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Recent observations by VLT/UVES (Asplundetal.2005). and Subaru/IIDS (Inoue el al. | Recent observations by VLT/UVES \citep{Asplund05} and Subaru/HDS (Inoue et al. |
2005. Aoki et al. in prep.) | 2005, Aoki et al, in prep.) |
have revealed that some very metal-poor stars possess surprisingly high abundances of Τα. much higher than expected [rom standard supernova. cosnic rav models (Ramatyοἱal.2000:Suzuki&Yoshii2001:Prantzos2006). | have revealed that some very metal-poor stars possess surprisingly high abundances of $^6$ Li, much higher than expected from standard supernova cosmic ray models \citep{Ramaty00,Suzuki01,Prantzos06}. |
. The measured values ave also higher than can be accommodated in the SN Ic production scenario discussed above. even in the case of an energetic explosion similar (ο SN 1995bw. as can beseen from the results of Nalunura&Shigevama (2004).. | The measured values are also higher than can be accommodated in the SN Ic production scenario discussed above, even in the case of an energetic explosion similar to SN 1998bw, as can beseen from the results of \citet{Nakamura04}. . |
However. besides spallation of C and O. the | However, besides spallation of C and O, the |
clusters have merged. alter approximately 100 Nr. | clusters have merged after approximately 100 Myr. |
The time within which half of the clusters meree increases with the CC size Crom 15 Myr lor model M.11.11.5.225 to 200 Myr for mocel 1.5.1150) aud decreases with CC mass (from 150 Myr for model 1100 to 65 Myr for model ΝΤ1100). | The time within which half of the clusters merge increases with the CC size (from 15 Myr for model 25 to 200 Myr for model 150) and decreases with CC mass (from 150 Myr for model 100 to 65 Myr for model 100). |
Τιe further time evolution of the merger object of model M_J1_11.5_1100 is plotted in Figure 6.. | The further time evolution of the merger object of model 100 is plotted in Figure \ref{fig_timeevol_M_R}. |
The encosecL mass of the merger objects is defined as the mass within 800 pc. | The enclosed mass of the merger objects is defined as the mass within 800 pc. |
The hal(-uiass raclius is the radius of the sphere within which half of the mass is enclosed. | The half-mass radius is the radius of the sphere within which half of the mass is enclosed. |
However. as observers can ouly cle‘ive a projected ballimass radius. we calculate also a value projected ou the sky defined as the »ojected radius within which ball of the mass is included. | However, as observers can only derive a projected half-mass radius, we calculate also a value projected on the sky defined as the projected radius within which half of the mass is included. |
The projected half-iuass radius is slightlyΙ snaller than the three-dimensioual half-inass radius (Table 3)) and. correspouds to the observed. efective radius. reg. i£ nass follows light. | The projected half-mass radius is slightly smaller than the three-dimensional half-mass radius (Table \ref{tbl-2}) ) and corresponds to the observed effective radius, $r_{\rm eff}$, if mass follows light. |
The effective radius and mass of the merger object decrease aud become fairly constaut alter about 7 Gyr. | The effective radius and mass of the merger object decrease and become fairly constant after about 7 Gyr. |
The structural parameters chauge ouly very slightly in the nest few Cir. | The structural parameters change only very slightly in the next few Gyr. |
We consider models with three CC masses. Af = 1.0. 1.5. and 2.0 «10° ML. and six CC Plummer r:cdi )etweell. ICC 25 ald 150 pe in steps of 25 pe to analyze the depeudence of the structure Jalaljeters of the merger oljects on the initjal CC mass ancl size. | We consider models with three CC masses, $M^{\rm CC}$ = 1.0, 1.5, and 2.0 $\times 10^{6}$ $_{\odot}$, and six CC Plummer radii between $R_{\rm pl}^{\rm CC}$ = 25 and 150 pc in steps of 25 pc to analyze the dependence of the structural parameters of the merger objects on the initial CC mass and size. |
In addition. two moclels with i = 2Q0 and 300 pe are calcuated for ACC 2.0 x109. | In addition, two models with $R_{\rm pl}^{\rm CC}$ = 200 and 300 pc are calculated for $M^{\rm CC}$ = 2.0 $\times~10^{6}$. |
All models have |e sale relative initial «‘ibutiou of star clusters. wlich is scaled according to the respective PIuimer radii (see Figure laa and b). | All models have the same relative initial distribution of star clusters, which is scaled according to the respective Plummer radii (see Figure \ref{figinimodel}a a and b). |
The velocities ofthe individual star clusters are also scaled withit CC mass aud size to keej» the CCs initially iu viria equilibrium. | The velocities of the individual star clusters are also scaled with CC mass and size to keep the CCs initially in virial equilibrium. |
For tal moclels he mereingeOm process leads to a stable object. | For all models the merging process leads to a stable object. |
The number of me'ged star clusters is between 10 ancl 20. | The number of merged star clusters is between 10 and 20. |
The properties of tlie merger objects atthe current. position of 22119 are cisplaved in Tae +)3.. | The properties of the merger objects atthe current position of 2419 are displayed in Table \ref{tbl-2}. |
Figure Jaa shows the eneosecL iuass. Alene. of tie merger objects as a function of the initial CC Plununer raditJf. | Figure \ref{figmassreff}a a shows the enclosed mass, $M_{\rm encl}$ , of the merger objects as a function of the initial CC Plummer radius, $R_{\rm pl}^{\rm CC}$. |
T1ο fraction of the initlal CC mass. whitch is bouud to the mereer object. deeeases amost linear with increasing CC size from about al ReyCC! = 25 pe for all Imoiree masses {ὁ valtes of 16.53. aud at "en 150 pe for CC masses of ALCO = 1.0. 15. and οκ108. 'espectively. | The fraction of the initial CC mass, which is bound to the merger object, decreases almost linear with increasing CC size from about at $R_{\rm pl}^{\rm CC}$ = 25 pc for all three masses to values of 46, 53, and at $R_{\rm pl}^{\rm CC}$ = 150 pc for CC masses of $M^{\rm CC}$ = 1.0, 1.5, and 2.0 $\times~10^{6}$, respectively. |
]ass-oss OCcurs either by thje. escape of individual 5ars from the diffuse 3.ellar com»)o0nent. whicl1 builcIs up curing the mergiug process (see Sect. 3.1)). | Mass-loss occurs either by the escape of individual stars from the diffuse stellar component, which builds up during the merging process (see Sect. \ref{timeevolution}) ), |
or by entire star ‘lusters escapiug the me'elug jyocess. | or by entire star clusters escaping the merging process. |
For all three 1asses. all 20 star clusters merge for compact --uodels up ο Reyδς — 5 dC. Le. mass-loss is ouly from the ciffuse component or these moclels. | For all three masses, all 20 star clusters merge for compact models up to $R_{\rm pl}^{\rm CC} $ = 75 pc, i.e. mass-loss is only from the diffuse component for these models. |
For je more extended models with /umμα.> LOO pe star clusters escape froin tlie uerpging process. | For the more extended models with $R_{\rm pl}^{\rm CC} \ge$ 100 pc star clusters escape from the merging process. |
For extended models with /umCC! = 150 pe six. live. aud two starclusters escape fo CC masses of ACC = 1.0. 1.5. and 2.0 x 10°. respectively. | For extended models with $R_{\rm pl}^{\rm CC}$ = 150 pc six, five, and two starclusters escape for CC masses of $M^{\rm CC}$ = 1.0, 1.5, and 2.0 $\times~10^{6}$ , respectively. |
Oue of the inost important eoals in modern astroplivsics is to uuderstaud the end of the cosmic dark ages. when the first stars aud galaxies transformed the simple carly uuiverse into a state of ever mucreasig colmplexity (Brouun&Larson2001:ChiudiFerrara 2009). | One of the most important goals in modern astrophysics is to understand the end of the cosmic dark ages, when the first stars and galaxies transformed the simple early universe into a state of ever increasing complexity \citep{bl04a, cf05, glover05, bl07, bromm09}. |
. The first stars. the so-called Population OI (Pop TT). were the source of livdrogen-ionizing UV photons. thus initiating the exteuded process of cosmic relonization (Fanetal.2006). | The first stars, the so-called Population III (Pop III), were the source of hydrogen-ionizing UV photons, thus initiating the extended process of cosmic reionization \citep{fck06}. |
. They also svuthesized the first heavy chemical elements. bevoud the hydrogen aud helium produced in the big bang. to be dispersed iuto the pristine iutergalactie medimm (ICM) through superuova (SN) explosions aud winds (οι,Madanetal.2001:&Abel 2008).. | They also synthesized the first heavy chemical elements, beyond the hydrogen and helium produced in the big bang, to be dispersed into the pristine intergalactic medium (IGM) through supernova (SN) explosions and winds \citep[e.g.,][]{mfr01,wa08b}. |
Au intriguing possibility for the first stars is that some of them may have died as a pairdustabilitv superuova (PISN). a peculiar fate predicted for progenitor niasses in the range 110AL.XM.&260 (Barkatctal.1967:Teeer&Woosley 2002). | An intriguing possibility for the first stars is that some of them may have died as a pair-instability supernova (PISN), a peculiar fate predicted for progenitor masses in the range $140~M_{\odot}\la M_{*}\la 260~M_{\odot}$ \citep{brs67,hw02}. |
. Current theory proposes a top-leavy initial mass function CALIF) for the first stars. with a characteristic mass M.z100AL. (AbelYoshidaetal. 2008). | Current theory proposes a top-heavy initial mass function (IMF) for the first stars, with a characteristic mass $M_{*}\ga 100~M_{\odot}$ \citep{abn02,bcl02,on07,yoh08}. |
. Together with the expectation that mass loss due to radiativelydriven winds is negligible at low metallicities (sudzitzki2002).. one arrives at the robust expectation that at least a fraction of Pop ITI stars should have died as PISNe. | Together with the expectation that mass loss due to radiatively–driven winds is negligible at low metallicities \citep{kudritzki02}, one arrives at the robust expectation that at least a fraction of Pop III stars should have died as PISNe. |
Compared to conventional core-collapse SNe. a PISN is distinguished by not leaving belind a compact renmnaut. | Compared to conventional core-collapse SNe, a PISN is distinguished by not leaving behind a compact remnant. |
Iustead. the exploding star is completely disrupted. audeff metals produced inside the Pop III star are released iuto the πππτοπο», leading to metal vields of y=Mz/AM.~0.5 (leger&Woosley 2002). | Instead, the exploding star is completely disrupted, and metals produced inside the Pop III star are released into the surroundings, leading to metal yields of $y=M_{\rm Z}/M_{*}\sim 0.5$ \citep{hw02}. |
. An abundant occurrence of PISNe in the carly universe could thus have rapidly established a bedrock of iuetals at least locally. of order Z>10?Z.. (Creifetal. 2007). | An abundant occurrence of PISNe in the early universe could thus have rapidly established a bedrock of metals, at least locally, of order $Z\ga 10^{-3}~Z_{\odot}$ \citep{greif07}. |
. Receutlv. the extremely Inninous SN 2007hi was observed iu a nearby galaxy. (CalYamictal.2009)... with characteristics. such as very large Ni masses. that seein to unauubigeuouslv poiut to a PISN origin. | Recently, the extremely luminous SN 2007bi was observed in a nearby galaxy \citep{gal-yam09}, with characteristics, such as very large Ni masses, that seem to unambiguously point to a PISN origin. |
This discovery ereatlv strenethens the possibility for fiudiug hese events at high redshifts as well. | This discovery greatly strengthens the possibility for finding these events at high redshifts as well. |
We here carry out cosmnological simulations tracing he detailed assembly process of a primordial galaxy. aking iuto account the feedback effects from Pop III star formation inside the nuiünihalo progenitor svstenis. | We here carry out cosmological simulations tracing the detailed assembly process of a primordial galaxy, taking into account the feedback effects from Pop III star formation inside the minihalo progenitor systems. |
Qur work exteuds the study bv οἱetal.(2008).. which ollowed the virialization of gas in the ealactic potential well under the idealized assumption of no such feedback. | Our work extends the study by \citet{greif08}, which followed the virialization of gas in the galactic potential well under the idealized assumption of no such feedback. |
Specifically, we include radiative feedback. leading to he build-up of IT regious around Pop IIT stars (ce...20072:Greifetal. 2009b).. as well as the mechanical and chemical feedback frou a single PISN that explodes in the earliest münibhalo progenitor. | Specifically, we include radiative feedback, leading to the build-up of H regions around Pop III stars \citep[e.g.,][]{wan04,abs06,yoshida07,greif09b}, as well as the mechanical and chemical feedback from a single PISN that explodes in the earliest minihalo progenitor. |
The latter feedback refers to the additional cooling that becomes available in nietal-euriched sas. allowing the gas to reach lower temperatures. and to possibly fragment iuto low-mass Population II (Pop II) stars. | The latter feedback refers to the additional cooling that becomes available in metal-enriched gas, allowing the gas to reach lower temperatures, and to possibly fragment into low-mass Population II (Pop II) stars. |
It has been suggested that | It has been suggested that |
We then can write a new generalized likelihood as: The Appenclix shows an easy wav used by ito find a maximum to log£. | We then can write a new generalized likelihood as: The Appendix shows an easy way used by to find a maximum to $\log\likelihood$. |
For simplicity we consider only circular orbits. | For simplicity we consider only circular orbits. |
Eccentricity introduces two elfects (Mazeh. Latham Stelanik 1996). the first of which is the dependence of A. on the eccentricity €. | Eccentricity introduces two effects (Mazeh, Latham Stefanik 1996), the first of which is the dependence of $K$ on the eccentricity $e$. |
Equation (3) should include an additional factor of (1—€?)|?. which causes A. to increase for increasing e. | Equation (3) should include an additional factor of $(1- e^2)^{-1/2}$, which causes $K$ to increase for increasing $e$. |
The other factor is the dependence of the detection threshold A54, on ο, | The other factor is the dependence of the detection threshold $K_{\rm min}$ on $e$ . |
Our simplibving assumption about the constancy of Av, throughout the sample breaks down when we consider eccentric orbits. | Our simplifying assumption about the constancy of $K_{\rm min}$ throughout the sample breaks down when we consider eccentric orbits. |
This is so because lor eccentric orbits the velocity variation tends to concentrate around the periastron passage. and (herelore A4, increases for increasing eccentricity. | This is so because for eccentric orbits the velocity variation tends to concentrate around the periastron passage, and therefore $K_{\rm min}$ increases for increasing eccentricity. |
These two ellects tend to cancel each other (Fischer Marcy 1992). the net effect depends on the characteristics of the observational search. | These two effects tend to cancel each other (Fischer Marcy 1992), the net effect depends on the characteristics of the observational search. |
By running numerical simulations Mazeh et ((1996) have louncl that if the detection limit depends on the ram.s. | By running numerical simulations Mazeh et (1996) have found that if the detection limit depends on the r.m.s. |
scatter of the observed raclial-velocity measurements. the two ellects cancel each other for any reasonable eccentricity. | scatter of the observed radial-velocity measurements, the two effects cancel each other for any reasonable eccentricity. |
We therefore chose not to include the eccentricity of the planets in our analvsis. | We therefore chose not to include the eccentricity of the planets in our analysis. |
To apply (to (he current. known sample of extrasolar planets we considered all known planets aud brown dwarls as of April 2001. | To apply to the current known sample of extrasolar planets we considered all known planets and brown dwarfs as of April 2001. |
We consider only G- or Ix-star. primaries and therefore excluded Gls 876 from the sample. | We consider only G- or K-star primaries and therefore excluded Gls 876 from the sample. |
Obviously. (he present sample in not complete. | Obviously, the present sample in not complete. |
In particular. not all planets with long periods and small induced rachal-velocity amplitucles were discovered and/or announced. | In particular, not all planets with long periods and small induced radial-velocity amplitudes were discovered and/or announced. |
To acquire some degree of completeness to our sample we have decided. somewhat arbitrarily. to exclude planets with periods longer than 1500 days and with racial-velocityamplitudes | To acquire some degree of completeness to our sample we have decided, somewhat arbitrarily, to exclude planets with periods longer than 1500 days and with radial-velocityamplitudes |
probabilities bv “chance”. | probabilities by "chance". |
Tustead of this. there are 16 such cases; which can still be by chance. | Instead of this, there are 16 such cases, which can still be by chance. |
The cispersion is v0.95«11=3.23. and 2«3.23=6.16»(1G1ll)- | The dispersion is $\sqrt{0.95\times 11}= 3.23$, and $2\times 3.23 = 6.46 > (16-11) = 5$. |
The probability to obtain 16 cases instead of 11 by chance is bigeer than 54. | The probability to obtain 16 cases instead of 11 by chance is bigger than $5\%$. |
Aeun. similarly to all GRBs aud to the short subclass. the null-hvpothesis of intrinsic isotropy of long CRBs should be rejected. | Again, similarly to all GRBs and to the short subclass, the null-hypothesis of intrinsic isotropy of long GRBs should be rejected. |
The results of paper may be stuarized as follows. | The results of paper may be summarized as follows. |
First. the ο=2. an=0 spherical harmonic shows that there is a clear anisotropy on the sienificance level >99.9%4 m the distribution of all 2281 CRBs. | First, the $n=2$, $m=0$ spherical harmonic shows that there is a clear anisotropy on the significance level $>99.9\%$ in the distribution of all 2281 GRBs. |
This fact is expected from the BATSE sky exposure function. and is interpreted as au artificial “iustrmuecntal” anisotropy. | This fact is expected from the BATSE sky exposure function, and is interpreted as an artificial "instrumental" anisotropy. |
Second. there is a clear iutriusic ausotropy of 253 “intermediate” CRBs at the >97% sienificauce leve a tothe 7 23021 tern. | Second, there is a clear intrinsic anisotropy of 253 "intermediate" GRBs at the $\geq 97\%$ significance level due to the $n=3$ $m=-1$ term. |
Third. both the [19 short CRBs axd the 998 lone GRBs. respectively. aud also all 2281 GRE) do exhibit anidsotropies on statistically high cnough siguificauce levels. | Third, both the 419 short GRBs and the 998 long GRBs, respectively, and also all 2281 GRBs do exhibit anisotropies on statistically high enough significance levels. |
All these results ave interesting. becase the departure from) dutrinsic isotropy just for the now “intermediate” subclass having the smallest muuber of CRBs ids surprising. | All these results are interesting, because the departure from intrinsic isotropy just for the new "intermediate" subclass having the smallest number of GRBs is surprising. |
Of course. the sienificauce level should still be iereased. because the 99% level (or even the 99.9% ) ls desirable. | Of course, the significance level should still be increased, because the $99\%$ level (or even the $99.9\%$ ) is desirable. |
However. the 97% significance level is already renirkable. aud hence surely should be /aunounuced. | However, the $97\%$ significance level is already remarkable, and hence surely should be announced. |
The vefinement of siguificance level may follow either from the methods used iu this paper aud iu Dalázzs ot al ( | The refinement of significance level may follow either from the methods used in this paper and in Balázzs et al. ( |
19908) (for cxample. a better estimation of lunaccuracv will notueed the drastic truncation »< or from wholly different statistical methods (I&eudall&Stuart1969.. Dagolvetal.1998)). | 1998) (for example, a better estimation of inaccuracy will notneed the drastic truncation $n <\sqrt{N}/3$ ), or from wholly different statistical methods \cite{kendall}, , \cite{bagoly}) ). |
There is considerable evidence from numerica experiments hat haloes formed in CDM simulations are well-»proximated. by two-coniponent power [aw 5oliles of the rm poxer(1rire) Cor pxre(1|(ré£r.)oO)1 | There is considerable evidence from numerical experiments that haloes formed in CDM simulations are well-approximated by two-component power law profiles of the form $\rho\propto r^{-\gamma}(1+r/r_s)^{\gamma-3}$ or $\rho\propto r^{-\gamma}(1+(r/r_s)^{3-\gamma})^{-1}$. |
he first form was presented bv Navarro. Frenk White 997. hereafter NEW) based on a suite of simulations with different initial density Ductuation spectra anc cosmological xuwameters. | The first form was presented by Navarro, Frenk White (1997, hereafter NFW) based on a suite of simulations with different initial density fluctuation spectra and cosmological parameters. |
They sugeest 5= Lis the universal exponent. | They suggest $\gamma=1$ is the universal exponent. |
Aoore et al. ( | Moore et al. ( |
1998) find. that the form of the inner profile (5) is resolution cependent. | 1998) find that the form of the inner profile $\gamma$ ) is resolution dependent. |
With 310" particles within he virtal raclius. this group finds -1.4. | With $3\times10^{6}$ particles within the virial radius, this group finds $\gamma=1.4$. |
Jing Suto (2000) find that t1e inner slope is not universal but. varies depending on envronment: they find 5=1.5.1.3 and 1.1 or galaxy-. grouj». and cluster-mass haloes. respectively. | Jing Suto (2000) find that the inner slope is not universal but varies depending on environment; they find $\gamma=1.5, 1.3$ and $1.1$ for galaxy-, group-, and cluster-mass haloes, respectively. |
Besides issues of n-bods resolution and. methodology. attempts to expl:ün the cliserepancy of these collisionless simulations and astronomical observations include new laws of physics (Sperge Steinhardt 2000) and the elfects of gas dissipation. (Vittles Couchman 1999. Frenk et al. | Besides issues of n-body resolution and methodology, attempts to explain the discrepancy of these collisionless simulations and astronomical observations include new laws of physics (Spergel Steinhardt 2000) and the effects of gas dissipation (Tittley Couchman 1999, Frenk et al. |
2000. Alvares. Shapiro Nartel 2000). | 2000, Alvares, Shapiro Martel 2000). |
The suggestion of some sort of universal profile. and more eenerallv t1ο physics of dissipationless collapse or violent relaxation has a long history. | The suggestion of some sort of universal profile and more generally the physics of dissipationless collapse or violent relaxation has a long history. |
Phe general problem of stellar dynamics in he presence of large Ductuations is very Πο and. muc-— of this focuses on first-principle forms [or the phase-space distribution. | The general problem of stellar dynamics in the presence of large fluctuations is very difficult and much of this focuses on first-principle forms for the phase-space distribution. |
The companion paper approaches this pronem as near-equilibrium evolution in à noisy environmen SCe Weinberg 2000. Paper 1. for a review of recent. theoretiCa work on this subject). | The companion paper approaches this problem as near-equilibrium evolution in a noisy environment (see Weinberg 2000, Paper 1, for a review of recent theoretical work on this subject). |
Here we apply the theory from a>er 1 to follow the evolution of haloes Wil ra number of dillerent. concentrations and. shapes. | Here we apply the theory from Paper 1 to follow the evolution of haloes with a number of different concentrations and shapes. |
The near-equilibritua resriction allows development of a solvaxe evoutionarv cquation given some initial condition. | The near-equilibrium restriction allows development of a solvable evolutionary equation given some initial condition. |
We find tha the profile rapidly assumes a double power law form with 5zc1.5 independent of the initial model. | We find that the profile rapidly assumes a double power law form with $\gamma\approx1.5$ independent of the initial model. |
So in he presence of noise. a cquasi-self-similar profile does appear anc in this sense represents the near-equilibrium: limit of vioent relaxation. | So in the presence of noise, a quasi-self-similar profile does appear and in this sense represents the near-equilibrium limit of violent relaxation. |
Ehe inner power-law profile evolves slowly ifthe noise is applied over long periods. | The inner power-law profile evolves slowly if the noise is applied over long periods. |
This evolution of the | This evolution of the |
Since the nebular abundance sttdies of the Magellanic Clouds it the nid 1970s (e.e..tet?) Peimber Torres-Peimbert 197I. 1976). it |as beei suggested that there 1velit be a correlation betwee ρακα mass and the metallicity of the inte‘stellar inediuim. | Since the nebular abundance studies of the Magellanic Clouds in the mid 1970's (e.g., Peimbert Torres-Peimbert 1974, 1976), it has been suggested that there might be a correlation between galaxian mass and the metallicity of the interstellar medium. |
This bas beer supported by observatio[unS ol H II regions in irregular galaxies wv Lequeux et ((1979). Taleit (1980). μα Davidso (1981). Skillman e ((1989). Richer McCall (19€)5). Miller (1996). ard vau Zee et ((1997b). | This has been supported by observations of H II regions in irregular galaxies by Lequeux et (1979), Talent (1980), Kinman Davidson (1981), Skillman et (1989), Richer McCall (1995), Miller (1996), and van Zee et (1997b). |
There have been suggestlous that he fuudainenta relationship may iol be between lass alc inetalliciy. but between sulace deusiyo All ineallicily (Moud. Iristiau. DaCosta 1933: Bothu Moule Loss: Ecuatuicds Phillips 19589. | There have been suggestions that the fundamental relationship may not be between mass and metallicity, but between surface density and metallicity (Mould, Kristian, DaCosta 1983; Bothun Mould 1988; Edmunds Phillips 1989). |
Based on obseryations of two dEs in the ASl group. Caldwell et ((1998) coicluded tha luninosv. al[1l| not su‘face brielituess. is the key )araineter determiing metallicity for the cEs. | Based on observations of two dEs in the M81 group, Caldwell et (1998) concluded that luminosity, and not surface brightness, is the key parameter determining metallicity for the dEs. |
henmieutt Skillman (2001). [1rawiug ¢1 observatious [rom the literature. couclidecl that Iumiuosiv. aud not surface brightuess. is probabvy the important parameter for cetermiuit& metallicity in dls aso. | Kennicutt Skillman (2001), drawing on observations from the literature, concluded that luminosity, and not surface brightness, is probably the important parameter for determining metallicity in dIs also. |
In Figure 8. we compa‘fe Οἱ) new observations of the Scuptor Cu(p «Is to the coimpilation of Local Group dds by lateo (1998). | In Figure 8, we compare our new observations of the Sculptor Group dIs to the compilation of Local Group dIs by Mateo (1998). |
We have added iu the oxygen abutdauce for the Seul)or Croup dl 4113 = ESO 2[2-C05 (see paper 1) reported by Àiller (1996) as 12 + log (O/H) = 1.07. | We have added in the oxygen abundance for the Sculptor Group dI A143 $=$ ESO 245-G05 (see paper 1) reported by Miller (1996) as 12 + log (O/H) $=$ 7.97. |
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