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Figure 3 shows the evolution of the radial inflow velocity of the gas. | Figure \ref{fig_vr} shows the evolution of the radial inflow velocity of the gas. |
Initially it is zero, and in the absence of cooling stays that way except in the outer region (r>100 kpc), where the cluster is not in perfect hydrostaticequilibrium!. | Initially it is zero, and in the absence of cooling stays that way except in the outer region $r > 100$ kpc), where the cluster is not in perfect hydrostatic. |
. As radiative cooling proceeds, a slow but steady inflow develops. | As radiative cooling proceeds, a slow but steady inflow develops. |
This is slow (~20 km/s) and nearly constant over a large range in radius, from 20 kpc | This is slow $\sim 20$ km/s) and nearly constant over a large range in radius, from 20 kpc |
Large scale magnetic fields. with strengths LO 107 C. have been observed in spiral and disc galaxies. galaxy clusters. and high redshift) concdensations κος (Ixronborg1994:Han&Wielebinski2002) ane references therein). | Large scale magnetic fields, with strengths $10^{-7}$ $10^{-5}$ G, have been observed in spiral and disc galaxies, galaxy clusters and high redshift condensations [see \cite{K,HW} and references therein]. |
The most promising explanation for the large scale galactic fields has been the dynamo mechanism. with the required seeds coming either from local astrophysical processes. such as battery ellects. or from. primordial magnetogenesis for recent reviews. see (Crasso&Rubinstein2001:Widroyμα 2002)]. | The most promising explanation for the large scale galactic fields has been the dynamo mechanism, with the required seeds coming either from local astrophysical processes, such as battery effects, or from primordial magnetogenesis [for recent reviews, see \cite{GR,Wid}] ]. |
The linear evolution of large scale magnetic LickWu. and their implications for structure formation has been studied by several authors sce. c.g. (Ruzmaikina&Ruz-Jarrow&Subramanian1998:PsagasMaartens 2000). | The linear evolution of large scale magnetic fields and their implications for structure formation has been studied by several authors [see, e.g., \cite{RR,W,PE,EF,TB,BS,TM}] ]. |
Certain aspects of the mildly nonlinear. clustering can be analvsecl in spherical symmetry. but. this approximation inevitably breaks down as the collapse proceeds and. any initially small anisotropics take over. | Certain aspects of the mildly nonlinear clustering can be analysed in spherical symmetry, but this approximation inevitably breaks down as the collapse proceeds and any initially small anisotropies take over. |
When magnetic fields are involved. the need. to incorporate anisotropic effects is particularly important. as the fields. are. themselves generically anisotropic sources. | When magnetic fields are involved, the need to incorporate anisotropic effects is particularly important, as the fields are themselves generically anisotropic sources. |
Numerical simulations of magnetic Ποιά evolution in galaxy clusters suggest. that anisotropies lead to aclelitional amplification of the field. | Numerical simulations of magnetic field evolution in galaxy clusters suggest that anisotropies lead to additional amplification of the field. |
Vidal cllects during mergers. lor example. increase the magnitude of the magnetic field. (Rocttigerοἱal1999). | Tidal effects during mergers, for example, increase the magnitude of the magnetic field \cite{RSB}. |
Shear flows in galaxy clusters can also amplify the field bevond the limits of spherical compression (Dolagetal1999:Dolagetal 2002).. | Shear flows in galaxy clusters can also amplify the field beyond the limits of spherical compression \cite{DBL1,DBL2}. |
The latter simulations also suggest that the final magnetic configuration is effectively independent of the fields initial set up or of the presence of a cosmological constant. | The latter simulations also suggest that the final magnetic configuration is effectively independent of the field's initial set up or of the presence of a cosmological constant. |
Llere we use the Zeldovich approximation to. look analytically into the οσο of gravitational anisotropy on seed magnetic fields in the milelly nonlinear regime. | Here we use the Zel'dovich approximation to look analytically into the effect of gravitational anisotropy on seed magnetic fields in the mildly nonlinear regime. |
The anisotropic collapse is driven by cold. dark matter (CDAD. | The anisotropic collapse is driven by cold dark matter (CDM). |
X previous analysis (Zel'dovichetal1983).. where the matter is purely barvonic. concluded: that the seeds must. be nesligiblv small in order to avoic the erowth of magnetic-. fields to levels which prevent pancake formation. | A previous analysis \cite{ZRS}, where the matter is purely baryonic, concluded that the seeds must be negligibly small in order to avoid the growth of magnetic fields to levels which prevent pancake formation. |
By contrast. this strict. constraint is avolded when CDM dominates. since the magnetic field couples only eravitationally with the CDM. | By contrast, this strict constraint is avoided when CDM dominates, since the magnetic field couples only gravitationally with the CDM. |
The field is frozen into the barvon fluid. ancl barvons are dragged by the CDM eravitational field. | The field is frozen into the baryon fluid, and baryons are dragged by the CDM gravitational field. |
Phe barvons are in a cillerent state of motion to the CDM and. unlike the ΔΙ. they feel the magnetic backreaction. | The baryons are in a different state of motion to the CDM and, unlike the CDM, they feel the magnetic backreaction. |
As a first approximation. however. we will ignore the magnetic backreaction ancl the relative motion of the barvons. elfectively considering a single Iuid. | As a first approximation, however, we will ignore the magnetic backreaction and the relative motion of the baryons, effectively considering a single fluid. |
This approximation also maintains the acceleration-free aud irrotational nature of the motion. which are Κον ingredients ofthe Zeldovich approximation. | This approximation also maintains the acceleration-free and irrotational nature of the motion, which are key ingredients of the Zel'dovich approximation. |
Our approach may be seen as a qualitative starting point for a more detailed analysis. | Our approach may be seen as a qualitative starting point for a more detailed analysis. |
We begin in Section 2 with a dynamical svstem description of the Zel'dovich approximation. which directly shows that. in a generic collapse. pancakes are the | We begin in Section 2 with a dynamical system description of the Zel'dovich approximation, which directly shows that, in a generic collapse, pancakes are the |
Therefore aLobjects were selected that have classificr values below 0.35. | Therefore all objects were selected that have classifier values below 0.35. |
Down to a Veo,=22.0 aud 23.0 mag we found 8723 aud 1775 galaxies respectively. | Down to a $V_{\rm tot} = 22.0$ and 23.0 mag we found 873 and 1775 galaxies respectively. |
The properties ofa] galaxies xiehter than 22.0 mae are compiled ina catalog. see Appendix A. This cutoff in the final sample was chosen for several reasons: the fudiug completcucss stars to drop. he scatter i the classifier values increases significanth. | The properties of all galaxies brighter than 22.0 mag are compiled in a catalog, see Appendix A. This cutoff in the final sample was chosen for several reasons: the finding completeness starts to drop, the scatter in the classifier values increases significantly. |
Also. at this magnitude we wotld not πα aly dEs that follow the surface brightuess magnitude relaion of Local Group or Virgo dwarf spheroidals (6.8. ]xornendy 1985.. DiusecliOO 199 1)) due to the luüt iu surface brightuess; | Also, at this magnitude we would not find any dEs that follow the surface brightness – magnitude relation of Local Group or Virgo dwarf spheroidals (e.g. Kormendy \cite{korm}, , Binggeli \cite{bing}) ) due to the limit in surface brightness. |
See Sect. | See Sect. |
7 for a discussioji of the COLLeteuess of dwarf galaxies iu the Fornax disace. | 7 for a discussion of the completeness of dwarf galaxies in the Fornax distance. |
Iu the backeround fields B3 and Bl we found with the sale selection criteria 668 and 1022 down to V=22.0 nag respectively. | In the background fields B3 and B4 we found with the same selection criteria 668 and 1022 down to $V = 22.0$ mag respectively. |
However. one has to be careful when colmparing these resuts with those of the other CCD fields. | However, one has to be careful when comparing these results with those of the other CCD fields. |
The pixel size is about 3 times larger than in the first run leading to a 9 times higher area covere by each xel. | The pixel size is about 3 times larger than in the first run leading to a 9 times higher area covered by each pixel. |
We simulated tus resolution for two fields of our first run by biuniug 3«23 pixel aud run SExtractor again. | We simulated this resolution for two fields of our first run by binning $3 \times 3$ pixel and run SExtractor again. |
Ou tfie one hand. some galaxies have been classified as voit sources due to their small augular sizes below the "new resolution. | On the one hand, some galaxies have been classified as point sources due to their small angular sizes below the “new” resolution. |
Ou the other laud. some new. "ealaxies? inve »cen gained due to the overlap of objects very close iu he high resolution iniage. | On the other hand, some new “galaxies” have been gained due to the overlap of objects very close in the high resolution image. |
Down to our magnitude linuit of y=22.0 niae loss and eain of galaxies are nearly balauced aud i the order of of the total galaxy counts. | Down to our magnitude limit of $V = 22.0$ mag loss and gain of galaxies are nearly balanced and in the order of of the total galaxy counts. |
Iu combination with the morphological appearance of he observed galaxies. the analysis of t101 SB profiles provides a reliable tool to classify them as cluster dwarf galaxies or backeround galaxies (c.g. Sandage Bineecli 1981)). | In combination with the morphological appearance of the observed galaxies, the analysis of their SB profiles provides a reliable tool to classify them as cluster dwarf galaxies or background galaxies (e.g. Sandage Binggeli \cite{sand}) ). |
Iu the surface brightucss versus magnitude diagrams (ji-V. diagram) the dEs follow a well defined sequence. | In the surface brightness versus magnitude diagram $\mu$ $V$ diagram) the dEs follow a well defined sequence. |
A erowth curve analysis for each ealaxv las been iade with increasing clliptical apertures using the position. ellipticity. and position anele of the SExtractor photometry results. | A growth curve analysis for each galaxy has been made with increasing elliptical apertures using the position, ellipticity, and position angle of the SExtractor photometry results. |
As local backeround the SExtractor value of the interpolated skv map for cach iuciviual eaaxv was taken (see Sect. | As local background the SExtractor value of the interpolated sky map for each individual galaxy was taken (see Sect. |
1). | 4). |
Iu. some cases overlapping Sars have been removed before the analysis x interpolatiie the galaxy surface brightuess profile from a unaffected rine outsie the stellar profile. | In some cases overlapping stars have been removed before the analysis by interpolating the galaxy surface brightness profile from a unaffected ring outside the stellar profile. |
Iu otLOY Cases, he region of au overlayplug star was masked out durius he fitΠιο process of fje surface briehtness profile. | In other cases, the region of an overlapping star was masked out during the fitting process of the surface brightness profile. |
Two nodeLindependent parameters have been deteruined for he eaaxies: the effective semiauajor axis e, (inajor axis of the ellipse that coutaius half of the toal light) and thenean effective surtace brigltuecss within the effecIVC seninajor axis. Glo)=Ves| 5dog(aug)| 2.5-log(27(1€)). with e=1bia (b= πο]ΙΟ axis. &= senmiüdnuajor axis). | Two model-independent parameters have been determined for the galaxies: the effective semi-major axis $a_{\rm eff}$ (major axis of the ellipse that contains half of the total light) and themean effective surface brightness within the effective semi-major axis, $\langle \mu_{\rm eff} \rangle = V_{\rm tot}
+ 5\cdot$ $(a_{\rm eff}) + 2.5\cdot$ $(2\pi(1 - \epsilon))$, with $\epsilon =
1 - b/a$ $b =$ semi-minor axis, $a =$ semi-major axis). |
These paraimcters. as well as the size of the full luajor axis Dog at the isophote of V—26 mag 7, are given in the photometric catalog (Appendix A). | These parameters, as well as the size of the full major axis $D_{26}$ at the isophote of $V = 26$ mag $^{-2}$, are given in the photometric catalog (Appendix A). |
The fia-Veor is shown iu Fig. | The $\mu_{\rm eff}$ $V_{\rm tot}$ is shown in Fig. |
3 (iamiddle panel). | 3 (middle panel). |
Qualitatively. this plot is comparable to the ρω lua diagersnn in the upper panel. | Qualitatively, this plot is comparable to the $\mu_{\rm peak}$ $V_{\rm tot}$ diagram in the upper panel. |
The sequence of dwarf galaxies is clearly separated from the location of ckerouinm ealaxies. | The sequence of dwarf galaxies is clearly separated from the location of background galaxies. |
However. tιο nucleated dEs. that are uidden iu the fi plot among the backerouud galaxies. all iu he range of the dE sequence when measuring Hag. | However, the nucleated dEs, that are hidden in the $\mu_{\rm peak}$ plot among the background galaxies, fall in the range of the dE sequence when measuring $\mu_{\rm eff}$. |
In hoh diagrams. there are some galaxies located )dow the bulk of backgrouid galaxies. falling iu the range of dwarfcllipticals. | In both diagrams, there are some galaxies located below the bulk of background galaxies, falling in the range of dwarfellipticals. |
Eacd individual galaxy in this region has oen Individually iispected. | Each individual galaxy in this region has been individually inspected. |
Most of them are ckerouinm spirals or galaxies that have close neighbours which distirb the correct pog calculation. | Most of them are background spirals or galaxies that have close neighbours which disturb the correct $\mu_{\rm eff}$ calculation. |
Furthermore. radial velocity nieasureimienuts have shown that nearly all ealaxies at the bright ji limi the dE sequence are judecd ckerouinm objects (see crosses in the middle panel of Fig. | Furthermore, radial velocity measurements have shown that nearly all galaxies at the bright $\mu$ limit of the dE sequence are indeed background objects (see crosses in the middle panel of Fig. |
3. | 3. |
The determination of he πο central surface xiehtness pp is critical for objects wih small angular diameters. | The determination of the true central surface brightness $\mu_0$ is critical for objects with small angular diameters. |
The centrally peaked light distributions are shared by seeius. which eads to a dinuuing of p. | The centrally peaked light distributions are blurred by seeing, which leads to a dimming of $\mu_0$. |
If he apparent core radius. where the surface intensity has decreased by a factor of 2 from its central ayparent value. is siualler than 20..(0=FWIIAL2.351). the SB profile cannot be deconvolved from theseciug xofile (Scluwcizcer 1981.. IXoxiueudy. 1985)). | If the apparent core radius, where the surface intensity has decreased by a factor of 2 from its central apparent value, is smaller than $2\sigma_\ast (\sigma_\ast = FWHM/2.354)$, the SB profile cannot be deconvolved from theseeing profile (Schweizer \cite{schw}, Kormendy \cite{korm}) ). |
In our observatiois the average secing dispersion is about σ=15 | In our observations the average seeing dispersion is about $\sigma_\ast =
0\farcs45$. |
Iu the following. the calculations are restricted subsample of galaxies. whose apparent core radii re, are larger than 079. | In the following, the calculations are restricted to a subsample of galaxies, whose apparent core radii $r_{\rm c,app}$ are larger than $0\farcs9$. |
The correcions given bv IKonueudy (1985)) were applied to derive frue core radii and true central surface brightuesses ουν. | The corrections given by Kormendy \cite{korm}) ) were applied to derive true core radii and true central surface brightnesses $\mu_{\rm 0,cor}$. |
Note that for objects with reap=079 the correction or the true ceutral surface brightuess is of the order of 2 uaenitudes. aud the true core radius is about a third of the apparent one. | Note that for objects with $r_{\rm c,app} = 0\farcs9$ the correction for the true central surface brightness is of the order of 2 magnitudes, and the true core radius is about a third of the apparent one. |
At the distance of the Fornax cluster V9 corresponds to about SO pc. | At the distance of the Fornax cluster $0\farcs9$ corresponds to about 80 pc. |
The Local Group dSpis. for example. lave core radii between 150 and 600 pe (Caldwell etal. 1992.. | The Local Group dSphs, for example, have core radii between 150 and 600 pc (Caldwell etal. \cite{cald92}, |
Mateo et al. 19933). | Mateo et al. \cite{mate}) ). |
Thus. the apparent surface brighnesses of dEs in the Fornax οuster shotId be nearly ideutical with their true ceutral srace brigitnesses. Whereas compact dwarts (ike M32) aud backeround galaxies are severely nuderestimated iu ticr nuesred posae | Thus, the apparent surface brightnesses of dEs in the Fornax cluster should be nearly identical with their true central surface brightnesses, whereas compact dwarfs (like M32) and background galaxies are severely underestimated in their measured $\mu_{\rm peak}$. |
The core radius of M32. for example. ix abot 500 times smaller than that of Local Caoup and Virgo dEs (7 1-2 pe. I&oxnoeudxy 1985)). | The core radius of M32, for example, is about 500 times smaller than that of Local Group and Virgo dEs $\simeq 1$ $2$ pc, Kormendy\cite{korm}) ). |
Fimre 3 (lower panel) shows the ουντοι for all ealaxies wit lcorrecions less than 1.5 mag. | Figure 3 (lower panel) shows the $\mu_{\rm 0,cor}$ $V_{\rm tot}$ for all galaxies with corrections less than 1.5 mag. |
The differeut svinbol sizes divide the sample iu degrees of resolution of the core. as given by the ratio reap/0 which can directly be rauslated into the correction in maguitudes Ape Ape<0.2 mag corresponds to reappsFs> 5. Apex0.5 mag to rea,fo.> δι and Apο ag to reappsts> 2. | The different symbol sizes divide the sample in degrees of resolution of the core, as given by the ratio $r_{\rm c,app}/\sigma_\ast$ , which can directly be translated into the correction in magnitudes $\Delta\mu$ : $\Delta\mu < 0.2$ mag corresponds to $r_{\rm c,app}/\sigma_\ast > 5$ , $\Delta\mu < 0.5$ mag to $r_{\rm c,app}/\sigma_\ast
> 3$ , and $\Delta\mu < 1.5$ mag to $r_{\rm c,app}/\sigma_\ast > 2$ . |
All nou-uneleated Fornax cluster dwarts are clearly separatedfroin the bulk of backgroundoO galaxies aud fit the | All non-nucleated Fornax cluster dwarfs are clearly separatedfrom the bulk of background galaxies and fit the |
Tweziéetal.(2008.hereafter108) further exteuded this elobal aualvsis of SDSS data by developing a photometric metallicity. estimator aud by utilizing a huge proper motion catalog based on SDSS aud Palomar Observatory Sky Survey data (Munetal.2001). | \citet[hereafter I08]{Ivezic2008} further extended this global analysis of SDSS data by developing a photometric metallicity estimator and by utilizing a large proper motion catalog based on SDSS and Palomar Observatory Sky Survey data \citep{Munn2004}. |
. I08 studied the dependence of the ietallicity. |Fe/TII] and rotatioua velocity. V. of disk stars on the distance from the Galactic plane aud detected eradicuts of both quantities over the distance ranging frou several hundred parsecs to several kiloparsecs. | I08 studied the dependence of the metallicity, [Fe/H] and rotational velocity, $V_{\phi}$, of disk stars on the distance from the Galactic plane and detected gradients of both quantities over the distance ranging from several hundred parsecs to several kiloparsecs. |
Such gradieuts would be expectcc in a thin/thick disk decomposition where the thick disk stars are a separate population defined bv a bulk rotational velocity lag aud a lower metallicity conmipare to those of the thin disk. | Such gradients would be expected in a thin/thick disk decomposition where the thick disk stars are a separate population defined by a bulk rotational velocity lag and a lower metallicity compared to those of the thin disk. |
However. such a model xvouk also predict a correlation between the metallicity arc the velocity lag. which is strouglv excluded. (76 level) by the IOS analysis (sec Figure 17. I08). | However, such a model would also predict a correlation between the metallicity and the velocity lag, which is strongly excluded $\sim$$7\sigma$ level) by the I08 analysis (see Figure 17, I08). |
Tn this work we turn to a more sophisticated Galactic description anu N body model o characterize stars withiu the SDSS volue aud solve this puzzle. | In this work we turn to a more sophisticated Galactic description — an $N$ –body model — to characterize stars within the SDSS volume and solve this puzzle. |
Over the past few decades. No body simulations have been used to provide supporting evidence for three distinct theories of thick disk formation: violent relaxation (Jones&Wyse1983).. substructure disruption (StatlerL988).. and heating bv satellites (Quinnctal.1993). | Over the past few decades, $N$ –body simulations have been used to provide supporting evidence for three distinct theories of thick disk formation: violent relaxation \citep{JonesWyse1983}, substructure disruption \citep{Statler1988}, and heating by satellites \citep{Quinn1993}. |
. Several works have receutly redressed these ideas. | Several works have recently redressed these ideas. |
Brooketal.(2001) aud Bouruaudotal.(2009) formed a thick disk in situ at high redshift duriug eas-rich mereecrs. where star formation is trigecred bv the rapid accretion of eas: this result is consistent with the thick disk forming through violent relaxation of the ealactic poteutial. | \citet{Brook2004} and \citet{Bournaud2009} formed a thick disk in situ at high redshift during gas-rich mergers, where star formation is triggered by the rapid accretion of gas; this result is consistent with the thick disk forming through violent relaxation of the galactic potential. |
Tn coutrast. Wazautzicdisctal.(2008).. Villalobos&Ποια (2008).. and Villalobosetal.(2010) investigated SHstructure disruption Wusing a cosmological voderived BHcllite accretion historv to perturb a Milky Wie-Hke disk: subhlialo-disk eucointers increased the scide height of this disk at all raclii effectively. foriing a tick disk. | In contrast, \citet{Kazantzidis2008}, \citet{Villalobos2008}, and \citet{Villalobos2010} investigated substructure disruption by using a cosmologically derived satellite accretion history to perturb a Milky Way-like disk; subhalo-disk encounters increased the scale height of this disk at all radii effectively forming a thick disk. |
Fiwally. Abadietal.{20103) showed that ]κ. tidally stiIppine/acercting satelites. the majority of the odest stars in the thick disk could have formed externally rather than in situ. | Finally, \citet{Abadi2003} showed that by tidally stripping/accreting satellites, the majority of the oldest stars in the thick disk could have formed externally rather than in situ. |
Tn this work. we study a new method of formation: radial iieration. | In this work, we study a new method of formation: radial migration. |
Radial migration due to scattering frou transicut spirals was first described by Sellwood&Diu-uev (2002). | Radial migration due to scattering from transient spirals was first described by \citet{Sellwood2002}. |
. In this model energy aud augular 1ioiicutui changes occur from mteractions with trausicut spiral armis. which move stars at the corotation resonance inward or outward iu radius while preserviusg their nearly-circular orbits. | In this model energy and angular momentum changes occur from interactions with transient spiral arms, which move stars at the corotation resonance inward or outward in radius while preserving their nearly-circular orbits. |
Roskaretal.(2008a.b.BRüsabhereafter) studied this phenomenon i N-body | Sunooth Particle Uvdrodvuanmuic (SPILT) simulations of disk formation. and showed that nüeratious are possible ou short timescales. | \citet[][R08ab hereafter]{Roskar2008, Roskar2008a} studied this phenomenon in N-body + Smooth Particle Hydrodynamic (SPH) simulations of disk formation, and showed that migrations are possible on short timescales. |
They explored the implications of radial mixing for stellar populations for a variety of stellar svstenis. iucludiug the solar neighborhood. | They explored the implications of radial mixing for stellar populations for a variety of stellar systems, including the solar neighborhood. |
ere we exteud thei work bv highliehtiug the vertical evolution that occurs as a result of migration. | Here we extend their work by highlighting the vertical evolution that occurs as a result of migration. |
We note that iu this paper. we are not testing the validity of the other models of formation. | We note that in this paper, we are not testing the validity of the other models of formation. |
Πονονο. recently. Salesetal.(2009) proposed uxing the eccentricity of orbits of stars iu the thick disk to coustrain the thick disks formation mechanism: they prescuted the eccentricity distuibutions that result from four body simulations: Abadietal.(2003)... (2008).. ROsb. aud Brooketal.(2001).. | However, recently, \citet{Sales2009} proposed using the eccentricity of orbits of stars in the thick disk to constrain the thick disk's formation mechanism; they presented the eccentricity distributions that result from four N-body simulations: \citet{Abadi2003}, , R08b, and \citet{Brook2004}. |
They found that the distributious that result from heating. radial nügration and mergers all had a stroug peal at low eccentricity (6~0.2.—0.3). while the distribution that results from accretion is centered at ligher orbital eccentries (<ée>~0.5). | They found that the distributions that result from heating, radial migration and mergers all had a strong peak at low eccentricity $\epsilon \sim 0.2 - 0.3$ ), while the distribution that results from accretion is centered at higher orbital eccentries $<$$\epsilon$$> \sim 0.5$ ). |
Building on this. Wilsonetal.(2010). studied the ecceutricity of orbits of stars in the thick disk observed in the Radial Velocity Experiment (RAVE) (Steinmetzetal.2006) and found these results to be incousistent with expectations for the pure accretion simulation. | Building on this, \citet{Wilson2010} studied the eccentricity of orbits of stars in the thick disk observed in the Radial Velocity Experiment (RAVE) \citep{Steinmetz2006} and found these results to be inconsistent with expectations for the pure accretion simulation. |
Ruchtietal.(2010) also leveraged. a ineasurements from RAVE to conclude that the a chhancement of the metal-poor thick disk iauplies that direct aceretion of stars frou dwarf galaxies did not plav a imajor role in the formation of the thick disk. | \citet{Ruchti2010} also leveraged $\alpha$ measurements from RAVE to conclude that the $\alpha$ enhancement of the metal-poor thick disk implies that direct accretion of stars from dwarf galaxies did not play a major role in the formation of the thick disk. |
Using SDSS DR. Dierickxetal.(2010). showed that the eccenutricitv of orbits of stars iu the thick disk implies the thick disk is unlikely to be fully populated by radially nuerated stars. | Using SDSS DR7, \citet{Dierickx2010} showed that the eccentricity of orbits of stars in the thick disk implies the thick disk is unlikely to be fully populated by radially migrated stars. |
We note that we cannot exclude that sole fraction of the thick disk is a fossil of a past more violent history. nor can this scenario explain thick disks in all galaxies. | We note that we cannot exclude that some fraction of the thick disk is a fossil of a past more violent history, nor can this scenario explain thick disks in all galaxies. |
However. in what follows. we show that a laree fraction of the stars in the thick disk could have formed im situ and arrived at their present location via radial iuieration. | However, in what follows, we show that a large fraction of the stars in the thick disk could have formed in situ and arrived at their present location via radial migration. |
The outline of this paper is as follows: in refssdnaul. we present two simulations. one with substantial migration aud the other with relatively little nueration. | The outline of this paper is as follows: in \\ref{s:simul}, we present two simulations, one with substantial migration and the other with relatively little migration. |
When we compare these two sinulations we can show that ποτάο can build a thick disk as first conceived by Calmore Reid (1983): a component witli a scale-height larger than that of tle thin disk. | When we compare these two simulations we can show that migration can build a thick disk as first conceived by Gilmore Reid (1983): a component with a scale-height larger than that of the thin disk. |
In refs:obs we qualitatively compare the \lilsy Way-like simulation (with iuieration) with the SDSS observations to show that they match cach other sufficiently well to pursue further comparison. | In \\ref{s:obs} we qualitatively compare the Milky Way-like simulation (with migration) with the SDSS observations to show that they match each other sufficiently well to pursue further comparison. |
Iu refs:solar we present a detailed comparison between the sinulation aud the local SDSS volume focusing on the reason for the lack of correlation between V,, aud [Fe/T]: in Appendix À we reconsider recent observational claims coucerning the lack of correlation between V, and |Fe/TI]. | In \\ref{s:solar} we present a detailed comparison between the simulation and the local SDSS volume focusing on the reason for the lack of correlation between $V_{\phi}$ and [Fe/H]; in Appendix \ref{s:spagna} we reconsider recent observational claims concerning the lack of correlation between $V_{\phi}$ and [Fe/H]. |
Iu refs:obsccompweuscthesimulationasaprocvy fortheALil ky egtoshi schocirich | In \\ref{s:obs_decomp} we use the simulation as a proxy for the Milky Way to show that classifying stars as members of the thin or thick disk by either velocity or metallicity leads to an apparent separation in the other property as observed. |
wecomparcourresultstorecenttheoreticalworktha yticstoincestiqgatehouwthesolarne: ghborhoodcouldhacvebeensha | | In \\ref{s:schoenrich} we compare our results to recent theoretical work that used semi-analytics to investigate how the solar neighborhood could have been shaped by radial migration and chemical evolution effects. |
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