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Using this value aud the SCOL velocity dispersion. we fud an L-band virial mass of M,=7.0c1.3«Lo? M... for AL82-F. This is significantly lower than the opically-derived SCL mass estimate of L240.1ς109 AL. de to more accurate measure of the cluster’s halt-lishbit radius (sce Section [.1)). | Using this value and the SG01 velocity dispersion, we find an $I$ -band virial mass of $M_I = 7.0 \pm 1.2 \times 10^5$ $_{\odot}$ for M82-F. This is significantly lower than the optically-derived SG01 mass estimate of $1.2 \pm 0.1 \times 10^6$ $_{\odot}$ due to more accurate measure of the cluster's half-light radius (see Section \ref{acsdata}) ). |
Phoonietrv was derived bv iutegratiug over the fitted Ning uodels anc applviug the appropriate conversiou factors for each iustrineut aud filter. | Photometry was derived by integrating over the fitted King models and applying the appropriate conversion factors for each instrument and filter. |
Details of the PSF fitting photometry and spectral cucrey distribution are presened in a companion paper €2).. | Details of the PSF fitting photometry and spectral energy distribution are presented in a companion paper \citep{vacca04}. |
The observed cluster i-banud muuimosities as defined iu ον, not corrected for extinction. are {Εν=3840.6.10° L. aud Lrigy=Loto.107 L.. | The observed cluster in-band luminosities as defined in \citet{mccrady03}, not corrected for extinction, are $L_{F814} = 3.8 \pm 0.6 \times 10^5$ $_{\odot}$ and $L_{F160} = 4.0 \pm 0.7 \times 10^5$ $_{\odot}$. |
By comparing the derived light-to-mmass ΑΓ) ratio of AIS2-F to population svuthesis models, we may characterize the IME of the cluster. | By comparing the derived light-to-mass $L/M$ ) ratio of M82-F to population synthesis models, we may characterize the IMF of the cluster. |
Critical to this analysis are the cluster age aud line-of-sight extinction. | Critical to this analysis are the cluster age and line-of-sight extinction. |
The spectruii of MBS2-F. imunecdiately places upper aud ower bounds ou he clusters age. | The spectrum of M82-F immediately places upper and lower bounds on the cluster's age. |
There is no evidence of nebular euussion in the Z7- or A-baud. which ucdicates he absence of O stars and a iium cluster aee of 6 που. | There is no evidence of nebular emission in the $H$ - or $K$ -band, which indicates the absence of O stars and a minimum cluster age of 6--7 Myr. |
? finds none of the features expected of AGB stars for the nuclear region of AIS82. setting an upper int of ~10? vears. | \citet{natascha98} finds none of the features expected of AGB stars for the nuclear region of M82, setting an upper limit of $\sim 10^8$ years. |
For more precise nuits ou the age. we turn to population svuthesis modelling. | For more precise limits on the age, we turn to population synthesis modelling. |
As noted in Sectki 3.. the F160W light of NIS2-F most closely matches cluplate spectral vpes in the ranee of INtTMOI. ? used Starburst99 population svuthesis models to determine the Hux-welehted average spectral type as a function of age for ορσνα] stellar populatious. | As noted in Section \ref{analysis}, the F160W light of M82-F most closely matches template spectral types in the range of K4I–M0I. \citet{gilbert02th} used Starburst99 population synthesis models to determine the flux-weighted average spectral type as a function of age for coeval stellar populations. |
Iu the Z7-baud. the cluster light is dominated by Ntλ0 stars for a brief period around 15 αν ancl during the ages of ~ LO60 Myr, | In the $H$ -band, the cluster light is dominated by K4–M0 stars for a brief period around 15 Myr and during the ages of $\sim$ 40–60 Myr. |
SCGOL used Starburst99 models to fit the Πὸ and Πο Τ absorption profiles iu optical specra. | SG01 used Starburst99 models to fit the $\delta$ and He I absorption profiles in optical spectra. |
Thev found that the best fits to the wines of the Balmer line aud depth of the helm liue suggest au age of 60+20 Myr. | They found that the best fits to the wings of the Balmer line and depth of the helium line suggest an age of $60 \pm 20$ Myr. |
This is cousistent with the dominant spectral type of the /7-baud light. aud we thereore adopt an age range of LO60 Myr for M82-F. SGHW used BV photometry to determine line-of-sight extinction of Ay=2.8. | This is consistent with the dominant spectral type of the $H$ -band light, and we therefore adopt an age range of 40–60 Myr for M82-F. SG01 used $BVI$ photometry to determine line-of-sight extinction of $A_V = 2.8$. |
Applving the ? extinction law (Ry= 3.1) to the this value gives «πιο.=0.53 aud Arawo=1.65. | Applying the \citet{cardelli89} extinction law $R_V =
3.1$ ) to the this value gives $A_{F160W} = 0.53$ and $A_{F814W} =
1.63$. |
As an independent test. we calculate svuthetic colors for the nem-IR dominant evolved IKlMO stars based ou the ? stellar spectral library. | As an independent test, we calculate synthetic colors for the near-IR dominant evolved K4–M0 stars based on the \citet{pickles98} stellar spectral library. |
The svuthetic [FLGOW| |F222M| color for a I&II star is 0.35. versus 0.52 for MOI. From our plotometiy. we find [FIGOW [F222MI] = 0.364FEO.01 for AIS2-F. at the blue cud of the| expected color rauge. | The synthetic [F160W] $-$ [F222M] color for a K4I star is 0.35, versus 0.52 for M0I. From our photometry, we find [F160W] $-$ [F222M] = $0.36 \pm 0.04$ for M82-F, at the blue end of the expected color range. |
This iniplies that. /7-baud extinction to MBS2-F. is quite simall. essentially neglieible. | This implies that $H$ -band extinction to M82-F is quite small, essentially negligible. |
The optical (600-800 111) Light of the cluster is expected to be dominated by KII stars (SCOL). | The optical (600-800 nm) light of the cluster is expected to be dominated by KII stars (SG01). |
Syuthetic F555W| [F5LIW] colors for these stars in the Pickles library range from 1.19 to 1.15. versus the measured 1.7140.01 for AIS2-F. For a standard (πι=3.1) interstellar extinction curve. this color excess gives «ρα=031 to 0.66 mae. | Synthetic [F555W] $-$ [F814W] colors for these stars in the Pickles library range from 1.19 to 1.45, versus the measured $1.71 \pm 0.04$ for M82-F. For a standard $R_V = 3.1$ ) interstellar extinction curve, this color excess gives $A_{F814W} =
0.34$ to 0.66 mag. |
To caleulate the light-to-nass ratio. we need to deredden the cluster. | To calculate the light-to-mass ratio, we need to deredden the cluster. |
We adopt extinction of dgio=0.0!uu aud πω=0.50.2 based on the svuthetic photometry results. | We adopt extinction of $A_{F160W} = 0.0 ^{+0.1}_{-0.0}$ and $A_{F814W}
= 0.5 \pm 0.2$ based on the synthetic photometry results. |
For sake of comparison. Figure ?7 reflects both these estimates aud the extinctions implied by the SCUL estimate of ely. | For sake of comparison, Figure \ref{sb99plot} reflects both these estimates and the extinctions implied by the SG01 estimate of $A_V$. |
The Calactic dust map of ? indicates Galactic extinction along the line of sight to AIS2-F of Ay=0.Ls. Aran=028 and Ariouy=0.09. | The Galactic dust map of \citet{schlegel98} indicates Galactic extinction along the line of sight to M82-F of $A_V = 0.48$, $A_{F814W} = 0.28$ and $A_{F160W} = 0.09$. |
This provides a for the estimated extinction. | This provides a cross-reference for the estimated extinction. |
We determine the huninmositv aud virial mass of the cluster independent of auv assumptions about the IME. | We determine the luminosity and virial mass of the cluster independent of any assumptions about the IMF. |
This is iu contrast to photometric mass deteriinations. which are based on cluster colors and ages aud mist herexe assuue an TMF. | This is in contrast to photometric mass determinations, which are based on cluster colors and ages and must therefore assume an IMF. |
Our inethod enables us to constrain the cluster IME by comparing observed lisht-to-lass ratios mn various wavebands to population svutlesis nodels (2?).. | Our method enables us to constrain the cluster IMF by comparing observed light-to-mass ratios in various wavebands to population synthesis models \citep{sternberg98,mccrady03}. |
Dv:ipplving the adopted exiuction corrections for F. we find de-redcdened lisht-toaunass ratios of L/M=L6xU (L. AL.) at L6;120 aud L/A/=0.9+0.2 at Span. Figure 77. compares he measured £/AL ratio to »pulatiou svuthesis model xedietious for two fiducial IME forms. | By applying the adopted extinction corrections for M82-F, we find de-reddened light-to-mass ratios of $L/M = 0.6 \pm 0.1$ $_\odot$ $_\odot$ ) at $\mu$ m and $L/M = 0.9 \pm 0.2$ at $\mu$ m. Figure \ref{sb99plot} compares the measured $L/M$ ratio to population synthesis model predictions for two fiducial IMF forms. |
We use Starburst99 version L0. with au istantancous burst. solar mieallicity (2). and the Ilillier Pauldrach atinosphere moels. | We used Starburst99 version 4.0, with an instantaneous burst, solar metallicity \citep{mcleod93} and the Hillier Pauldrach atmosphere models. |
The derived. £/AL ratio of MBS2-F is too Heh for the standard ? IMFE over the full range of sCllar masses (0.100. ADL. | The derived $L/M$ ratio of M82-F is too high for the standard \citet{kroupa01} IMF over the full range of stellar masses (0.1–100 $_{\odot}$ ). |
Rather it appears that the clusters IAIF is eCficient in low-1ass stars. | Rather it appears that the cluster's IMF is deficient in low-mass stars. |
For example. the L/ÀA ratio is roughly cousisteut with a? IME truncated at a lower uass of about 2 AL... | For example, the $L/M$ ratio is roughly consistent with a \citet{salpeter55} IMF truncated at a lower mass of about 2 $_{\odot}$. |
If the cluster age were approximately 15 Myr. the derived LIAM vatios would ος consistent with the Iroupa IME: an independent determination ¢ft the clusters age could rule out this possibility. | If the cluster age were approximately 15 Myr, the derived $L/M$ ratios would be consistent with the Kroupa IMF; an independent determination of the cluster's age could rule out this possibility. |
Ii our analysis we assunue a pariculax form for the IMF ιο, Ixxoupa or Salpeter power avs) and inodifv the lower-nass cutoff to eeucrate the observed L/M ratio at the adopted age. | In our analysis we assume a particular form for the IMF (i.e., Kroupa or Salpeter power laws) and modify the lower-mass cutoff to generate the observed $L/M$ ratio at the adopted age. |
This method allows us to test the top-heavy IME hypothesis 1| fhe observed £/AL ratio is fit with au IMF extending «own to 0.1 M. as in cluster AICC-9 in ον there is no need to invoke an abnormal IME. | This method allows us to test the top-heavy IMF hypothesis — if the observed $L/M$ ratio is fit with an IMF extending down to 0.1 $_{\odot}$, as in cluster MGG-9 in \citet{mccrady03}, there is no need to invoke an abnormal IMF. |
Iun the present case. a torial (Ikroupa) IATF cannot explain the relatively large L/AM ratio of the cluster. | In the present case, a normal (Kroupa) IMF cannot explain the relatively large $L/M$ ratio of the cluster. |
An elevated lowerauass cutoff is not required. however. | An elevated lower-mass cutoff is not required, however. |
The observed L/M ratios could )o caused at the clusters age by flattening the slope of t16 IME. which would. change the relative proportions of stellar masses. | The observed $L/M$ ratios could be caused at the cluster's age by flattening the slope of the IMF, which would change the relative proportions of stellar masses. |
Although our cata cο not distinguish between changes to the lower mass cutoff or IME slope. it is evident that the IAIF is different than he Iroupa form. and different frou, other nearby SSCs in the M82 nuclear starburst. | Although our data do not distinguish between changes to the lower mass cutoff or IMF slope, it is evident that the IMF is different than the Kroupa form, and different from other nearby SSCs in the M82 nuclear starburst. |
Q.2 dex. compared to those with low SSFR at fixed stellar mass. | 0.2 dex, compared to those with low SSFR at fixed stellar mass. |
The fact that their metallicities are estimated by a comparable but independent method (see Kewley&Dopita(2002). for details). further strengthens our result. | The fact that their metallicities are estimated by a comparable but independent method (see \citet{kd} for details), further strengthens our result. |
Around of all red star-forming galaxies are classified as star-forming or composite galaxies from their spectra. | Around of all red star-forming galaxies are classified as star-forming or composite galaxies from their spectra. |
This is in agreement with the results of Graves.Faber.&Schiavon(2009) who also find ~30% of nearby red sequence galaxies from SDSS show presence of emission lines in their spectrum. | This is in agreement with the results of \citet{gfs} who also find $\sim$ of nearby red sequence galaxies from SDSS show presence of emission lines in their spectrum. |
They analyse LINER like emission in these galaxies to show that they are systematically younger by 2.5-3 Gyrs than their quiescent counterparts without emission at fixed rotational velocity. | They analyse LINER like emission in these galaxies to show that they are systematically younger by 2.5-3 Gyrs than their quiescent counterparts without emission at fixed rotational velocity. |
At >~0.1. Tremontietal.(2004) have shown that the metallicity is well correlated with the stellar mass of the galaxy. | At $z\!\sim\!0.1$, \citet{tremonti} have shown that the metallicity is well correlated with the stellar mass of the galaxy. |
These quantities correlate well for 8.S<log Μ΄ 10.5 M.. but the relation flattens thereafter. | These quantities correlate well for $\leq$ log $^{*}\leq$ 10.5 $_\odot$ , but the relation flattens thereafter. |
The fact that 60% of our red galaxies have log M* in the range 10.6- M. supports the hypothesis that a signiticant fraction of the red star-forming galaxies obtain their red colours from the residual older generations of stars. | The fact that $\geq$ of our red galaxies have log $^*$ in the range 10.6-10.9 $_\odot$ supports the hypothesis that a significant fraction of the red star-forming galaxies obtain their red colours from the residual older generations of stars. |
We have also shown that ~S0% of all the red star-forming galaxies are likely to be attenuated by 70.2 mag in the SDSS :- (Fig. 4). | We have also shown that $\sim$ of all the red star-forming galaxies are likely to be attenuated by $>$ 0.2 mag in the SDSS $z$ -band (Fig. \ref{dust}) ), |
and have bimodal Hs EW and D,, 4000 distributions (Figs. 5..6 | and have bimodal $_\delta$ EW and $_n$ 4000 distributions (Figs. \ref{hd}, |
and 7)) which are very different from that of the typical red sequence galaxies. | \ref{d4000} and \ref{hd-d4}) ) which are very different from that of the typical red sequence galaxies. |
As discussed in refanalysis! and refanalysis2.. the red star-forming class of galaxies remain bimodal in the distributions of A.. Hs and D,,4000. even if the selection criterion based on SFR/M® is changed to a significantly different value (e.g. ~15-percentile. corresponding to log SFR/M” = -10 7). | As discussed in \\ref{analysis1} and \\ref{analysis2}, the red star-forming class of galaxies remain bimodal in the distributions of $_z$, $_\delta$ and $_n$ 4000, even if the selection criterion based on $^*$ is changed to a significantly different value (e.g. $\sim$ 15-percentile, corresponding to log $^*$ = -10 $^{-1}$ ). |
This clearly indicates that the component of red star-forming galaxies that are similar in properties to the red sequence galaxies are not merely a result of scatter from the red sequence (quadrant lI: Fig. 2)). Butcher&Oe | This clearly indicates that the component of red star-forming galaxies that are similar in properties to the red sequence galaxies are not merely a result of scatter from the red sequence (quadrant 1; Fig. \ref{ssf-gr}) ). |
mler(1984). found that clusters at moderate to high redshift contain an exeess of blue galaxies. compared to their low-redshift counterparts. | \citet{bo} found that clusters at moderate to high redshift contain an excess of blue galaxies, compared to their low-redshift counterparts. |
For the redshift regime of our sample ἐς~ 0.1. they found a uniform blue fraction f, of ~0.03 in all clusters within Κο. the radius containing of the cluster's red sequence population. | For the redshift regime of our sample $z\!\sim\!0.1$ ), they found a uniform blue fraction $_b$ of $\sim$ 0.03 in all clusters within $_{30}$, the radius containing of the cluster's red sequence population. |
This exercise has been repeated several times in recent years for many samples of clusters (Butcher&Oemlerorisetal. 2004). | This exercise has been repeated several times in recent years for many samples of clusters \citep{bo,m3,e2,d3}. |
. Most studies of this kind detine blue’ galaxies in erms of a broadband colour that is bluer by 0.2 mag than that of he cluster's red sequence galaxies. | Most studies of this kind define `blue' galaxies in terms of a broadband colour that is bluer by 0.2 mag than that of the cluster's red sequence galaxies. |
From the above discussion. it is clear that this cut leaves out a significant fraction of star-forming galaxies that have redder broadband colours. and classifies some j»ussively evolving galaxies as ‘star-forming’. | From the above discussion, it is clear that this cut leaves out a significant fraction of star-forming galaxies that have redder broadband colours, and classifies some passively evolving galaxies as `star-forming'. |
Margoniner&Carvalho(2000). found the blue fraction of galaxies to be f,20.03-70.09. similar to Butcher&Oemler (1984).. or 44 Abell clusters at 0.03::2::0.38 within a fixed cluster-centric distance of 0.7 Mpc. which translates into the mean Ro for most Abell clusters below 2=0.1. | \citet{m3} found the blue fraction of galaxies to be $_b$ $\pm$ 0.09, similar to \citet{bo}, , for 44 Abell clusters at $\leq$ $\leq$ 0.38 within a fixed cluster-centric distance of 0.7 Mpc, which translates into the mean $_{30}$ for most Abell clusters below $z\lesssim 0.1$. |
However. Margonineretal.(2001) measured the fraction to be (1.24-E0.07).—0.01 for clusters at >=0.25. again for galaxies within a fixed cluster-centric aperture of 0.7 Mpe. | However, \citet{m1} measured the fraction to be $(1.24\pm 0.07)z-0.01$ for clusters at $z\!\leq\!
0.25$, again for galaxies within a fixed cluster-centric aperture of 0.7 Mpc. |
Elsewhere (Ellingsonetal.2001:DePropris 2004)3.. various multiples of cluster-centrie radius. sealed with rooo. have been used to show the effect of chosen aperture size on the blue fraction. | Elsewhere \citep{e2,d3}, various multiples of cluster-centric radius, scaled with $_{200}$, have been used to show the effect of chosen aperture size on the blue fraction. |
Ellingsonetal.(2001) suggest that the origin of the Butcher-Oemler effect lies in the fact that the relative fraction of "blue" galaxies on the outskirts of the clusters at higher redshifts is higher than in their local counterparts. | \citet{e2}
suggest that the origin of the Butcher-Oemler effect lies in the fact that the relative fraction of `blue' galaxies on the outskirts of the clusters at higher redshifts is higher than in their local counterparts. |
If confirmed. this result ean have a eritical impact on the study ofevolution of galaxy properties. | If confirmed, this result can have a critical impact on the study of evolution of galaxy properties. |
It is worth mentioning here that even though most of these studies use optical spectra for most or all of their sample (except for Butcher&Oemler(1984))). the spectroscopic information is used only for assigning cluster membership. | It is worth mentioning here that even though most of these studies use optical spectra for most or all of their sample (except for \citet{bo}) ), the spectroscopic information is used only for assigning cluster membership. |
Recently. with the increasing availability of large multi-wavelength datasets. the possibility of the Buteher-Oemler effect being observed in other fundamental galaxy properties (e.g.shology.Gotoetal.2003b).. or in non-optical data (e.g.mid-IR.Saintonge.Tran&Holden 2008).. is being explored. | Recently, with the increasing availability of large multi-wavelength datasets, the possibility of the Butcher-Oemler effect being observed in other fundamental galaxy properties \citep[\eg~morphology,][]{goto03b}, or in non-optical data \citep[\eg~mid-IR, ][]{saintonge}, is being explored. |
The study of Wolf.Gray&Meisenheimer(2005) unraveled the oesence of young stars and dust in the red sequence galaxies in a cluster pair at z—0.17. while Bildfelletal.(2008). find that he colour profile of of the brightest cluster galaxies (BCOs) urn bluer by 0.5—1.0 mag towards their centres. compared to the average colours seen in the cluster's red sequence. | The study of \citet{wolf} unraveled the presence of young stars and dust in the red sequence galaxies in a cluster pair at $\sim$ 0.17, while \citet{bildfell} find that the colour profile of of the brightest cluster galaxies (BCGs) turn bluer by 0.5–1.0 mag towards their centres, compared to the average colours seen in the cluster's red sequence. |
All of this ws serious implications for Butcher-Oemler like effects based on global photometric properties. | All of this has serious implications for Butcher-Oemler like effects based on global photometric properties. |
Our sample provides some additional insight to this discussion. | Our sample provides some additional insight to this discussion. |
The presence of metals. and of colour gradients relatec to SFR in a galaxy. will affect the estimated fraction of “star-forming” galaxies in low redshift clusters. | The presence of metals, and of colour gradients related to SFR in a galaxy, will affect the estimated fraction of `star-forming' galaxies in low redshift clusters. |
The fact that nearby cluster galaxies are more metal-rich than those in high-redshif clusters. will artificially enhance the fraction of passive galaxies that are selected solely on the basis of their colour. | The fact that nearby cluster galaxies are more metal-rich than those in high-redshift clusters, will artificially enhance the fraction of passive galaxies that are selected solely on the basis of their colour. |
Such an effec is usually not accounted for while comparing the blue/star-forming galaxy fraction in clusters at different redshifts. | Such an effect is usually not accounted for while comparing the blue/star-forming galaxy fraction in clusters at different redshifts. |
A quantitative analysis of this effect requires a consistent study of an unbiasec sample of mutually comparable galaxy clusters spanning a wide redshift range. for which both photometric and spectroscopic multi-wavelength data is available. | A quantitative analysis of this effect requires a consistent study of an unbiased sample of mutually comparable galaxy clusters spanning a wide redshift range, for which both photometric and spectroscopic multi-wavelength data is available. |
We leave this for a later consideration. | We leave this for a later consideration. |
We divide a sample of 76.000 galaxies found in or near rich Abell clusters (z<0.12) in the SDSS spectroscopic catalogue. into four populations. defined by placing simple limits in (gο) colour (derived from photometry) and specitiestar formation rate (derived from the detailed modelling of optical spectra). | We divide a sample of $>$ 6,000 galaxies found in or near rich Abell clusters $\leq$ 0.12) in the SDSS spectroscopic catalogue, into four populations, defined by placing simple limits in $(g-r)^{0.1}$ colour (derived from photometry) and specificstar formation rate (derived from the detailed modelling of optical spectra). |
While most blue galaxies have evidence of star formation. and red galaxies do not. this identifies two significant populations of blue passive galaxies and red star-forming galaxies. | While most blue galaxies have evidence of star formation, and red galaxies do not, this identifies two significant populations of blue passive galaxies and red star-forming galaxies. |
We trace the spectroscopic and photometric properties of galaxies in these four sub-samples in detail. using data available from SDSS DR4. | We trace the spectroscopic and photometric properties of galaxies in these four sub-samples in detail, using data available from SDSS DR4. |
The main results of our analysisean be summarised as follows: | The main results of our analysiscan be summarised as follows: |
reflection is required. the inclination dependence of the iron line properties will be important for disentangling the enhancement mechanisms. | reflection is required, the inclination dependence of the iron line properties will be important for disentangling the enhancement mechanisms. |
As shown in this work. if source motion is responsible or the enhancement. the EW of the line will decrease signicantly as one considers sources at highere inclination (note that the inclination of the disk can be measured robustly rom the iron line profile). | As shown in this work, if source motion is responsible for the enhancement, the EW of the line will decrease significantly as one considers sources at higher inclination (note that the inclination of the disk can be measured robustly from the iron line profile). |
This cllect will be much stronger ian the inclination. dependence of the line just based on imb-darkening (George Fabian 1901). | This effect will be much stronger than the inclination dependence of the line just based on limb-darkening (George Fabian 1991). |
A careful analysis of existing. datasets might. allow suc. ra trend to be addressed. | A careful analysis of existing datasets might allow such a trend to be addressed. |
CSR. gratefully acknowledges support from the National Science Foundation under grant ÀST9529175. | CSR gratefully acknowledges support from the National Science Foundation under grant AST9529175. |
ACT thanks the Roval Society for support. | ACF thanks the Royal Society for support. |
21 76% (Fuji .A. and the “COSC coincidence” problem(Steinhar | $24\%$ $76\% $ \citep{1}, $,\Lambda,$ and the "cosmic coincidence" problem\citep{2}. |
dt Observational data indicates that O4=0.763 aud ο=0.237. so large value of O4 obviously predicts that the uuiverse is accelerating today. rather than decelerating as had long been believed. | Observational data indicates that $\Omega_{\Lambda}=0.763$ and $\Omega_{m}=0.237$, so large value of $\Omega_{\Lambda}$ obviously predicts that the universe is accelerating today, rather than decelerating as had long been believed. |
The observation evidence tells us that rate of expansion m the high-z region is slower than that in our neigliborhood. | The observation evidence tells us that rate of expansion in the high-z region is slower than that in our neighborhood. |
Iu his condition. where as variation of the pa with respect to the time is equa to zero. this is provide a problem in cosinoloeyv. called fine tauniug. the quiutessence (cosmon-fBeld) solved this problem. bv using coupling between scalar feld and dar matter (FujiaucNish-loka1981:Weoiulvere 1989). | In this condition, where as variation of the $\rho _{\Lambda }$ with respect to the time is equal to zero, this is provide a problem in cosmology, called fine tanning, the quintessence (cosmon-field) solved this problem, by using coupling between scalar field and dark matter \citep{3b,3}. |
. There are several cliffcrent theories. which rave been proposed by people. to interpret the accelerating universe. such as. holograjhic DE model (Li2001:Setare20072:IriaudFeli 2011).. ageerajhic DE models (Cai2007:WeiaudCai2008:Irunet aud scalar field models of DE. which incluiue quintessence field (NojiriandOdinSOV. 2003).. quiutui field (Elazaldeetal.2001:Fenge 2005).. oiautoui field (Caldwellctal.2003) aud many others. | There are several different theories, which have been proposed by people, to interpret the accelerating universe, such as, holographic DE model \citep{4,41a,4b,41d}, agegraphic DE models \citep{5,a,b,c} and scalar field models of DE, which including quintessence field \citep{61}, quintum field \citep{8a,8}, phantom field \citep{6} and many others. |
While the quantity of cosmological coustaut is ron zero. the DE component is more generally uoceled as quiutesscwee mnechauisin. | While the quantity of cosmological constant is non zero, the DE component is more generally modeled as quintessence mechanism. |
I is a scalar field rolling down a fia potential. | It is a scalar field rolling down a flat potential. |
Iu quintessence 1uechauisui. he field ju Uuceative pressure aud therefore acts o accelerate expansion (Weimberg0051. | In quintessence mechanism, the field has negative pressure and therefore acts to accelerate expansion \citep{10}. |
Ikhouryv aud Wellan {2003) have introduced another ια nechamisin. which called chameleou iechanisi. | Khoury and Weltman (2003) have introduced another kind mechanism, which called chameleon mechanism. |
Iu lis mechamisin the scalar field acquire a mass whose naguitude depexls ou the local matter density (Braxetal.2 XL. | In this mechanism the scalar field acquire a mass whose magnitude depends on the local matter density \citep{11}. |
.. Also it is a παν to related an effective nass for scalar field o. | Also it is a way to related an effective mass for scalar field $\phi$. |
Scalar field is expausion field. and cun © obtained from string theory (Orti/u.2001:001)... Also the chameleon niechauisui Is a wav to eive an effective mass to a Πο scalar field via field sef interaction and interaction between | Scalar field is expansion field, and can be obtained from string theory \citep{12,12a}.. Also the chameleon mechanism is a way to give an effective mass to a light scalar field via field self interaction and interaction between |
(INomatsuetal.2009).. (Springeletal.2005).. (Bertone | \citep{art-Komatsuetal2009}, \citep{art-Springel2005Natur.435..629S}. \citep{rev-BertoneHS2005}. |
etal.2005).. 10.770? 10.om? (Ahmedetal.2009).. | \citep{art-Bernabeietal2010AIPC.1223...50B,art-CDMSII2010science,art-Aalseth2010,art-Fitzpatrick2010}
$10^{-32} {\rm cm^2}$ \citep{art-XENON10_SD2008,art-COUPP2008}) $10^{-40} {\rm cm}^2$ \citep{art-CDMSII_SI2009}. |
(Dergstrón)000) Sarkar2010) iuportaut complement to ¢lect detection searches for ight WIAMPs. | \citep{art-Bergstrom2009NJPh} \citep{art-Frandsen2010PhRvL} important complement to direct detection searches for light WIMPs. |
These particles are trapped iu the Suus interior when they collide with nuclei aud lose (linear) nonientuin. and drift iuto the Suus core. | These particles are trapped in the Sun's interior when they collide with nuclei and lose (linear) momentum, and drift into the Sun's core. |
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