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An obvious candidate are blazars. where the dust extinetion could be derived either as done here from Galactic dust maps. or from reddening of the optical power-law itself (where the caveats are that the o0ata in different bands be taken simultaneously and that one =eeds to account for intrinsic curvature of the spectra). | An obvious candidate are blazars, where the dust extinction could be derived either as done here from Galactic dust maps, or from reddening of the optical power-law itself (where the caveats are that the data in different bands be taken simultaneously and that one needs to account for intrinsic curvature of the spectra). |
Indeed. for a smaller sample. the X-ray and some of the optical-UV data already exists. | Indeed, for a smaller sample, the X-ray and some of the optical-UV data already exists. |
Blazars night be bright enough for UV spectroscopy as well. allowing H» and Ένα measurements to be made. | Blazars might be bright enough for UV spectroscopy as well, allowing $_2$ and $\alpha$ measurements to be made. |
Such a programme. while observationally intensive for a large enough sample to probe a significant fraction of the Galaxy. could represent one of the best probes of the Galactic ISM vet devised. | Such a programme, while observationally intensive for a large enough sample to probe a significant fraction of the Galaxy, could represent one of the best probes of the Galactic ISM yet devised. |
Studies of metals-to-dust and gas-to-dust also exist outside the local group. | Studies of metals-to-dust and gas-to-dust also exist outside the local group. |
Metals-to-dust ratios from foreground lensing galaxies at zx| have been obtained using multiply-imaged quasars (??).. | Metals-to-dust ratios from foreground lensing galaxies at $z\lesssim1$ have been obtained using multiply-imaged quasars \citep{2006ApJ...637...53D, 2009ApJ...692..677D}. |
Results show metals-to-dust ratios consistent with Galactic values. | Results show metals-to-dust ratios consistent with Galactic values. |
Those objects are typically relatively high mass. evolved systems. | Those objects are typically relatively high mass, evolved systems. |
? showed in an analysis of GRB afterelows that GRB host galaxies. which are typically young- and star-forming with low metallicities and hard radiation environments (?????2?).. have substantially lower dust-to-gas ratios than the local group even after accounting for metallicity. | \citet{2011arXiv1102.1469Z} showed in an analysis of GRB afterglows that GRB host galaxies, which are typically young and star-forming with low metallicities and hard radiation environments \citep{2004A&A...425..913C,2006ApJ...653L..85C,2009ApJ...691..182S,2010MNRAS.405...57S,2002ApJ...566..229C,2008ApJ...683..321F,2010arXiv1010.1783W}, have substantially lower dust-to-gas ratios than the local group even after accounting for metallicity. |
An analysis of the dust-to-metals ratio and metallicity of sightlines through the Galaxy has been presented. | An analysis of the dust-to-metals ratio and metallicity of sightlines through the Galaxy has been presented. |
The Galactic metal column densities were determined using the lower bound of the distribution of soft X-ray absorptions of the afterglows of a large sample of GRBs detected by the satellite. | The Galactic metal column densities were determined using the lower bound of the distribution of soft X-ray absorptions of the afterglows of a large sample of GRBs detected by the satellite. |
The corresponding extinction and gas column densities were found using the dust and maps of ?. and ? respectively. | The corresponding extinction and gas column densities were found using the dust and maps of \citet{1998ApJ...500..525S} and \citet{2005A&A...440..775K} respectively. |
The metal to atomic hydrogen relation is well reproduced with a metallicity ~1.75 times the solar metallicity of ?.. | The metal to atomic hydrogen relation is well reproduced with a metallicity $\sim1.75$ times the solar metallicity of \citet{2009ARA&A..47..481A}. |
The best-fitting relation between metal and dust column densities is Ny/Ay2.2.245x107! cem mmagé! (using AG89 abundances). | The best-fitting relation between metal and dust column densities is $N_{H_X}/A_V = 2.2_{-0.3}^{+0.4}\times10^{21}$ $^{-2}$ $^{-1}$ (using AG89 abundances). |
Previous observations are consistent with this result. suggesting that the metallicity for a typical ISM sightline is ddex higher than the current best value for the solar metallicity. | Previous observations are consistent with this result, suggesting that the metallicity for a typical ISM sightline is dex higher than the current best value for the solar metallicity. |
It is therefore suggested that a better reproduction of the Galactic soft X-ray absorption will be provided with a metallicity ~20% than the solar metallicity of AGS89, | It is therefore suggested that a better reproduction of the Galactic soft X-ray absorption will be provided with a metallicity $\sim20\%$ than the solar metallicity of AG89. |
However. it is also found that a linear representation does not reproduce the gas-to-dust relationship very well. | However, it is also found that a linear representation does not reproduce the gas-to-dust relationship very well. |
A gas-to-dust relationship with Ny,=2.00x0.14AT77" provides a much better fit to the data. | A gas-to-dust relationship with $N_\ion{H}{i} = 2.00\pm0.14\times A_V^{0.77\pm0.07}$ provides a much better fit to the data. |
It is very likely that this ts predominantly a metallicity-gradient effect. and it is therefore concluded that while the gas-to-dust relation may be as given above. the best proxy for the Galactic soft X-ray absorption should be given by the dust column density with a relation of Ny./Av=2.2x107! cem mmag! or cem mmag7! for an Ry=3.1. | It is very likely that this is predominantly a metallicity-gradient effect, and it is therefore concluded that while the gas-to-dust relation may be as given above, the best proxy for the Galactic soft X-ray absorption should be given by the dust column density with a relation of $N_{\rm H_X}/A_V =
2.2\times10^{21}$ $^{-2}$ $^{-1}$ or $N_{\rm H_X}/E(B-V) =
6.8\times10^{21}$ $^{-2}$ $^{-1}$ for an $_V=3.1$ . |
) and MeII (2800) lines in the IUE spectrum support its classification as a hot post-ACD star. | ) and MgII (2800) lines in the IUE spectrum support its classification as a hot post-AGB star. |
Using solar metallicitv Iurnucz (1991) models we doetermuue Tac19000 + 1000I aud log ο = 2.5 £ 0.5. | Using solar metallicity Kurucz (1994) models we determined $_{\rm eff}$ =19000 $\pm$ 1000K and log g = 2.5 $\pm$ 0.5. |
ZIIRÀS 22025|5219 (LSIIE Its IRAS colors were found to be similar to PNe iux Parthasarathy et al. ( | IRAS 22023+5249 (LSIII Its IRAS colors were found to be similar to PNe and Parthasarathy et al. ( |
2001) classified it as a hot post-AGB star. | 2001) classified it as a hot post-AGB star. |
The photometry of the star was obtained frou the Twcho-2 Catalogue (Που ct al. | The photometry of the star was obtained from the Tycho-2 Catalogue (Hog et al., |
2000). | 2000). |
It is Inter in Wackerlings (1970) catalog of carly-type euission-liune y.ars. | It is listed in Wackerling's (1970) catalog of early-type emission-line stars. |
The D spectral type of the star was obtained from the Simbad database. | The B spectral type of the star was obtained from the Simbad database. |
We compared the UV spectra of this star with that of a D2I standard star (ITD11117) aud found it to be simular. | We compared the UV spectrum of this star with that of a B2I standard star (HD41117) and found it to be similar. |
The Fa;/F.; (=3.03) ratio suggests obscuration of the hot ceutral star and lends support to the modeled circumstellar extinction value of QUE | The $_{\rm fir}$ $_{\rm star}$ (=3.03) ratio suggests obscuration of the hot central star and lends support to the modelled circumstellar extinction value of 44. |
There was uo difference between the 1993 and 1995 UV(IUE) spectra. | There was no difference between the 1993 and 1995 UV(IUE) spectra. |
ZIIRÀS 22195|5131 (LSIIE It is classified asa PN in the Strasboure-ESO catalogue of Galactic planetary jebulae (Acker et al. | IRAS 22495+5134 (LSIII It is classified as a PN in the Strasbourg-ESO catalogue of Galactic planetary nebulae (Acker et al., |
1992). | 1992). |
Tvleuda and Stasinska (1991) reported an augular diameter of0.5"., expausion velocity of 10 + aud V—12.08. | Tylenda and Stasinska (1994) reported an angular diameter of, expansion velocity of 10 $^{-1}$ and V=12.08. |
The V auaenitude from the Twcho-2 Catalogue (Που et al. | The V magnitude from the Tycho-2 Catalogue (Hog et al., |
2000) is 11778. | 2000) is 78. |
Haudler (1999) found it to be variable with amplitude variations of 33i the Jolson V banc. | Handler (1999) found it to be variable with amplitude variations of 3 in the Johnson V band. |
Wlüle the longc» term variations (several davs) were non periodic. the short terii variations were quasi-periodic with tine scales of ether 8.9 or LL3 hours. | While the long term variations (several days) were non periodic, the short term variations were quasi-periodic with time scales of either 8.9 or 14.3 hours. |
Variations in the stellar iass-loss coupled with stellar pulsation may explain the observed long aud short-term variability. | Variations in the stellar mass-loss coupled with stellar pulsation may explain the observed long and short-term variability. |
From the Wipparchos catalog. the Ilipparchos maguitucdes atf iunaxinmna ancl nuuinuun are 129111 aud 17 respectively. | From the Hipparchos catalog, the Hipparchos magnitudes at maximum and minimum are 11 and 47 respectively. |
Frou: the logarithiuic extinction at I> using radio flux at δέ (cHO.11). Waler(1983) obtained Vy=0M228. | From the logarithmic extinction at $\beta$ using radio flux at 5Gz (c=0.41), Kaler(1983) obtained 28. |
Using the Πα το ILJ ratio (¢-=0.55). Tvleuda et al. ( | Using the $\alpha$ to $\beta$ ratio (c=0.55), Tylenda et al. ( |
1992) obtained E(B-Vj=0.37. | 1992) obtained E(B-V)=0.37. |
These values are in eood agreement with the imterstellar E(B-V)=0.33 derived from the feature in the UV. | These values are in good agreement with the interstellar E(B-V)=0.33 derived from the feature in the UV. |
We analvsed the UV(IUE) spectra of 15 hot post-ACD candidates. | We analysed the UV(IUE) spectra of 15 hot post-AGB candidates. |
Iu 11 cases (IRAS13266-5551 (CPD-55 5588). IRAÁSII331-6135. (IIlen3-1013). TRASI6206-5956 (SAO 213756). IRASL707Lis15 (Ion3-1217). IRÁSITT311-1921 (IIe2-1128). IRASI7123-1755. (Ποτο). TRASLS023-3109 (LSS 1631). IRASIS062]2110. (SAO 857606). IRAS18371-3159. (LSE 63). TRAS22023(5219 (LSIII 15221) aud TRAS22195|5131 ΠΠ 15112)). the UV. spectra revealed obscuration of the hot central stars due to cireunistellar dust. | In 11 cases (IRAS13266-5551 (CPD-55 5588), IRAS14331-6435 (Hen3-1013), IRAS16206-5956 (SAO 243756), IRAS17074-1845 (Hen3-1347), IRAS17311-4924 (Hen3-1428), IRAS17423-1755 (Hen3-1475), IRAS18023-3409 (LSS 4634), IRAS18062+2410 (SAO 85766), IRAS18371-3159 (LSE 63), IRAS22023+5249 (LSIII +5224) and IRAS22495+5134 (LSIII +5142)), the UV spectra revealed obscuration of the hot central stars due to circumstellar dust. |
While IRÀÁS17123-1755 (ITen3-1175) was nof detected at all in a 35 nünute exposure. the UV contiuua of the remaining 10 stars were found to be considerably reddened. | While IRAS17423-1755 (Hen3-1475) was not detected at all in a 35 minute exposure, the UV continua of the remaining 10 stars were found to be considerably reddened. |
We found that the circiunstellar extinction in these 10 stars varies linearly as A3, | We found that the circumstellar extinction in these 10 stars varies linearly as $\lambda^{-1}$. |
A A| law for the circumstellar extinction was also found iu the case of the post-ACD star. IER1010 (Waters et al. | A $\lambda^{-1}$ law for the circumstellar extinction was also found in the case of the post-AGB star, HR4049 (Waters et al., |
1989, Monier Parthasarathy. 1999). | 1989, Monier Parthasarathy, 1999). |
In the coutext of Mie scattering (Spitzer. 19758). linear extinction arises frou dust eraius simall compared to the wavelength of light. | In the context of Mie scattering (Spitzer, 1978), linear extinction arises from dust grains small compared to the wavelength of light. |
The shortest wavelength of light at which the extinction is linear can give an estimate of the size of the smallest erains in the ciremustellar euvironment of these stars (A = 251. where. a is the radius of the cust grain). | The shortest wavelength of light at which the extinction is linear can give an estimate of the size of the smallest grains in the circumstellar environment of these stars $\lambda$ = $\pi$ a, where, a is the radius of the dust grain). |
However our IVE SWE observations are limited to (A! x 8.71). | However our IUE SWP observations are limited to $\lambda^{-1}$ $\approx$ $\mu^{-1}$ ). |
Shortward of the spectra are noisy and often contaminated by Lyman o. | Shortward of the spectra are noisy and often contaminated by Lyman $\alpha$. |
Taking as the shortest observed waveleugth at which the extinction is linear in At. we may infer an upper nut of a oz or the radii of the simall erams. | Taking as the shortest observed wavelength at which the extinction is linear in $\lambda^{-1}$, we may infer an upper limit of a $\approx$ for the radii of the small grains. |
Waters ct al. ( | Waters et al. ( |
1989) speculate that the destruction of these grains in the vicinity of the hot ceutral stars of PPNe and PNe max eive rise to sinaller eraius aud polvaromatic hivdrocarbons (PAIIS). | 1989) speculate that the destruction of these grains in the vicinity of the hot central stars of PPNe and PNe may give rise to smaller grains and polyaromatic hydrocarbons (PAHs). |
PATI catures at 8.2. 8.6 aud 11.3 µ have been detected in the ecirciuustellaz euvironment of several post-AGB stars. PPNe and PNe (see ee. | PAH features at 8.2, 8.6 and 11.3 $\mu$ have been detected in the circumstellar environment of several post-AGB stars, PPNe and PNe (see eg. |
Beintema et al.. | Beintema et al., |
1996). | 1996). |
It would be interesting to study the iutrared spectra of our hot post-ACB candidates to know more about the chemical compositions (carhon-rich or oxvgenarich nature) of the dust eraius and the evolutiou of these eraius iu the circumstellar environment of these stars. | It would be interesting to study the infrared spectra of our hot post-AGB candidates to know more about the chemical compositions (carbon-rich or oxygen-rich nature) of the dust grains and the evolution of these grains in the circumstellar environment of these stars. |
Variation of IRAS16206-5956 (SAO 213756) and IRÁSISOG2|2110 (SAO 55766) in the UV iav be due to stellar pulsatious audor due to variable circumstellar extinction similar to that observed in the case of TTR049 (Waters et al.. | Variation of IRAS16206-5956 (SAO 243756) and IRAS18062+2410 (SAO 85766) in the UV may be due to stellar pulsations and/or due to variable circumstellar extinction similar to that observed in the case of HR4049 (Waters et al., |
1989. Monier Parthasarathy. 1999). | 1989, Monier Parthasarathy, 1999). |
Significant circuiistellar extinction was not observed in the case of 1531. IRASI7L60-3111 (SAO 209306) and IBAS18379-L707 (LSS 5112). | Significant circumstellar extinction was not observed in the case of IRAS17203-1534, IRAS17460-3114 (SAO 209306) and IRAS18379-1707 (LSS 5112). |
The effective temperatures aud eravitics of these three stars were estinutfed using ERurucz model atmospheres. | The effective temperatures and gravities of these three stars were estimated using Kurucz model atmospheres. |
Fay Feray 1.0 iu the case of IRAST1331-6135 (IIcus-1013). IRAST7211-192| (Ten3-1128). IRASI17123-1755 (Houw1175). IRÁS18062|2110 (SAO 85766). aud TRAS22023|5219 (LSIII 15221) indicates the presence of dusty disks around these stars. | $_{\rm fir}$ $_{\rm star}$ $>>$ 1.0 in the case of IRAS14331-6435 (Hen3-1013), IRAS17311-4924 (Hen3-1428), IRAS17423-1755 (Hen3-1475), IRAS18062+2410 (SAO 85766) and IRAS22023+5249 (LSIII +5224) indicates the presence of dusty disks around these stars. |
From the UV(IUE) spectra we found that 7 (IRASI258 Ες (CPD-55 5588). TRASL7203-1531. TRASITT311- (Heuc-1128).. IRASIL71GO-311L (SAO 209306), IRASISO23-3109 (LSS 1631) and IRAS22023|5219 (LSI 15221)) of the 15 hot post-AGD candidates have stellar wind velocities in excess of | iudicatiug mass-loss. | From the UV(IUE) spectra we found that 7 (IRAS12584-4837, IRAS13266-5551 (CPD-55 5588), IRAS17203-1534, IRAS17311-4924 (Hen3-1428), IRAS17460-3114 (SAO 209306), IRAS18023-3409 (LSS 4634) and IRAS22023+5249 (LSIII +5224)) of the 15 hot post-AGB candidates have stellar wind velocities in excess of $^{-1}$ indicating post-AGB mass-loss. |
to 6800Ix. note the four stars with A(Li)e-3.3 between 6900 and 7200lx. It. is conceivable that these somewhat high A(Li) might somehow be related to upwards diffusion: however. the RM93-precdicted Peak is not as wide as 300K. and the Li-depleted stars in the same Ty; range would require a dillerent explanation. | to 6800K, note the four stars with $\sim$ 3.3 between 6900 and 7200K. It is conceivable that these somewhat high A(Li) might somehow be related to upwards diffusion; however, the RM93-predicted Peak is not as wide as 300K, and the Li-depleted stars in the same $T_{eff}$ range would require a different explanation. |
Finally. although it is becoming increasingly clear that slow mixing affects the radiative lavers below the SCZs of low mass stars. the challenge remains (o decipher what combination of specific physical mechanisms are al work. | Finally, although it is becoming increasingly clear that slow mixing affects the radiative layers below the SCZs of low mass stars, the challenge remains to decipher what combination of specific physical mechanisms are at work. |
We conclude that the study of the timing and morphology of the formation of of the Li eap provides a powerful new method to study the physical mechanisms occuring inside stars. | We conclude that the study of the timing and morphology of the formation of of the Li gap provides a powerful new method to study the physical mechanisms occuring inside stars. |
Results from applying this method to M35 join a variety of other evidence that points to the action ol slow mixing in low mass stars. including (ae Li-Be depletion correlation 1993).. higher Li in short period Gclally locked binaries (ΙΟΓΙΟ.Thorburn.etal.Delivannisetal.1994:Han& 1995).. moderately. rapid Li depletion in AIG? subgiants (Delivannis.King.&Doesgaard1997:SillsDelivannis2000).. the existence of a moderate De Gap (Boeseaarcd&Ixing2002) though (so far) no B Gap 1993).. the continued depletion of Li during the main sequence (Jeffries1997:Jellriesetal. 2002).. and the Li dispersions at fixed 7;;; (Thorburn et al.: 1993b)). | Results from applying this method to M35 join a variety of other evidence that points to the action of slow mixing in low mass stars, including the Li-Be depletion correlation \citep{d98}, higher Li in short period tidally locked binaries \citep[``SPTLBs'',][]{thdp, con94, rd95}, moderately rapid Li depletion in M67 subgiants \citep{d97, sd00}, the existence of a moderate Be Gap \citep{bk02} though (so far) no B Gap \citep{b98}, , the continued depletion of Li during the main sequence \citep{j97, d00, j02}, and the Li dispersions at fixed $T_{eff}$ (Thorburn et al.; \citealt{s93b}) ). |
This work has been supported by the National Science Foundation under Grants and AST-0206202. | This work has been supported by the National Science Foundation under Grants AST-9812735and AST-0206202. |
GES model that used in this study as well as the Nu-Raucall cloud. seleme that it used for cloud simulation and the ANP model that we used to derive forecast of atmospheric seeing. | GFS model that used in this study as well as the Xu-Randall cloud scheme that it used for cloud simulation and the AXP model that we used to derive forecast of atmospheric seeing. |
In Section 3. we describe the observations we used to evaluate the GES model. | In Section 3, we describe the observations we used to evaluate the GFS model. |
Section { preseuts tle details and discussions of evaluation methodology and result: while Section 5 gives the couclucdiug remarks of this study. | Section 4 presents the details and discussions of evaluation methodology and result; while Section 5 gives the concluding remarks of this study. |
The GFS inodel provides output in two erids with dillerent. spatial resolution: grkl 003 at 17x1. aud grid 00 αι 0.5°«0.57. | The GFS model provides output in two grids with different spatial resolution: grid 003 at $1^{\circ}\times1^{\circ}$, and grid 004 at $0.5^{\circ}\times0.5^{\circ}$. |
To provide best-possible forecast. we use the later in our study. | To provide best-possible forecast, we use the later in our study. |
Model outputs from the period of January. 1. 2008 to December 31. 2009 at three-hourly interval for 0«r€ 72h at 00Z initialization is retrieved from the National Operational Model Archive Distribution System (NOMADS:seeRutledgeetal.2006). [or evaluation. | Model outputs from the period of January 1, 2008 to December 31, 2009 at three-hourly interval for $0<\tau\leq72$ h at 00Z initialization is retrieved from the National Operational Model Archive Distribution System \citep[NOMADS; see][]{rut06} for evaluation. |
The GES dataset coutaius approximately L10 fields. supplying forecast fields both of geueral meteorological interests (sucli as temperature. huiidity. wind direction aud speed. etc.) | The GFS dataset contains approximately 140 fields, supplying forecast fields both of general meteorological interests (such as temperature, humidity, wind direction and speed, etc.) |
aud. [or special purpose. including cloud cover fraction ou different layer (low. mid. high. couvective. aud of total atmospheric columnu). | and for special purpose, including cloud cover fraction on different layer (low, mid, high, convective, and of total atmospheric column). |
Although the atmospheric seeing is not among the output fields. it can be derived indirectly as every of the required meteorological variables are given. | Although the atmospheric seeing is not among the output fields, it can be derived indirectly as every of the required meteorological variables are given. |
Iu the GES model. cloud cover fraction for each grid box is computed using the cloud scheme presented by Xu&Randall(1996)... which is shown as eq. [ | In the GFS model, cloud cover fraction for each grid box is computed using the cloud scheme presented by \citet{xu96}, which is shown as eq. [ |
1]. | 1]. |
In this equation. RA is the relative liumidity. q* aud q are the saturatiou specilic humidity and gene, Is a prescribed miuimuau threshold value of q.. | In this equation, $RH$ is the relative humidity, $q^{*}$ and $q_{c}$ are the saturation specific humidity and $q_{cmin}$ is a prescribed minimum threshold value of $q_{c}$. |
Dependiug ou the euviroumental temperature. q and gy, are caleulated with respect to water pliase or ice phase (Yaug. personal communication). | Depending on the environmental temperature, $q^{*}$ and $q_{c}$ are calculated with respect to water phase or ice phase (Yang, personal communication). |
Cloud cover fraction cau therefore be calculated for any layer as long as the RA. q auc g. are known aud qo; Is sultably prescribed. | Cloud cover fraction can therefore be calculated for any layer as long as the $RH$, $q^{*}$ and $q_{c}$ are known and $q_{cmin}$ is suitably prescribed. |
We uote that the calculation is done as part of the moclel simulation at NCEP. so the cloud fields are used as-is from the GES datasets. | We note that the calculation is done as part of the model simulation at NCEP, so the cloud fields are used as-is from the GFS datasets. |
The GES imoclel divides the whole atmospheric columu tuto 26 layers. | The GFS model divides the whole atmospheric column into 26 layers. |
The total cloud cover lor eutireatmospheric column is derived uider the asstunption that clouds in all layers are maximally randomly overlapped (Yangetal. 2005).. | The total cloud cover for entireatmospheric column is derived under the assumption that clouds in all layers are maximally randomly overlapped \citep{yan05}. . |
only weakly on the observed. redshift. | only weakly on the observed redshift. |
“Phe dependence arises through the angular cliameter distance (since the total SZ flux scales approximately as in the absence of instrumental beam elects). | The dependence arises through the angular diameter distance (since the total SZ flux scales approximately as $d_{\rm A}^{-2}$ in the absence of instrumental beam effects). |
The minimumαντ mass lor detection is also seen to be lower for objects with higher τι since these objects are denser and so produce larger SZ Ilux for given mass. C | The minimum mass for detection is also seen to be lower for objects with higher $z_{\rm f}$ since these objects are denser and so produce larger SZ flux for given mass. ( |
Ehe lines become horizontal at 2=zn. since no cluster can have formed after it was observed.) | The lines become horizontal at $z_{\rm f}=z_{\rm obs}$ since no cluster can have formed after it was observed.) |
In this subsection we explore the theoretical uncertainties inherent in current analytic models of SZ cluster properties. | In this subsection we explore the theoretical uncertainties inherent in current analytic models of SZ cluster properties. |
Our aim is to assess the relative importance of these uncertainties and highlight where further work is needed to produce an accurate moclel. | Our aim is to assess the relative importance of these uncertainties and highlight where further work is needed to produce an accurate model. |
Alany previous analytic studies of the SZ elfect have used the Press-Schechter mass function (Press&Schechter 1974).. but more accurate fitting formulaes are now known. | Many previous analytic studies of the SZ effect have used the Press-Schechter mass function \citep{ps74}, but more accurate fitting formulaes are now known. |
To explore the differences. intriduced. by. these alternative formulae. we have implemented both the Press-Schechter mass function. and two others. based. upon N-body simulations developed. by J2000 and. ST. | To explore the differences intriduced by these alternative formulae, we have implemented both the Press-Schechter mass function and two others based upon N-body simulations developed by J2000 and ST. |
The J2000 mass function is currentiv the best mateh to numerical simulations over a wide range of mass scales. ancl should be taken as giving the best estimate of the halo mass function. | The J2000 mass function is currently the best match to numerical simulations over a wide range of mass scales, and should be taken as giving the best estimate of the halo mass function. |
We show results for the other mass functions simply. for comparison. | We show results for the other mass functions simply for comparison. |
The results for all three mass functions for the experiment are displaved in Fig. 4.. | The results for all three mass functions for the experiment are displayed in Fig. \ref{fig:Pydep}. |
The main cillerences between the models show up in the redshift distribution of clusters. | The main differences between the models show up in the redshift distribution of clusters. |
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