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(1)

Reasons to measure Oxygen abundances

in the Universe

Nikos  Prantzos  

Institut G·$VWURSK\VLTXH de Paris

(2)

Galactic Chemical Evolution ingredients

   Stellar  properties  

     (function  of  mass  M  and  metallicity  Z)  

-­ Lifetimes

Yields

(quantities of elements ejected) -­ Masses of residues (WD, NS, BH)

   

Collective  Stellar  Properties  

-­ Star Formation Rate (SFR) -­ Initial Mass Function (IMF)

Gas  Flows    

-­ Infall

-­ Outflow

-­ Radial inflows in disks LOW MASS

MASSIVE

INTERMEDIATE MASS

Initial Mass Function

Lifetimes of stars (Myr)

SN PN

He3, He4, N14, C12, C13, O17, F19, s-­

C12 to Fe-­peak, weak s-­, r-­, p-­

(3)

  STELLAR YIELDS M(M

ٖ

)  

  Yields of intermediate and low mass stars

  Yields of massive stars

(main  oxygen  producers)    

(4)

SN19 87 A

Yields of Oxygen : Physical uncertainties

(

stellar physics: mass loss, convection ; nuclear physics«

(5)

IMF

SN19 87 A

Interpolation in metallicity (especially in

the range -‐4 < [Fe/H] < -‐2

Extrapolation in mass (M>40 Mʘ for WW95 and CL04)

What is the fate of stars in that

mass range?

«DQGXQFHUWDLQWLHVIURPWKH:$<WKH\LHOGV$5(86('«

(6)

«DQGLQVRPHFDVHVWKRVHXQFHUWDLQWLHVPD\QRWHYHQWKHPRVWLPSRUWDQW

Filled : Zٖ

Open: <10-­4 Zٖ

Uncertainties in BOTH O yields

(C12(Į,Ȗ)  rate,  convection,..)   AND

Fe yields

(H[SORVLRQ³PDVVFXW´

(7)

Closed  box        

or  Outflow  with  ROUT  =  k    SFR   and    SFR  =  Ȟ    mG  

Infall  with  mG    =  constant   and    SFR  =  Ȟ    mG  

  Initial

conditions

     m

G

 =  m

T

 =  1          m

T

 =    m

G    

 =  m

G,0      

 

Gas Mass

m

G  

Total mass

 

m

T  

H,He,Metals  

X  -­  X

0  

Deuterium

 

X  /  X

0  

e à÷(1àR+k)t m G;0

1+k

1+k e

à÷(1+k)t

m G;0 + (÷m G;0 à R)t p i 1àR+k 1 àR ln m

G

ð ñ 1

p

i

ð 1 à e

(1àû1)

ñ m G

1àR+kR

1 à R(1 à e

11à1=ûàR

)

Analytical  solutions  with  Instantaneous  Recycling  Approximation  (IRA),     as  a  function  of   Return  fraction  R   and   yield  p  

In  all  cases:  Gas  fraction  ı  =  m

G  

/  m

T      

and      stars  m

S

 =  m

T

 -­  m

G

     

(8)

Application:  Constraining  the  oxygen  yield  in  solar  neighborhood    

:  

(9)

Oxygen  in  the  local  Galaxy    

Important  to  know  that  the   young,  local   Galaxy     has  a   homogeneous  chemical  composition  

Stasinska  et  al.  2012  

(10)

Age -­ Metallicity Metallicity distribution

O/Fe vs Fe/H

Solar  

Neighborhood    

Constraints:  

  Local    

column  densities     Of  gas,  stars,  SFR  

 

as  well  as  

O/Fe vs Fe/H (requires SNIa

for late Fe) Age ² Metallicity

(uncertain)

Metallicity distribution

(requires slow infall)

(11)

Abundances at Solar system formation

(Massive stars: Woosley+Weaver 1995; Intermediate mass stars: van den Hoek+Gronewegen 1997;

SNIa: Iwamoto et al. 2000)

The Solar Neighborhood

Goswami and NP 2000

Is the Sun representative of the local ISM 4.5 Gyr ago ?

(12)

The Galactic Bulge

NP in Stasinska et al. 2012

Barbuy 2010

Bensby et al. 2010

Zoccali et al. 2008

Age -­ Metallicity Metallicity

distribution

One-­zone  models  of    bulge  evolution     involve    a  closed box model,   a  short timescale of  star  formation    

(~ 1 Gyr)  

and  an  enhanced SF efficiency

Bulge Thick disk

Thin disk

Minniti and Zoccali 2008

The late decrease of į/Fe at [Fe/H] > -­0.5 suggests a rapid evolution of the bulge

(more rapid than in the case of thin or thick disks)

(13)

Gonzalez et al. 2011

Recent observations of the bulge display a gradient of the mean metallicity and of [Ƚ/Fe] with distance from galactic plane

Bulge regions away from the plane are less evolved chemically

Low latitude

High latitude

High [Ƚ/Fe]

Low Fe/H

Low [Ƚ/Fe]

High Fe/H

Observations OK with monolithic  collapse  or  early  mergers  

Short timescale disfavors  secular  evolution  (e.g. through the action of a bar)

(14)

Stasinska  et  al.  2012  

The Galactic Halo

Important to know observational trends

of O/Fe with Fe/H (  to  understand  

how  the  O  and/or  Fe  yields   vary  with  metallicity)

and the amount of dispersion (to    probe  the  degree  

of  homogeneity  

of  chemical  evolution)  

(15)

NP in Stasinska et al. 2012

The Galactic Halo

(16)

Vertical structure of the Solar neighborhood: the issue of the THICK DISK

,VLWD´SXIIHG-­XSµSDUWRIWKHGLVN

which was formed early on ? (In  which  case  its  evolution  is  the    

HDUO\HYROXWLRQRIWKHGLVN «  

Or has it been formed elsewhere (satellite galaxies) and accreted later?

(In  which  case  its  evolution  is     Independent  of    the    

HYROXWLRQRIWKHWKLQGLVN «  

Or has it been formed from old stars of the inner disk

which migrated here?

(In  which  case  its  evolution  is     coupled  to  the  one    

RIWKHWKLQGLVN «  

TOTAL STELLAR DISK THIN

DISK

GAS THICK DISK

(17)

Reddy 2010

x   Thin disk

Metal rich thick disk Metal weak Thin disk Thick disk

Halo

The į/Fe (or O/Fe) ratio evolves differently in the thin and thick disks What implications for their evolution ?

Selection criteria important for attributing stars to thin or thick disks

(18)

0.   0.1   0.2   0.3   0.4   0.5   Lee et al. 2011

(17000 SEGUE stars)

Adopting chemical criteria (high vs low į/Fe) instead of  kinematical  ones for thin/thick separation leads to a clear distinction of their spatial/kinematical properties

Lee et al. 2011

(19)

However, in that case,

after proper weighting

of the iso[į/Fe]

SRSXODWLRQV«

«LWLVIRXQGWKDWWKHYHUWLFDO

surface density profile in the solar neighborhood

is smooth,

i.e. it does not display the

FKDUDFWHULVWLF´GRXEOHH[SRQHQWLDOµ SURILOHRIWKLQWKLFNGLVNV«

(Bovy  et  al.  2011)

Bovy et al. 2011

(20)

     

       

   

 

The Milky Way disk

Inside-Out formation by infall and radially varying SFR efficiency

required to reproduce

observed SFR, gas and colour profiles (Scalelengths: RB#4 kpc, RK#2.6 kpc)

(Boissier  and  Prantzos  1999) Hou et al. 2000

and predict abundance gradients

Hou et al. 2000

(21)

Stasinska  et  al.  2012  

(22)

Stasinska  et  al.  2012  

The Ogygen profile in the Milky Way disk

What is its true value and shape? Does it flatten (steepen) in the inner or outer disk ?

It constrains combination of SFR(R) + infall(R) + radial inflows

(23)

Present day gradient : dlog(O/H)/dR ׽ -­ 0.045 dex/kpc

Metallicity profile of Milky Way disk

12 Gyr

7.5 Gyr 3 Gyr

Ideally,

combined observations of D and O

across the Galactic disk would offer

a clear picture of the radial variations

of SFR + infall (NP 1996)

(24)

   

     

     

Maciel et al. 2005

Evolution of the abundance profile

difficult  to  measure    

Large  uncertainties  in  both       abundance  profile      

(abundances,  distances)     and  age  determinations  

Fu et al. 2009 Hou et al. 2000

(25)

Theoretical evolution of abundance gradients (Pikkington  et  al.  2012)

(26)

Yuan et al. 2011

Observation of abundance profiles at high redshift in lensed galaxies

7KHDEXQGDQFHSURILOHVDSSHDUWREHPXFKVWHHSHUWKDQW\SLFDOSURILOHVRIORFDOGLVNV«

SDSSJ120601.69+514227.8 ´WKH&ORQHµ 

Jones et al. 2010

SP1149 Z~1.5

Z~2

Models  from  

Boissier  and  Prantzos  2000   Yuan et al. 2011

(27)

The value of the abundance gradient is crucial regarding the action of a galactic bar

A bar drives (metal-‐poor) gas from the outer disk inwards, and this radial inflow tends to reduce abundance gradients

The Milky Way bar ?

The importance of the Galactic bar for chemical evolution

is not well established yet

M 101 ² non barred Strong  gradient  

NGC 1365 ² barred Weak  gradient  

Dors and Copetti (2005)

(28)

Stasinska  et  al.  2012  

Jin  et  al.  2009  

A  small   abundance    

gradient   in  M31    

(29)

Jin  et  al.  2009  

Comparative evolution of the Milky Way and Andromeda

SFR efficiency = MW x 2

(30)

The abundance gradient decreases (in absolute value)

with galaxy mass/luminosity when expressed in dex/kpc but is ~constant when expressed in

dex/scalength or dex/R25

(Garnett  et  al.  1997´KRPRORJRXVHYROXWLRQµ  But: what drives such a homologuous evolution?

Garnett  et  al.  1997  

(31)

Tremonti  et  al.  2004  

Zahid  et  al.  2011  

Mass-­metallicity  relation  

It is shaped by

the complex interplay between infall, outflow and SFR in the mass-­dependent galactic potential wells

with a mass-­dependent scatter:

95% of stars within

~1 dex at logM=8

~0.5 dex at logM=10

(32)

Mass-­metallicity     relation  

Its  upper  part    (VC  >  100  km/s)   can  be  readily  explained  

by  more  efficient  gas  processing    

Z  =  y  ln(1/ı)  

in  more  massive  disks,   that  are  observed  to  have    

lower  gas  fractions  

ı

,     on  average  

(Boissier  and  Prantzos  2000)    

(33)

Its lower part (VC  <  100  km/s) concerns gas rich galaxies (IJ~1) and requires an effective  yield      

y

eff  

=  Z  /  ln(1/ı)  

decreasing  at  low  masses,  i.e.important outflows (Garnett  2002,  Tremonti  et  al.  2004)

or a galaxy mass-­dependent IMF (Koppen  et  al.  2007)

Mass-­metallicity  relation  

Tremonti  et  al.  2004   Garnett  2002  

(34)

Mannucci    et  al.  2010  

A  fundamental  relation   between    

Z,  M *  and  SFR  

(Mannucci  et  al.  2010)  

Also  valid  at  higher  redshifts  z  

(35)

Yates  et  al.  2012:  

   

Observational     and  theoretical   (semi-­analytical)    

investigation   of  the  evolution  

of  the     M-­Z  relation  

including     the  effects      of  the  SFR  

 

Unsuccessful   at  high  redshifts  

(36)

The more massive AND metallic a star is the more mass it loses in stellar wind

and less H (even He) is left at the SN explosion 61,,ĺ61,Eĺ61,F  

For a given IMF, the number ratios N(II)/N(Ibc) or N(Ic)/N(Ib)

will depend on

the limiting masses MIb, MIc, which are functions of metallicity

By determining N(II)/N(Ibc) or N(Ic)/N(Ib)

as a function of metallicity Z in galaxies one may determine MIb, MIc

and constrain stellar models

(Prantzos  and  Boissier  2003)  

SN rates vs metallicity

How to determine galaxian metallicity

for a large number of SN?

(37)

 

By  using  the   metallicity  vs    

luminosity   relation  

one  may  correlate     SN  ratios    

to  the   global    

galaxian  metallicity   (Prantzos  and  Boissier  2003)  

     

By  using  the  

metallicity  gradient     vs  luminosity  

relation  

one  may  correlate  SN   ratios    

to  the  

local  galaxian   metallicity  

(Boissier  and  Prantzos  2009)    

                         

(38)

There is indeed a correlation of N(Ibc)/N(II) with metallicity Z

either on a global level

(Prantzos  and  Boissier  2003)  

or on a local level (Boissier  and  Prantzos  2009)

(39)

This correlation constrains Single and binary models RIPDVVLYHVWDUHYROXWLRQ«

«DQGLPSRVHVPDVVOLPLWV MIb, MIc, as function of metallicity

to single star models

Boissier  and  Prantzos  2009  

(40)

Modjaj  et  al.  2011  

Leloudas  et  al.  2011  

Better:  Direct  measurements  of  O  abundance  at  the  SN  position    

But  it  will  take  time  !  

(41)

Conclusions

Measurements of oxygen abundances in various sites (stars and ISM) of galaxies are necessary,

to constrain models of galactic evolution

but, in most cases, this is done within a theoretical framework involving several other poorly known parameters.

In some cases (SN and GRB host galaxies) measurements

constrain directly pre-­SN and pre-­GRB models

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