;Jimenez et al. 2004;Maraston 2005), ail sampled at AX = 20 Â. For the 767 galaxies in their randomly selected test sample, the resulting mean star for mation fractions (analogous to our /i vector) differ by factors of a few for the youngest and oldest âges, and close to a full order of magnitude for /* ~0.1 -1 Gyr. Recovering SFHs in this intermediate âge range is particularly hard, as discussed byMathis et al. (2006). Indeed, the experiments reported by Panter et al. often find a suspiciously large 1 -Gyr component. The behaviour of our starlight fits is also anomalous in this âge range. This is clearly seen in the mean SFHs shown in Figs 8-10, particularly in the xs représentation, which shows a hump at ~1 Gyr. Interestingly, starlight experiments with a base of evolutionary synthesis models using the MILES library of Sanchez-Blazquez et al. (2006, instead of the STELIB library used in the base adopted for this and previous SEAGal studies) do not produce this hump at ~1 Gyr. In fact, the whole mean SFFIs derived with this new set of models differs systematically from those shown in Figs 8-10. The mass assembly function /)*(/*), for instance, rises more slowly and converges at somewhat smaller r* than those obtained in this work. There are also systematic différences in stellar extinction, which © 2007 The Authors. Journal compilation © 2007 RAS, MNRAS 381, 263-279

Papaqui J. P., Schoenell W" 2007. MNRAS,375, L16 (SEAGal V) © 2007 The Authors. Journal compilation © 2007 RAS, MNRAS 381, 263-279

This paper has been typeset from a TfX/UTjX file prepared by the author. © 2007 The Authors. Journal compilation © 2007 RAS, MNRAS 381, 263-279

© 2009 The Authors. Journal compilation © 2009 RAS

© 2009 The Authors. Journal compilation © 2009 RAS, MNRAS

This paper has been typeset from a TpX/IéTpX file prepared by the author. © 2009 The Authors. Journal compilation © 2009 RAS, MNRAS

© 2008 The Authors. Journal compilation © 2008 RAS

© 2008 The Authors. Journal compilation © 2008 RAS, MNRAS 391, L29-L33

This paper has been typeset from a TgX/LTgX file prepared by the author. © 2008 The Authors. Journal compilation © 2008 RAS, MNRAS 391, L29-L33

ACKNOWLEDGMENTS

We thank the Brazilian agencies CNPq and FAPESP for their sup port. We are greatly in debt to several colleagues and institutions around the globe who hâve contributed to this project by allowing access to their computers. The Sloan Digital Sky Survey is a joint project of The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, the Johns Hopkins University, the Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Princeton University, the United States Naval Observatory and the University of Washington. Funding for the project has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the Na tional Aeronautics and Space Administration, the National Science Foundation, the US Department of Energy, the Japanese Monbukagakusho and the Max Planck Society. ACKNOWLEDGMENTS The starlight project is supported by the Brazilian agencies CNPq, CAPES, FAPESP, by the France-Brazil CAPES-COFECUB program, and by Observatoire de Paris. We thank Laerte Sodré, Cristina Chiappini and Nikos Prantzos for their helpful comments on the manuscript. We are also grateful to the referee for interesting sug gestions. The Sloan Digital Sky Survey is a joint project of The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, the Johns Hopkins University, the Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy, the Max-Planck-Institute for Astrophysics, New Mexico State Univer sity, Princeton University, the United States Naval Observatory and the University of Washington. Funding for the project has been provided by the Alfred P. Sloan Foundation, the Participating In stitutions, the National Aeronautics and Space Administration, the National Science Foundation, the US Department of Energy, the Japanese Monbukagakusho and the Max Planck Society. ACKNOWLEDGMENTS This work was supported by the CAPES-COFECUB program. The STARLIGHT project is supported by the Brazilian agencies CNPq, CAPES and FAPESP. The Sloan Digital Sky Survey is a joint project of The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, the Johns Hopkins Univer sity, the Los Alamos National Laboratory, the Max Planck Institute for Astronomy, the Max Planck Institute for Astrophysics, New Mexico State University, Princeton University, the United States Naval Observatory and the University of Washington. Funding for the project has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the US Depart ment of Energy, the Japanese Monbukagakusho and the Max Planck

Society. We thank the referee for insightful comments.

doi : 10.1111/j. 1745-3933.2009.00664.x The évolution of the mass-metallicity relation in SDSS galaxies uncovered by astropaleontology N. Vale Asari,1,2* G. Stasinska,2 R. Cid Fernandes,1 J. M. Gomes1,3 M. Schlickmann,1 A. Mateus4 and W. Schoenell1 (the SEAGal collaboration)! JZLpêncfice D Artigo: Can retired galaxies mimic active galaxies? Clues from the Sloan Digital Sky Survey Artigo por Stasiriska et al. (2008) (DOI: 10.1111/j. 1745-3933.2008.00550.x). Também dispomvel em versào preprint (arXiv:0809.1341). These considérations imply that détectable retired galaxies must hâve metallicities of the order of solar or larger, and be massive.

As shown in Figs 7(a) and (b), this is in agreement with the mean stellar metallicities (Z*) and masses (A/*) obtained using starlight for those galaxies.

Finally, why are there no SF galaxies with such high metallici ties in the BPT diagram? The first idea that cornes to mind is that, because of downsizing, such galaxies hâve already stopped forming stars. There is another possibility, of different nature: nebular metallicities above ~3 Zq lower the électron température so much that [Ohi]Z5007 cannot be excited (this does not happen in the right wing, since the harder ionizing radiation field produces more efficient heating and compensâtes the cooling caused by the large métal abundance). Interestingly, there are emission-line galaxies in the SDSS that could well be SF with metallicities larger than 2-3 Zq, i.e. the ancestors of the most metal-rich retired galaxies: they Note that the concept of the retired galaxies applies also to galax ies that belong to the red sequence studied by Graves et al. (2007) and hâve such a weak Hft that they do not appear in the BPT dia gram, and [N ii]À6584/Ha ratios similar to those of L1NER galaxies.

As a matter of fact, there is a numerous population of such galaxies in the SDSS (over 40 000 using the sélection criteria of Section 2.1).

The weaker Ha équivalent widths of these galaxies are easily reproduced by our models for retired galaxies.

CONCLUSION

We hâve shown that retired galaxies can account for a good part of the LINER branch in the seagull's right wing. How much exactly is difficult to say, given the uncertainties. If, as we argue, a significant fraction of the LINER galaxies are in fact retired galaxies -and not active galaxies as generally claimed -the perception of the local Universe would be drastically changed. Nuclear activity would not be as common as thought, and some of the work based on SDSS data and related to the AGN population will hâve to be reconsidered.