ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
Also, SiC feature at around 11.3 wm can be seen in carbon stars.	Also, SiC feature at around 11.3 $\mu$ m can be seen in carbon stars.
Interestingly, '13 um feature’ is seen in M-type and/or S-type giants.	Interestingly, so-called '13 $\mu$ m feature' is seen in M-type and/or S-type giants.
This feature is probably due to aluminum oxides (e.g., Posch et al. 1999)).	This feature is probably due to aluminum oxides (e.g., Posch et al. \cite{posch1999}) ).
Sloan et al. (2003b))	Sloan et al. \cite{sloan2003b}) )
suggested that this feature tends to be stronger in systems with lower infrared excesses and thus lower mass-loss rates (e.g., Onaka et al. 1989;;	suggested that this feature tends to be stronger in systems with lower infrared excesses and thus lower mass-loss rates (e.g., Onaka et al. \cite{onaka1989};
Kozasa Sogawa 1997)).	Kozasa Sogawa \cite{kozasa1997}) ).
Interestingly, Sloan et al. (1996))	Interestingly, Sloan et al. \cite{sloan1996}) )
reported that the 13 um feature is preferentially detected in semi-regular or irregular variables.	reported that the 13 $\mu$ m feature is preferentially detected in semi-regular or irregular variables.
Based on these available data of galactic counterparts, we suggest that low mass-loss rate M-type giants,	Based on these available data of galactic counterparts, we suggest that low mass-loss rate M-type giants,
post-starburst galaxies can be further divided iuto k \| a anda \| k type galaxies. (	post-starburst galaxies can be further divided into k + a and a + k type galaxies. (
Fraux 1993: Dressler et al.	Franx 1993; Dressler et al.
1999: Pogeianti ct al.	1999; Poggianti et al.
1999).	1999).
Northern. southern aud casteru chuups are mostly populated by k type galaxies aud a fow k \| a and ai \| k galaxies are also seen im these clumps.	Northern, southern and eastern clumps are mostly populated by k type galaxies and a few k + a and a + k galaxies are also seen in these clumps.
The classification as post-starburst galaxies of two out of ten cluster 1iucnibers may be uncertain due to the possible alteration of the IL; equivalent width by the tellure correction.	The classification as post-starburst galaxies of two out of ten cluster members may be uncertain due to the possible alteration of the $_{\delta}$ equivalent width by the telluric correction.
Alost passive galaxies have red colours (I.ἷνςs= 5) aud most of them populate the cluster red sequence (CRS) in the colouranaguitude diagram (CAID).	Most passive galaxies have red colours $R-K_s \simeq 5$ ) and most of them populate the cluster red sequence (CRS) in the colour-magnitude diagram (CMD).
In Fig. 9..	In Fig. \ref{rcs},
we show the distribution of cluster 1iemibers that have eround based LRIS and Sofl photometry iu two CAIDs.	we show the distribution of cluster members that have ground based LRIS and SofI photometry in two CMDs.
Open circles are passive galaxies aud solid triangles are star-forming galaxies.	Open circles are passive galaxies and solid triangles are star-forming galaxies.
The solid lines correspond to linear fits (ROAL=0.027\|5.1E in upper panel aud PofF=0.022F\|3.39 in lower panel) to the data within the regions defined in the caption of Fie. 9..	The solid lines correspond to linear fits $R-K_\mathrm{s} = -0.027 \ K_\mathrm{s} + 5.44$ in upper panel and $V-I = -0.022 \ I + 3.39$ in lower panel) to the data within the regions defined in the caption of Fig. \ref{rcs}.
Dashed lues indicate the 1-0 level above and below the fits.	Dashed lines indicate the $\sigma$ level above and below the fits.
If we compare the CRS of the RA. ovs dv, diagram to the models computed by Kodama Arimoto (1997). we infer that the colours of the passive galaxies in the CRS are cousisteut with an carly formation epoch 2).	If we compare the CRS of the $R-K_\mathrm{s}$ vs $K_\mathrm{s}$ diagram to the models computed by Kodama Arimoto (1997), we infer that the colours of the passive galaxies in the CRS are consistent with an early formation epoch $z \simgreat 2$ ).
However. we note the presence of a tail of fainter (ht.κ18.5) passive galaxies extending to ιο (L0XROON.X5.7) colours. off the CRS.	However, we note the presence of a tail of fainter $K_s \simgreat 18.5$ ) passive galaxies extending to bluer $4.0 \simless R-K_s \simless 5.7$ ) colours, off the CRS.
Passive galaxies are also observed outside the hieli-deusitv peaks. however the outskirts of aare dominated by star-forming galaxies.	Passive galaxies are also observed outside the high-density peaks, however the outskirts of are dominated by star-forming galaxies.
The segregation between star-forming and passive ealaxies in lis in agreement with a picture in which the star formation activity in galaxies 1s suppressed diving the iufall of these ealaxies into the highest density zoucs.	The segregation between star-forming and passive galaxies in is in agreement with a picture in which the star formation activity in galaxies is suppressed during the infall of these galaxies into the highest density zones.
An analysis of tle physical properties of the star-forming galaxy population in bbased on the combination of VLT spectroscopic aud HST/ACS imagine data. together with implications for cluster galaxy evolution. will be preseuted iu Honieier et al.	An analysis of the physical properties of the star-forming galaxy population in based on the combination of VLT spectroscopic and HST/ACS imaging data, together with implications for cluster galaxy evolution, will be presented in Homeier et al.
2001.	2004.
In Fie. 10.	In Fig. \ref{colcol},
we show the distribution of cluster members m two coloum-colour diagrams.	we show the distribution of cluster members in two colour-colour diagrams.
The sviibols are the same as those in Fie. 9..	The symbols are the same as those in Fig. \ref{rcs}.
For reference. we also slow the tracks as a function of redshift for au elliptical (Ell) and three late-type (She. Scd and Br) galaxies using the teiiplate library of Coleman. Wu. Weedman (1980). whose spectral cuerey distributious (SEDs) have becu extended to. the rear-IR and far-UV using Druzual Charlot (1993. with recent updates) models;	For reference, we also show the tracks as a function of redshift for an elliptical (Ell) and three late-type (Sbc, Scd and Irr) galaxies using the template library of Coleman, Wu, Weedman (1980), whose spectral energy distributions (SEDs) have been extended to the near-IR and far-UV using Bruzual Charlot (1993, with recent updates) models.
The tracks iuclude zeorrectious oulv.	The tracks include k-corrections only.
The expected colours for the differeut galaxy types at the cluster redshift are indicatec with stars.	The expected colours for the different galaxy types at the cluster redshift are indicated with stars.
Passive galaxies are located in the upper right part of thi colour-colour diagrams (see Fig. 10)).	Passive galaxies are located in the upper right part of both colour-colour diagrams (see Fig. \ref{colcol}) ).
The observe offset between the locus of these galaxies iux their expected colours from the models is consistent with passive evolution of earbv-tvpe galaxies formed at 2~1.9.	The observed offset between the locus of these galaxies and their expected colours from the models is consistent with passive evolution of early-type galaxies formed at $z \sim 1.9$.
Finally. we note that the brightest cluster ΠΟΡΟΙ of the southern clump (ID=387) is redder than the brightes cluster ποσο of the northern chuup (ID-701).	Finally, we note that the brightest cluster member of the southern clump (ID=387) is redder than the brightest cluster member of the northern clump (ID=701).
The differences in the colours are A(RA.)=OdT ane AVI)=(5 with crrors of 6(RWoy=0015 and οΕΠ=0.012 respectively.	The differences in the colours are $\Delta(R-K_\mathrm{s})=0.17$ and $\Delta(V-I)=0.15$ with errors of $\delta(R-K_\mathrm{s})=0.015$ and $\delta(V-I)=0.012$ respectively.
These differences in colour ect bigeer if we compare ID=3a7 with the pair of xieht galaxies (ID=1166 aud ID=1167) located in the core of the northern chump.	These differences in colour get bigger if we compare ID=387 with the pair of bright galaxies (ID=1466 and ID=1467) located in the core of the northern clump.
Calaxy ID=3s87 turus ou o he 0.25 magnitudes redder in RoA, and ~0.2 naenitudes redder in V.{ than both galaxies at the core of the northern subcluster.	Galaxy ID=387 turns out to be $\sim 0.25$ magnitudes redder in $R-K_\mathrm{s}$ and $\sim 0.24$ magnitudes redder in $V-I$ than both galaxies at the core of the northern subcluster.
Assiuuiug that these ealaxies ive smaular metallicities. such a cifference iu colour may )o an indication of a difference in stellar ages.	Assuming that these galaxies have similar metallicities, such a difference in colour may be an indication of a difference in stellar ages.
Towever. nore accurate piomoetry. sucli as. e.g...o ACS photometry.	However, more accurate photometry, such as, e.g., ACS photometry,
core.	core.
For these models a partial reflection at the bottom of the helium-rich. envelope reduces the amplitude and. hence radiative damping in the core.	For these models a partial reflection at the bottom of the helium-rich envelope reduces the amplitude and hence radiative damping in the core.
Since only few e-modes are excited and the evolutionary speed is fast. these g-mocdes are not very important observationalls.	Since only few g-modes are excited and the evolutionary speed is fast, these g-modes are not very important observationally.
Our original goals were to understand oscillations observed in the helium-rich. subdwarl 127116. ancl in. the DOGLITIG variables.	Our original goals were to understand oscillations observed in the helium-rich subdwarf $-14^{\circ}116$, and in the PG1716 variables.
We therefore carried out a broad review of Fe-bump driven pulsational instability in low-massstars’.	We therefore carried out a broad review of Fe-bump driven pulsational instability in low-mass.
ὃν dist considering radial modes. Νο have demonstrated: the essential rolle of chemical composition such that instability increases with the contrast. between the iron-bump. opacity and. other opacity sources.	By first considering radial modes, we have demonstrated the essential rôlle of chemical composition such that instability increases with the contrast between the iron-bump opacity and other opacity sources.
This increased Contrast may be achieved either. hy increasing the iron abundance. confirming earlier work by Charpinetetal.(2001). or by reducing the hyelrogen abundance (Jelfery&SaloLOOS).	This increased contrast may be achieved either by increasing the iron abundance, confirming earlier work by \citet{Cha01} or by reducing the hydrogen abundance \citep{Jef98}.
. At least one of these is necessary to excite oscillations in all of the Fe-bump pulsators discovered to date.	At least one of these is necessary to excite oscillations in all of the Fe-bump pulsators discovered to date.
We have further demonstrated that the bluc-edge for racial instability is allectecl by the mean molecular weight in the stellar envelope. so that increasing the iron abundance alone provides a bluer instability region than increasing the iron abundance in concert with all elements heavier than helium.	We have further demonstrated that the blue-edge for radial instability is affected by the mean molecular weight in the stellar envelope, so that increasing the iron abundance alone provides a bluer instability region than increasing the iron abundance in concert with all elements heavier than helium.
The former is required to explain the locus of the EC14026 instability region. making the general assumption that the properties of non-racial pamocdes moces are closely linked to the corresponding radial mode of the same racial orcler.	The former is required to explain the locus of the EC14026 instability region, making the general assumption that the properties of non-radial p-modes modes are closely linked to the corresponding radial mode of the same radial order.
Furthermore. the blue-edge also depends on the radial order of the oscillations. so that higher-order modes may be found in hotter stars.	Furthermore, the blue-edge also depends on the radial order of the oscillations, so that higher-order modes may be found in hotter stars.
By comparing theoretical and observed periods for EC14026 stars. we have shown that low-order or funcamental mode racial oscillations are only likely to be seen in the coolest EC14026 stars and that the hottest stars must oscillate in higher-order moces.	By comparing theoretical and observed periods for EC14026 stars, we have shown that low-order or fundamental mode radial oscillations are only likely to be seen in the coolest EC14026 stars and that the hottest stars must oscillate in higher-order modes.
However. we have been unable to explain the oscillations in the helium-rich star 4116: the observed. periods are simply too long compared with the L/AJ ratio obtained [rom spectroscopy.	However, we have been unable to explain the oscillations in the helium-rich star $-14^{\circ}116$; the observed periods are simply too long compared with the $L/M$ ratio obtained from spectroscopy.
A lower surface gravity would. definitely help.	A lower surface gravity would definitely help.
Considering non-radial pulsations. we have focused on e-mocde instability and in particular on instability in modes with spherical degree /«4.	Considering non-radial pulsations, we have focused on g-mode instability and in particular on instability in modes with spherical degree $l<4$.
We have discovered a small g-mode instability island on the blue horizontal branch which does not require iron enhancement or hydrogen depletion.	We have discovered a small g-mode instability island on the blue horizontal branch which does not require iron enhancement or hydrogen depletion.
Llowever. since it requires Z=0.02 and a partial rellection of he wave at the interface between the hydrogen-rich envelope and the helium core. it will take some observational cllort o verily whether there are any. real pulsating horizontal-anch stars which correspond with these models.	However, since it requires $Z=0.02$ and a partial reflection of the wave at the interface between the hydrogen-rich envelope and the helium core, it will take some observational effort to verify whether there are any real pulsating horizontal-branch stars which correspond with these models.
Effective emperatures are expected to be between 100000 and 200000 Ix. and. pulsation periods between 1: and 4 hours.	Effective temperatures are expected to be between 000 and 000 K, and pulsation periods between 1 and 4 hours.
With a factor of 10. enhancement of iron on a xickeround. metallicity Z= 0.02. a large instability region develops. even for /=1. which extends from <<18000 to 24000 IX (/= 3).	With a factor of 10 enhancement of iron on a background metallicity $Z=0.02$ , a large instability region develops, even for $l=1$, which extends from $<13\,000$ to $\sim 24\,000$ K $l=3$ ).
With such iron enhancement. there is always an overlap between p- and e-mode instability regions so that stars near the boundary would be expected to exhibit both modes simultaneously.	With such iron enhancement there is always an overlap between p- and g-mode instability regions so that stars near the boundary would be expected to exhibit both modes simultaneously.
Depleting the envelope hydrogen abundance tends to shift both the g-mode blue edge and the racialsp-moce red edge to lower temperatures.	Depleting the envelope hydrogen abundance tends to shift both the g-mode blue edge and the radial/p-mode red edge to lower temperatures.
Many of these results. have been demonstrate individually (Saio1993:Jeffery&Saio1998:Charpineteal.1996:Fontaineet 2003): this investigation has place hem in a more general context. as well as delivering some new results.	Many of these results have been demonstrated individually \citep{Sai93,Jef98,Cha96,Fon03}; this investigation has placed them in a more general context, as well as delivering some new results.
Of these. the rolle that mean molecular weigh dlavs in determining the instability boundaries for racial anc mode oscillations has implications for understanding the oscillations in IC'14026 variables.	Of these, the rôlle that mean molecular weight plays in determining the instability boundaries for radial and p-mode oscillations has implications for understanding the oscillations in EC14026 variables.
We have mace no cllor o justify the adopted abundances on physical grounds. we jwe simply sought parametric solutions which satisfy the observations.	We have made no effort to justify the adopted abundances on physical grounds, we have simply sought parametric solutions which satisfy the observations.
Racliative levitation (Chaver.Fontaine&We-seniael1995) is known to operate in extreme horizontal-anch stars and. provides a natural explanation. for the required. iron. enhancements.	Radiative levitation \citep{Cha95} is known to operate in extreme horizontal-branch stars and provides a natural explanation for the required iron enhancements.
The fact that it operates selectively accords well with our deduction that only iron should be enhanced.	The fact that it operates selectively accords well with our deduction that only iron should be enhanced.
Our discovery that e-moce oscillations mav be excited in blue horizontal branch stars with envelopes of “normal” composition has two-foldconsequences.	Our discovery that g-mode oscillations may be excited in blue horizontal branch stars with envelopes of “normal” composition has two-foldconsequences.
In. addition to the observational question already posed. it may. be less hard	In addition to the observational question already posed, it may be less hard
reach the observed mean value of e=0.35 within about 105 vears.	reach the observed mean value of $e=0.35$ within about $10^6$ years.
Investigations of eccentricity evolution due (o planet-disk interactions are plagued bv several major uncertainties.	Investigations of eccentricity evolution due to planet-disk interactions are plagued by several major uncertainties.
herefore. such an issue to be considered in another paper o appear elsewhere.	therefore, such an issue to be considered in another paper to appear elsewhere.
An interesting result concerning the clistribution ofdark matter and barvonic matter can be seen in Tables 2-5.	An interesting result concerning the distribution of dark matter and baryonic matter can be seen in Tables 2-5.
Phe inal diameter of the dark matter is always greater than he final diameter of the barvonic matter. when all physical orocesses considered in the present paper are included in the calculations.	The final diameter of the dark matter is always greater than the final diameter of the baryonic matter, when all physical processes considered in the present paper are included in the calculations.
To see the influence of the several physical processes on he evolution of a given negative density perturbation. in ρασας on the final diameters of the voids. we performec some calculations studying individually the effects of all the our cooling-heating mechanisms and also the influence of he photon drag (Table 6).	To see the influence of the several physical processes on the evolution of a given negative density perturbation, in particular on the final diameters of the voids, we performed some calculations studying individually the effects of all the four cooling-heating mechanisms and also the influence of the photon drag (Table 6).
Note that the dissipative barvonic processes do no define a particular characteristic void scale: their main effec is to segregate the dark matter and barvonic matter.	Note that the dissipative baryonic processes do not define a particular characteristic void scale; their main effect is to segregate the dark matter and baryonic matter.
From Tables 4 and 6 we note that the barvonic ciameter is à little bit lower than the dark matter diameter when al physical processes are included. and this result remains when we disregard each one of the cooling-heating mechanisms.	From Tables 4 and 6 we note that the baryonic diameter is a little bit lower than the dark matter diameter when all physical processes are included, and this result remains when we disregard each one of the cooling-heating mechanisms.
Another interesting result is obtained when we disregard all the processes.	Another interesting result is obtained when we disregard all the processes.
In this case. the final diameter of the barvonic matter is almost the same as the dark matter component.	In this case, the final diameter of the baryonic matter is almost the same as the dark matter component.
When we disregard. all the processes and. the pressure of the Universe. the final diameters are the same for both components. as expected. since the two [LIuids in these circumstances are subject only to gravity.	When we disregard all the processes and the pressure of the Universe, the final diameters are the same for both components, as expected, since the two fluids in these circumstances are subject only to gravity.
In Table 6 we see the influence of cach physical process one at a time.	In Table 6 we see the influence of each physical process one at a time.
When we disregard the photon drag. but maintain all the other processes and the pressure of the Universe. the final radius of the barvonic (dark) component &oes [rom 6.84(6.94)h'A\Ipe to 6.93(6.99)ἡ?Mpe.	When we disregard the photon drag, but maintain all the other processes and the pressure of the Universe, the final radius of the baryonic (dark) component goes from $6.84\; (6.94)\, h^{-1}{\rm Mpc}$ to $6.93\; (6.99)\, h^{-1}{\rm Mpc}$.
This occurs because the photon drag acts against the expansion of the negative density perturbations. inhibiting the growth of peculiar expansion. velocities.	This occurs because the photon drag acts against the expansion of the negative density perturbations, inhibiting the growth of peculiar expansion velocities.
The cooling mechanisms Lp. L, and Lg» have very similar effects. on the void. evolution. and their principal inlluence is related to the thermal evolution of the matter inside the void regions.	The cooling mechanisms $L_{\rm R}$, $L_\alpha$ and $L_{\rm H2}$ have very similar effects on the void evolution, and their principal influence is related to the thermal evolution of the matter inside the void regions.
The Compton heating-cooling when disregarded: reduces. slightly the dimension. of the void. oecause at high. redshifts this physical. process maintains he temperature of the matter inside the void close to the radiation temperature.	The Compton heating-cooling when disregarded reduces slightly the dimension of the void, because at high redshifts this physical process maintains the temperature of the matter inside the void close to the radiation temperature.
The above results also appear when we perform. the o)esent analysis for other models described in Table 1.	The above results also appear when we perform the present analysis for other models described in Table 1.
In xwticular. independent. of the spectra used. the physical orocesses produce very similar results concerning the final diameters of the barvonic and dark matter.	In particular, independent of the spectra used, the physical processes produce very similar results concerning the final diameters of the baryonic and dark matter.
We note that even using other spectra. like those used. by de Araujo Opher (1993) for example. this cillerence in diameter due to the presence of physical processes is maintained.	We note that even using other spectra, like those used by de Araujo Opher (1993) for example, this difference in diameter due to the presence of physical processes is maintained.
With all the physical processes included. there exists a transition region with thickness of up to ~2.55.“Alpe. which depends on the mass of the perturbation. the normalization of the spectrum and the respective density parameters for the dark ancl barvonic components as well as the cosmological constant.	With all the physical processes included, there exists a transition region with thickness of up to $\sim 2.5\, h^{-1}{\rm Mpc}$, which depends on the mass of the perturbation, the normalization of the spectrum and the respective density parameters for the dark and baryonic components as well as the cosmological constant.
Phese transition regions have sizes comparable to the clametors of galaxies ane clusters of ealaxies.	These transition regions have sizes comparable to the diameters of galaxies and clusters of galaxies.
The model 26. for cxample. has a transition region of ~2.5htMpe thickness for a perturbation. with mass LOMM...	The model 26, for example, has a transition region of $\sim 2.5\, h^{-1}{\rm Mpc}$ thickness for a perturbation with mass $10^{15}{\rm M}_\odot$.
Within this region the density contrast of the cold dark matter is Opp=0.628 (its diameter is 53.5hf!Mpe).	Within this region the density contrast of the cold dark matter is $\delta_{\rm DF}=-0.628$ (its diameter is $53.5\, h^{-1}{\rm Mpc}$ ).
‘Thus. the mass density of the dark matter is py.=0.372044. where paris the ambient mass density of the dark matter (that is. the mass density of the dark matter present in the Universe).	Thus, the mass density of the dark matter is $\rho_{\rm d}=0.372\rho_{\rm du}$, where $\rho_{\rm du}$ is the ambient mass density of the dark matter (that is, the mass density of the dark matter present in the Universe).
As the bouncary of the barvonic matter ds. at 48.7!Mpe. within the transition region the density contrast of the harvonie matter is 05=0. and so pi,=pia.	As the boundary of the baryonic matter is at $48.7\, h^{-1}{\rm Mpc}$ , within the transition region the density contrast of the baryonic matter is $\delta_{\rm B}=0 $, and so $\rho_{\rm b}=\rho_{\rm bu}$.
The relation between the barvonic and. dark. components present within this transition region is given by ‘Thus. in this region the amount of barvonic matter is approximately equal to the amount of non-barvonie dark matter. although this universe model has three times more dark matter (ο= 0.15) than baryonic matter (3,=0.05).	The relation between the baryonic and dark components present within this transition region is given by Thus, in this region the amount of baryonic matter is approximately equal to the amount of non-baryonic dark matter, although this universe model has three times more dark matter $\Omega_{\rm d}=0.15$ ) than baryonic matter $\Omega_{\rm b}=0.05$ ).
Consequently. galaxies formed near voids and. within he transition region could. during their formation process. jwe a lower content of cold dark matter than those galaxies ormed in regions lar from voids.	Consequently, galaxies formed near voids and within the transition region could, during their formation process, have a lower content of cold dark matter than those galaxies formed in regions far from voids.
This result. leads us to conclude that. if one uses an analysis of the content of dark o barvonic matter for the galaxies in different regions. one could infer very dillerent. results. rellecting therefore only ocal values.	This result leads us to conclude that, if one uses an analysis of the content of dark to baryonic matter for the galaxies in different regions, one could infer very different results, reflecting therefore only local values.
Let us consider another example.	Let us consider another example.
The transition zone of model 26. for a perturbation with mass LOYAL.. gives f£pa(0/04)~2.3. although in such a universe model the amount of dark matter is three times greater than the amount of barvonic matter.	The transition zone of model 26, for a perturbation with mass $10^{13}{\rm M}_\odot$, gives $\rho_{\rm b} /\rho_{\rm d} \sim 7 (\Omega_{\rm b}/\Omega_{\rm d}) \sim 2.3$, although in such a universe model the amount of dark matter is three times greater than the amount of baryonic matter.
lt is worth stressing that galaxies formed. out of the transition zone and that escape to there could. have a diferent relation between the baryonic and dark components than that found in the transition zone.	It is worth stressing that galaxies formed out of the transition zone and that escape to there could have a different relation between the baryonic and dark components than that found in the transition zone.

Also, SiC feature at around 11.3 wm can be seen in carbon stars.

Also, SiC feature at around 11.3 $\mu$ m can be seen in carbon stars.

Interestingly, '13 um feature’ is seen in M-type and/or S-type giants.

Interestingly, so-called '13 $\mu$ m feature' is seen in M-type and/or S-type giants.

This feature is probably due to aluminum oxides (e.g., Posch et al. 1999)).

This feature is probably due to aluminum oxides (e.g., Posch et al. \cite{posch1999}) ).

Sloan et al. (2003b))

Sloan et al. \cite{sloan2003b}) )

suggested that this feature tends to be stronger in systems with lower infrared excesses and thus lower mass-loss rates (e.g., Onaka et al. 1989;;

suggested that this feature tends to be stronger in systems with lower infrared excesses and thus lower mass-loss rates (e.g., Onaka et al. \cite{onaka1989};

Kozasa Sogawa 1997)).

Kozasa Sogawa \cite{kozasa1997}) ).

Interestingly, Sloan et al. (1996))

Interestingly, Sloan et al. \cite{sloan1996}) )

reported that the 13 um feature is preferentially detected in semi-regular or irregular variables.

reported that the 13 $\mu$ m feature is preferentially detected in semi-regular or irregular variables.

Based on these available data of galactic counterparts, we suggest that low mass-loss rate M-type giants,

post-starburst galaxies can be further divided iuto k | a anda | k type galaxies. (

post-starburst galaxies can be further divided into k + a and a + k type galaxies. (

Fraux 1993: Dressler et al.

Franx 1993; Dressler et al.

1999: Pogeianti ct al.

1999; Poggianti et al.

1999).

Northern. southern aud casteru chuups are mostly populated by k type galaxies aud a fow k | a and ai | k galaxies are also seen im these clumps.

Northern, southern and eastern clumps are mostly populated by k type galaxies and a few k + a and a + k galaxies are also seen in these clumps.

The classification as post-starburst galaxies of two out of ten cluster 1iucnibers may be uncertain due to the possible alteration of the IL; equivalent width by the tellure correction.

The classification as post-starburst galaxies of two out of ten cluster members may be uncertain due to the possible alteration of the $_{\delta}$ equivalent width by the telluric correction.

Alost passive galaxies have red colours (I.ἷνςs= 5) aud most of them populate the cluster red sequence (CRS) in the colouranaguitude diagram (CAID).

Most passive galaxies have red colours $R-K_s \simeq 5$ ) and most of them populate the cluster red sequence (CRS) in the colour-magnitude diagram (CMD).

In Fig. 9..

In Fig. \ref{rcs},

we show the distribution of cluster 1iemibers that have eround based LRIS and Sofl photometry iu two CAIDs.

we show the distribution of cluster members that have ground based LRIS and SofI photometry in two CMDs.

Open circles are passive galaxies aud solid triangles are star-forming galaxies.

Open circles are passive galaxies and solid triangles are star-forming galaxies.

The solid lines correspond to linear fits (ROAL=0.027|5.1E in upper panel aud PofF=0.022F|3.39 in lower panel) to the data within the regions defined in the caption of Fie. 9..

The solid lines correspond to linear fits $R-K_\mathrm{s} = -0.027 \ K_\mathrm{s} + 5.44$ in upper panel and $V-I = -0.022 \ I + 3.39$ in lower panel) to the data within the regions defined in the caption of Fig. \ref{rcs}.

Dashed lues indicate the 1-0 level above and below the fits.

Dashed lines indicate the $\sigma$ level above and below the fits.

If we compare the CRS of the RA. ovs dv, diagram to the models computed by Kodama Arimoto (1997). we infer that the colours of the passive galaxies in the CRS are cousisteut with an carly formation epoch 2).

If we compare the CRS of the $R-K_\mathrm{s}$ vs $K_\mathrm{s}$ diagram to the models computed by Kodama Arimoto (1997), we infer that the colours of the passive galaxies in the CRS are consistent with an early formation epoch $z \simgreat 2$ ).

However. we note the presence of a tail of fainter (ht.κ18.5) passive galaxies extending to ιο (L0XROON.X5.7) colours. off the CRS.

However, we note the presence of a tail of fainter $K_s \simgreat 18.5$ ) passive galaxies extending to bluer $4.0 \simless R-K_s \simless 5.7$ ) colours, off the CRS.

Passive galaxies are also observed outside the hieli-deusitv peaks. however the outskirts of aare dominated by star-forming galaxies.

Passive galaxies are also observed outside the high-density peaks, however the outskirts of are dominated by star-forming galaxies.

The segregation between star-forming and passive ealaxies in lis in agreement with a picture in which the star formation activity in galaxies 1s suppressed diving the iufall of these ealaxies into the highest density zoucs.

The segregation between star-forming and passive galaxies in is in agreement with a picture in which the star formation activity in galaxies is suppressed during the infall of these galaxies into the highest density zones.

An analysis of tle physical properties of the star-forming galaxy population in bbased on the combination of VLT spectroscopic aud HST/ACS imagine data. together with implications for cluster galaxy evolution. will be preseuted iu Honieier et al.

An analysis of the physical properties of the star-forming galaxy population in based on the combination of VLT spectroscopic and HST/ACS imaging data, together with implications for cluster galaxy evolution, will be presented in Homeier et al.

2001.

2004.

In Fie. 10.

In Fig. \ref{colcol},

we show the distribution of cluster members m two coloum-colour diagrams.

we show the distribution of cluster members in two colour-colour diagrams.

The sviibols are the same as those in Fie. 9..

The symbols are the same as those in Fig. \ref{rcs}.

For reference. we also slow the tracks as a function of redshift for au elliptical (Ell) and three late-type (She. Scd and Br) galaxies using the teiiplate library of Coleman. Wu. Weedman (1980). whose spectral cuerey distributious (SEDs) have becu extended to. the rear-IR and far-UV using Druzual Charlot (1993. with recent updates) models;

For reference, we also show the tracks as a function of redshift for an elliptical (Ell) and three late-type (Sbc, Scd and Irr) galaxies using the template library of Coleman, Wu, Weedman (1980), whose spectral energy distributions (SEDs) have been extended to the near-IR and far-UV using Bruzual Charlot (1993, with recent updates) models.

The tracks iuclude zeorrectious oulv.

The tracks include k-corrections only.

The expected colours for the differeut galaxy types at the cluster redshift are indicatec with stars.

The expected colours for the different galaxy types at the cluster redshift are indicated with stars.

Passive galaxies are located in the upper right part of thi colour-colour diagrams (see Fig. 10)).

Passive galaxies are located in the upper right part of both colour-colour diagrams (see Fig. \ref{colcol}) ).

The observe offset between the locus of these galaxies iux their expected colours from the models is consistent with passive evolution of earbv-tvpe galaxies formed at 2~1.9.

The observed offset between the locus of these galaxies and their expected colours from the models is consistent with passive evolution of early-type galaxies formed at $z \sim 1.9$.

Finally. we note that the brightest cluster ΠΟΡΟΙ of the southern clump (ID=387) is redder than the brightes cluster ποσο of the northern chuup (ID-701).

Finally, we note that the brightest cluster member of the southern clump (ID=387) is redder than the brightest cluster member of the northern clump (ID=701).

The differences in the colours are A(RA.)=OdT ane AVI)=(5 with crrors of 6(RWoy=0015 and οΕΠ=0.012 respectively.

The differences in the colours are $\Delta(R-K_\mathrm{s})=0.17$ and $\Delta(V-I)=0.15$ with errors of $\delta(R-K_\mathrm{s})=0.015$ and $\delta(V-I)=0.012$ respectively.

These differences in colour ect bigeer if we compare ID=3a7 with the pair of xieht galaxies (ID=1166 aud ID=1167) located in the core of the northern chump.

These differences in colour get bigger if we compare ID=387 with the pair of bright galaxies (ID=1466 and ID=1467) located in the core of the northern clump.

Calaxy ID=3s87 turus ou o he 0.25 magnitudes redder in RoA, and ~0.2 naenitudes redder in V.{ than both galaxies at the core of the northern subcluster.

Galaxy ID=387 turns out to be $\sim 0.25$ magnitudes redder in $R-K_\mathrm{s}$ and $\sim 0.24$ magnitudes redder in $V-I$ than both galaxies at the core of the northern subcluster.

Assiuuiug that these ealaxies ive smaular metallicities. such a cifference iu colour may )o an indication of a difference in stellar ages.

Assuming that these galaxies have similar metallicities, such a difference in colour may be an indication of a difference in stellar ages.

Towever. nore accurate piomoetry. sucli as. e.g...o ACS photometry.

However, more accurate photometry, such as, e.g., ACS photometry,

core.

For these models a partial reflection at the bottom of the helium-rich. envelope reduces the amplitude and. hence radiative damping in the core.

For these models a partial reflection at the bottom of the helium-rich envelope reduces the amplitude and hence radiative damping in the core.

Since only few e-modes are excited and the evolutionary speed is fast. these g-mocdes are not very important observationalls.

Since only few g-modes are excited and the evolutionary speed is fast, these g-modes are not very important observationally.

Our original goals were to understand oscillations observed in the helium-rich. subdwarl 127116. ancl in. the DOGLITIG variables.

Our original goals were to understand oscillations observed in the helium-rich subdwarf $-14^{\circ}116$, and in the PG1716 variables.

We therefore carried out a broad review of Fe-bump driven pulsational instability in low-massstars’.

We therefore carried out a broad review of Fe-bump driven pulsational instability in low-mass.

ὃν dist considering radial modes. Νο have demonstrated: the essential rolle of chemical composition such that instability increases with the contrast. between the iron-bump. opacity and. other opacity sources.

By first considering radial modes, we have demonstrated the essential rôlle of chemical composition such that instability increases with the contrast between the iron-bump opacity and other opacity sources.

This increased Contrast may be achieved either. hy increasing the iron abundance. confirming earlier work by Charpinetetal.(2001). or by reducing the hyelrogen abundance (Jelfery&SaloLOOS).

This increased contrast may be achieved either by increasing the iron abundance, confirming earlier work by \citet{Cha01} or by reducing the hydrogen abundance \citep{Jef98}.

. At least one of these is necessary to excite oscillations in all of the Fe-bump pulsators discovered to date.

At least one of these is necessary to excite oscillations in all of the Fe-bump pulsators discovered to date.

We have further demonstrated that the bluc-edge for racial instability is allectecl by the mean molecular weight in the stellar envelope. so that increasing the iron abundance alone provides a bluer instability region than increasing the iron abundance in concert with all elements heavier than helium.

We have further demonstrated that the blue-edge for radial instability is affected by the mean molecular weight in the stellar envelope, so that increasing the iron abundance alone provides a bluer instability region than increasing the iron abundance in concert with all elements heavier than helium.

The former is required to explain the locus of the EC14026 instability region. making the general assumption that the properties of non-racial pamocdes moces are closely linked to the corresponding radial mode of the same racial orcler.

The former is required to explain the locus of the EC14026 instability region, making the general assumption that the properties of non-radial p-modes modes are closely linked to the corresponding radial mode of the same radial order.

Furthermore. the blue-edge also depends on the radial order of the oscillations. so that higher-order modes may be found in hotter stars.

Furthermore, the blue-edge also depends on the radial order of the oscillations, so that higher-order modes may be found in hotter stars.

By comparing theoretical and observed periods for EC14026 stars. we have shown that low-order or funcamental mode racial oscillations are only likely to be seen in the coolest EC14026 stars and that the hottest stars must oscillate in higher-order moces.

By comparing theoretical and observed periods for EC14026 stars, we have shown that low-order or fundamental mode radial oscillations are only likely to be seen in the coolest EC14026 stars and that the hottest stars must oscillate in higher-order modes.

However. we have been unable to explain the oscillations in the helium-rich star 4116: the observed. periods are simply too long compared with the L/AJ ratio obtained [rom spectroscopy.

However, we have been unable to explain the oscillations in the helium-rich star $-14^{\circ}116$; the observed periods are simply too long compared with the $L/M$ ratio obtained from spectroscopy.

A lower surface gravity would. definitely help.

A lower surface gravity would definitely help.

Considering non-radial pulsations. we have focused on e-mocde instability and in particular on instability in modes with spherical degree /«4.

Considering non-radial pulsations, we have focused on g-mode instability and in particular on instability in modes with spherical degree $l<4$.

We have discovered a small g-mode instability island on the blue horizontal branch which does not require iron enhancement or hydrogen depletion.

Llowever. since it requires Z=0.02 and a partial rellection of he wave at the interface between the hydrogen-rich envelope and the helium core. it will take some observational cllort o verily whether there are any. real pulsating horizontal-anch stars which correspond with these models.

However, since it requires $Z=0.02$ and a partial reflection of the wave at the interface between the hydrogen-rich envelope and the helium core, it will take some observational effort to verify whether there are any real pulsating horizontal-branch stars which correspond with these models.

Effective emperatures are expected to be between 100000 and 200000 Ix. and. pulsation periods between 1: and 4 hours.

Effective temperatures are expected to be between 000 and 000 K, and pulsation periods between 1 and 4 hours.

With a factor of 10. enhancement of iron on a xickeround. metallicity Z= 0.02. a large instability region develops. even for /=1. which extends from <<18000 to 24000 IX (/= 3).

With a factor of 10 enhancement of iron on a background metallicity $Z=0.02$ , a large instability region develops, even for $l=1$, which extends from $<13\,000$ to $\sim 24\,000$ K $l=3$ ).

With such iron enhancement. there is always an overlap between p- and e-mode instability regions so that stars near the boundary would be expected to exhibit both modes simultaneously.

With such iron enhancement there is always an overlap between p- and g-mode instability regions so that stars near the boundary would be expected to exhibit both modes simultaneously.

Depleting the envelope hydrogen abundance tends to shift both the g-mode blue edge and the racialsp-moce red edge to lower temperatures.

Depleting the envelope hydrogen abundance tends to shift both the g-mode blue edge and the radial/p-mode red edge to lower temperatures.

Many of these results. have been demonstrate individually (Saio1993:Jeffery&Saio1998:Charpineteal.1996:Fontaineet 2003): this investigation has place hem in a more general context. as well as delivering some new results.

Many of these results have been demonstrated individually \citep{Sai93,Jef98,Cha96,Fon03}; this investigation has placed them in a more general context, as well as delivering some new results.

Of these. the rolle that mean molecular weigh dlavs in determining the instability boundaries for racial anc mode oscillations has implications for understanding the oscillations in IC'14026 variables.

Of these, the rôlle that mean molecular weight plays in determining the instability boundaries for radial and p-mode oscillations has implications for understanding the oscillations in EC14026 variables.

We have mace no cllor o justify the adopted abundances on physical grounds. we jwe simply sought parametric solutions which satisfy the observations.

We have made no effort to justify the adopted abundances on physical grounds, we have simply sought parametric solutions which satisfy the observations.

Racliative levitation (Chaver.Fontaine&We-seniael1995) is known to operate in extreme horizontal-anch stars and. provides a natural explanation. for the required. iron. enhancements.

Radiative levitation \citep{Cha95} is known to operate in extreme horizontal-branch stars and provides a natural explanation for the required iron enhancements.

The fact that it operates selectively accords well with our deduction that only iron should be enhanced.

Our discovery that e-moce oscillations mav be excited in blue horizontal branch stars with envelopes of “normal” composition has two-foldconsequences.

Our discovery that g-mode oscillations may be excited in blue horizontal branch stars with envelopes of “normal” composition has two-foldconsequences.

In. addition to the observational question already posed. it may. be less hard

In addition to the observational question already posed, it may be less hard

reach the observed mean value of e=0.35 within about 105 vears.

reach the observed mean value of $e=0.35$ within about $10^6$ years.

Investigations of eccentricity evolution due (o planet-disk interactions are plagued bv several major uncertainties.

Investigations of eccentricity evolution due to planet-disk interactions are plagued by several major uncertainties.

herefore. such an issue to be considered in another paper o appear elsewhere.

therefore, such an issue to be considered in another paper to appear elsewhere.

An interesting result concerning the clistribution ofdark matter and barvonic matter can be seen in Tables 2-5.

An interesting result concerning the distribution of dark matter and baryonic matter can be seen in Tables 2-5.

Phe inal diameter of the dark matter is always greater than he final diameter of the barvonic matter. when all physical orocesses considered in the present paper are included in the calculations.

The final diameter of the dark matter is always greater than the final diameter of the baryonic matter, when all physical processes considered in the present paper are included in the calculations.

To see the influence of the several physical processes on he evolution of a given negative density perturbation. in ρασας on the final diameters of the voids. we performec some calculations studying individually the effects of all the our cooling-heating mechanisms and also the influence of he photon drag (Table 6).

To see the influence of the several physical processes on the evolution of a given negative density perturbation, in particular on the final diameters of the voids, we performed some calculations studying individually the effects of all the four cooling-heating mechanisms and also the influence of the photon drag (Table 6).

Note that the dissipative barvonic processes do no define a particular characteristic void scale: their main effec is to segregate the dark matter and barvonic matter.

Note that the dissipative baryonic processes do not define a particular characteristic void scale; their main effect is to segregate the dark matter and baryonic matter.

From Tables 4 and 6 we note that the barvonic ciameter is à little bit lower than the dark matter diameter when al physical processes are included. and this result remains when we disregard each one of the cooling-heating mechanisms.

From Tables 4 and 6 we note that the baryonic diameter is a little bit lower than the dark matter diameter when all physical processes are included, and this result remains when we disregard each one of the cooling-heating mechanisms.

Another interesting result is obtained when we disregard all the processes.

In this case. the final diameter of the barvonic matter is almost the same as the dark matter component.

In this case, the final diameter of the baryonic matter is almost the same as the dark matter component.

When we disregard. all the processes and. the pressure of the Universe. the final diameters are the same for both components. as expected. since the two [LIuids in these circumstances are subject only to gravity.

When we disregard all the processes and the pressure of the Universe, the final diameters are the same for both components, as expected, since the two fluids in these circumstances are subject only to gravity.

In Table 6 we see the influence of cach physical process one at a time.

In Table 6 we see the influence of each physical process one at a time.

When we disregard the photon drag. but maintain all the other processes and the pressure of the Universe. the final radius of the barvonic (dark) component &oes [rom 6.84(6.94)h'A\Ipe to 6.93(6.99)ἡ?Mpe.

When we disregard the photon drag, but maintain all the other processes and the pressure of the Universe, the final radius of the baryonic (dark) component goes from $6.84\; (6.94)\, h^{-1}{\rm Mpc}$ to $6.93\; (6.99)\, h^{-1}{\rm Mpc}$.

This occurs because the photon drag acts against the expansion of the negative density perturbations. inhibiting the growth of peculiar expansion. velocities.

This occurs because the photon drag acts against the expansion of the negative density perturbations, inhibiting the growth of peculiar expansion velocities.

The cooling mechanisms Lp. L, and Lg» have very similar effects. on the void. evolution. and their principal inlluence is related to the thermal evolution of the matter inside the void regions.

The cooling mechanisms $L_{\rm R}$, $L_\alpha$ and $L_{\rm H2}$ have very similar effects on the void evolution, and their principal influence is related to the thermal evolution of the matter inside the void regions.

The Compton heating-cooling when disregarded: reduces. slightly the dimension. of the void. oecause at high. redshifts this physical. process maintains he temperature of the matter inside the void close to the radiation temperature.

The Compton heating-cooling when disregarded reduces slightly the dimension of the void, because at high redshifts this physical process maintains the temperature of the matter inside the void close to the radiation temperature.

The above results also appear when we perform. the o)esent analysis for other models described in Table 1.

The above results also appear when we perform the present analysis for other models described in Table 1.

In xwticular. independent. of the spectra used. the physical orocesses produce very similar results concerning the final diameters of the barvonic and dark matter.

In particular, independent of the spectra used, the physical processes produce very similar results concerning the final diameters of the baryonic and dark matter.

We note that even using other spectra. like those used. by de Araujo Opher (1993) for example. this cillerence in diameter due to the presence of physical processes is maintained.

We note that even using other spectra, like those used by de Araujo Opher (1993) for example, this difference in diameter due to the presence of physical processes is maintained.

With all the physical processes included. there exists a transition region with thickness of up to ~2.55.“Alpe. which depends on the mass of the perturbation. the normalization of the spectrum and the respective density parameters for the dark ancl barvonic components as well as the cosmological constant.

With all the physical processes included, there exists a transition region with thickness of up to $\sim 2.5\, h^{-1}{\rm Mpc}$, which depends on the mass of the perturbation, the normalization of the spectrum and the respective density parameters for the dark and baryonic components as well as the cosmological constant.

Phese transition regions have sizes comparable to the clametors of galaxies ane clusters of ealaxies.

These transition regions have sizes comparable to the diameters of galaxies and clusters of galaxies.

The model 26. for cxample. has a transition region of ~2.5htMpe thickness for a perturbation. with mass LOMM...

The model 26, for example, has a transition region of $\sim 2.5\, h^{-1}{\rm Mpc}$ thickness for a perturbation with mass $10^{15}{\rm M}_\odot$.

Within this region the density contrast of the cold dark matter is Opp=0.628 (its diameter is 53.5hf!Mpe).

Within this region the density contrast of the cold dark matter is $\delta_{\rm DF}=-0.628$ (its diameter is $53.5\, h^{-1}{\rm Mpc}$ ).

‘Thus. the mass density of the dark matter is py.=0.372044. where paris the ambient mass density of the dark matter (that is. the mass density of the dark matter present in the Universe).

Thus, the mass density of the dark matter is $\rho_{\rm d}=0.372\rho_{\rm du}$, where $\rho_{\rm du}$ is the ambient mass density of the dark matter (that is, the mass density of the dark matter present in the Universe).

As the bouncary of the barvonic matter ds. at 48.7!Mpe. within the transition region the density contrast of the harvonie matter is 05=0. and so pi,=pia.

As the boundary of the baryonic matter is at $48.7\, h^{-1}{\rm Mpc}$ , within the transition region the density contrast of the baryonic matter is $\delta_{\rm B}=0 $, and so $\rho_{\rm b}=\rho_{\rm bu}$.

The relation between the barvonic and. dark. components present within this transition region is given by ‘Thus. in this region the amount of barvonic matter is approximately equal to the amount of non-barvonie dark matter. although this universe model has three times more dark matter (ο= 0.15) than baryonic matter (3,=0.05).

The relation between the baryonic and dark components present within this transition region is given by Thus, in this region the amount of baryonic matter is approximately equal to the amount of non-baryonic dark matter, although this universe model has three times more dark matter $\Omega_{\rm d}=0.15$ ) than baryonic matter $\Omega_{\rm b}=0.05$ ).

Consequently. galaxies formed near voids and. within he transition region could. during their formation process. jwe a lower content of cold dark matter than those galaxies ormed in regions lar from voids.

Consequently, galaxies formed near voids and within the transition region could, during their formation process, have a lower content of cold dark matter than those galaxies formed in regions far from voids.

This result. leads us to conclude that. if one uses an analysis of the content of dark o barvonic matter for the galaxies in different regions. one could infer very dillerent. results. rellecting therefore only ocal values.

This result leads us to conclude that, if one uses an analysis of the content of dark to baryonic matter for the galaxies in different regions, one could infer very different results, reflecting therefore only local values.

Let us consider another example.

The transition zone of model 26. for a perturbation with mass LOYAL.. gives f£pa(0/04)~2.3. although in such a universe model the amount of dark matter is three times greater than the amount of barvonic matter.

The transition zone of model 26, for a perturbation with mass $10^{13}{\rm M}_\odot$, gives $\rho_{\rm b} /\rho_{\rm d} \sim 7 (\Omega_{\rm b}/\Omega_{\rm d}) \sim 2.3$, although in such a universe model the amount of dark matter is three times greater than the amount of baryonic matter.

lt is worth stressing that galaxies formed. out of the transition zone and that escape to there could. have a diferent relation between the baryonic and dark components than that found in the transition zone.

It is worth stressing that galaxies formed out of the transition zone and that escape to there could have a different relation between the baryonic and dark components than that found in the transition zone.